{
    "0601/astro-ph0601272_arXiv.txt": {
        "abstract": "We consider the properties of a hyperaccretion model for gamma-ray bursts (GRBs) at late times when the mass supply rate is expected to decrease with time. We point out that the region in the vicinity of the accretor and the accretor itself can play an important role in determining the rate of accretion, and its time behavior, and ultimately the energy output. Motivated by numerical simulations and theoretical results, we conjecture that  the energy release can be repeatedly stopped and then restarted by the magnetic flux accumulated around the accretor. We propose that the episode or episodes when the accretion resumes correspond to X-ray flares discovered recently in a number of GRBs. ",
        "introduction": "Gamma-Ray Bursts (GRB) are generally believed to be powered by hyperaccretion onto a compact, stellar mass object. The total amount of the available fuel is considered to be the key factor determining the burst duration. Within merger scenarios for short-duration GRBs, a neutron star (NS) is  accreted onto another NS or onto a stellar mass black hole (BH; e.g, Paczy\\'nski 1986, 1991; Eichler et al. 1989; Narayan et al. 1992; Fryer et al. 1999). Within the collapsar model for long-duration GRBs, up to 20 $\\MSUN$ of a stellar envelope collapses onto the star's core which is a NS or a BH (e.g., Woosley 1993; Paczy\\'nski 1998; MacFadyen \\& Woosley 1999; Popham, Woosley, \\& Fryer 1999; Proga et al. 2003). For short- and long-duration events, the accretion rate, $\\MDOT_a$ must be of order of 1~$\\MSUN~{\\rm s^{-1}}$, yielding a duration of less than a few seconds for the former and a duration as long as tens to hundreds of seconds for the latter. These duration estimates are made under the assumption that all the available fuel is accreted during the GRB activity at a time-averaged constant rate.   Recent GRB observations obtained with {\\it Swift} motivate us to review the above assumption and some other aspects of  GRB models. In particular, early X-ray afterglow lightcurves of nearly half of the long-duration GRBs show X-ray flares (Burrows et al. 2005; Romano et al. 2006; Falcone et al. 2006). X-ray flares are also found to follow the short-duration GRB 050724 (Barthelmy et al. 2005) whose host galaxy is early-type, which is consistent with the merger origin. The flares generally rise and fall rapidly, with typical rising and falling time scales much shorter than the epoch when the flare occurs. This time behavior strongly supports the ``internal'' origin of the flares (Burrows et al. 2005; Zhang et al. 2006; Fan \\& Wei 2005), in contrast to the ``external'' origin of the power-law decay afterglows. The internal model not only offers a natural interpretation of the rapid rise and decay behavior of the flares, but also demands a very small energy budget (Zhang et al. 2006). Within this picture, the data require a {\\it restart} of  the GRB central engine (i.e., a restart of accretion).  Fragmentations in the collapsing star (King et al. 2005) or in the outer parts of the accretion disc (Perna et al. 2006) have been suggested to be responsible for the observed episodic flaring behavior. These two flare models appeal to one of the basic ingredients of an accretion powered engine -- the mass accretion rate -- and conjecture that the episodic energy output is driven by  changes in the mass supply and subsequently accretion rate. In this picture, the inner part of the accreting system {\\em passively} responds to changes in the accretion flow at larger radii.  Here, we point out that the region in the vicinity of the accretor and the accretor itself can play an important role of determining the rate and time behavior of the accretion and the energy output. In particular, we conjecture that  the energy release can be repeatedly stopped and then restarted, provided the mass supply rate decreases with time even if the decrease is smooth. For both merger and collapsar GRB models, a decrease of the mass supply rate is expected, especially in the late phase of activity, because the mass density decreases with increasing radius. In our model, we appeal to the fact that, as mass is being accreted onto a BH, the magnetic flux is accumulating in the vicinity of the BH. Eventually, this magnetic flux must become dynamically important and affect the inner accretion flow, unless the magnetic field is very rapidly diffused. In the remaining part of the paper we list and discuss theoretical arguments and results from a variety of numerical magnetohydrodynamic (MHD) simulations of accretion flows that support our model. We also provide analytic estimates to show that our model can quantitatively account for the observed features of the flares.   \\section[]{Magnetic model for GRBs and their flares}  \\subsection{Insights from numerical models}  Generally, our model for the flares is based on the results from the numerical simulation of an MHD collapsar model for GRBs carried out by Proga et al. (2003) and the results from a number of simulations of radiatively inefficient accretion flows (RIAFs) onto a BH (Proga \\& Begelman 2003, PB03 hereafter; Igumenshchev, Narayan, \\&Abramowicz 2003, INA03 hereafter). Proga et al. (2003) performed  time-dependent axisymmetric MHD simulations of the collapsar model.  These MHD simulations included a realistic equation of state, neutrino cooling, photodisintegration of helium, and resistive heating. The progenitor was assumed to be spherically symmetric but with spherical symmetry broken by the introduction of a small, latitude-dependent angular momentum and a weak split-monopole magnetic field. The main conclusion from the simulations is that, within the collapsar model, MHD effects alone are able to launch, accelerate and sustain a strong polar outflow. The MHD outflow provides favorable initial conditions for the subsequent production of a baryon-poor fireball (e.g., Fuller, Pruet \\& Abazajian 2000; Beloborodov 2003; Vlahakis \\& K$\\ddot{\\rm o}$nigl 2003; M\\'{e}sz\\'{a}ros 2002), or a magnetically dominated ``cold fireball'' (Lyutikov \\& Blandford 2002), though the specific toroidal magnetic field geometry Proga et al. derived differs from some of these models (e.g., Vlahakis \\& K$\\ddot{\\rm o}$nigl 2003; Lyutikov \\& Blandford 2002). The latest Swift UV-Optical Telescope (UVOT) observations indicate that the early reverse shock emission is generally suppressed ( Roming et al. 2005), which is consistent with the suggestion that at least some GRBs are Poynting-flux-dominated outflows (Zhang \\& Kobayashi 2005).  To study the extended GRB activity, one would like to follow the collapse of the entire star. However, such studies are beyond current computer and model limits. Therefore, we explore instead the implications of the published simulations and consider the physics of the collapsing star to infer the properties and physical conditions in the vicinity of a BH during the late phase of evolution, i.e., when a significant fraction of the total available mass is accreted.  The long time evolution of axisymmetric MHD accretion flows was studied by PB03 who explored simulations very similar to those performed by Proga et al. (2003) but with much simpler micro physics (i.e., an adiabatic equation of state, no neutrino cooling or photodisintegration of helium). Proga et al. (2003) found that despite the more sophisticated micro physics of the MHD collapsar simulations the flow cooling is dominated by advection not neutrino cooling. As a result, the early phase of the time evolution, and the dynamics of the innermost flow, are very similar in both the  RIAF simulations and the collapsar simulations. In particular, after an initial transient behavior, the flow settles into a complex convolution of several distinct, time-dependent flow components including an accretion torus, its corona and outflow, and an inflow and outflow in the polar funnel (see the left panel in Fig. 1 for a schematic picture of such a flow). The accretion through the torus is facilitated by the magnetorotational instability (MRI, e.g., Balbus \\& Halwey 1991) which also dominates the overall dynamics of the inner flow.  In the remaining part of the paper, we will assume that the late evolution of the MHD collapsar simulations is similar to the late evolution of the RIAFs simulations. This assumption is justifiable because the flows in the collapsar and RIAFs simulations are similar during the early phase of the evolution (i.e., their dymanics and cooling are dominated by MRI and advection, respectively)   The late evolution of RIAFs shows that the torus accretion can be interrupted for a short time by a strong poloidal magnetic field in the vicinity of a BH. This result is the main motivation for this paper, as it shows that  the extended GRB activity may be a result of an accretion flow modulated by the ``magnetic-barrier'' and gravity. Because this barrier halts the accretion flow intermittently (see Figs.~6 \\& 8 in PB03), the accretion rate is episodic (see Fig.3 of PB03). This potentially gives a natural mechanism for flaring variability in the magnetic-origin models of GRBs as we first mentioned in Fan, Zhang \\& Proga (2005; see the middle panel of Fig. 1 here, for a cartoon picture of the accretion halted by the magnetic-barrier.)  The importance of accumulating of the magnetic flux has been explored and observed by others in various astrophysical contexts (e.g., Bisnovatyi-Kogan \\& Ruzmaikin 1974, 1976; Narayan, Igumenshchev \\& Abramowicz 2003; INA03). In particular, INA03 carried out a three-dimensional (3D) MHD simulation (their model B) to late model times. They found that the magnetic flux accumulates, initially near the BH and then farther out, and the field becomes dynamically dominant. At late times, mass is able to accrete only via narrow streams, in a highly nonaxisymmetric manner (see also Narayan et al. 2003).  The main difference between PB03's and INA03's results is the extent and duration of the magnetic dominance. In PB03, the magnetic dominance is a {\\em transient} whereas in INA03 is a {\\em persistent} state. The reason for this difference is the treatment of the magnetic field: for the initial conditions, PB03 used the split-monopole magnetic field and any changes in the magnetic flux near the BH during the evolution are due to the chaotic, small-scale fields generated in the disc. The detailed analysis show that the disc properties in PB03's simulations are determined by MRI. In particular, MRI is responsible for the complex field structure and for the disc toroidal field being one or even two orders of magnitude higher than the poloidal field (see figs. 9 and 10 in PB03 and fig. 3 in Proga et al. 2003). On the other hand, in their model B, INA03 set up a poloidal field configuration in the injected gas in such a way that the portion of the material that accretes always carries in the same sign of the vertical component of the magnetic field. The simulations carried out by PB03 and INA03 differ also in the assumed geometry (axisymmetric versus fully 3D). INA03 and PB03 do not explore all cases including the case where the external or initial field has zero net flux or the field with the poloidal component changing sign on length scales much smaller than the size of the mass reservoir \\footnote{ In the case where the initial or external flux has zero-net flux, a large scale coherent field might in some circumstances be generated by MRI (e.g., Livio, Pringle, \\& King  2003). If so the central magnetic flux could vary with time but still be dynamical signifacant for some periods of time.}. Additionally, these simulations also do not give definitive answers to the problems for which they were designed.  Nevertheless, they give interesting insights into the general problem of MHD accretion flows. In particular, they suggest that magnetic fields can provide an important parameter determining the time scale for the accretion; i.e., it can be significantly longer than the local dynamical time scale. This can have important implications for the observed X-flares in GRBs, as we argue here, and for X-ray spectral states for BH binaries as discussed by Spruit \\& Uzdensky (2005, SU05 hereafter). In fact, the work by SU05 describes very well the general physics and theory of magnetic flux accumulated by an accretion flow. Therefore we now turn our attention to some theoretical aspects of the problem as presented by SU05.  \\subsection{Theory of the magnetic barrier and accretion flow}  SU05 considered a new mechanism of efficient inward transport of the large-scale magnetic field through a turbulent accretion disc. The key element of the mechanism is concentration of the external field into patches of field comparable in strength to  the MRI turbulence in the disc. They focused on how to increase the magnetic flux at the center in the context of BH binaries. In particular, they argue that the capture of external magnetic flux by accretion disc and its subsequent compression in the inner regions of the disc may explain both changes in the radiation spectrum and jet activity in those objects. However, their model and physical arguments are generic and applicable to our problem.  One can expect that as the strength of the magnetic field increases at the center, the field may eventually suppress MRI turbulence and reduce the mass accretion rate and the power in the outflow. This should be the case especially for GRBs because the mass inflow rate at the late time is most likely much lower than at the early time. The disc may become a Magnetically-Dominated Accretion Flow (MDAF) as proposed by Meier (2005) or the fields in the polar funnel can expand toward the equator and reconnect as in PB03's simulations. In the latter, the torus is pushed outward by the magnetic field. At this time, the gas starts to pile up outside the barrier; eventually it can become unstable to interchange instabilities at the barrier outer edge as suggested by SU05 or  the gas in the torus can squash the magnetic field (compare Fig.~5 and 6 in PB05 or the middle and right panel in Fig.1 here).  When interchange instabilities operate, magnetic flux from the bundle mixes outward into the disc while the disc material enters the barrier. In the accretion disk context, interchange instabilities have been studied by a few authors (e.g., Spruit et al. 1995; Lubow \\& Spruit 1995; Stehle 1996; Stehle \\& Spruit 2001; Li \\& Narayan 2004). These studies showed that the onset of small-scale modes typical of interchanges (as in Rayleigh-Taylor instabilities) takes place only at rather large field strengths, due to a stabilizing effect of the Keplerian shear. The interchange instability operates at moderate field strengths, but only at low shear rates (less than Keplerian). However for most of the time, we expect high shear rates in a torus because a low shear torus quickly becomes Keplerian due to MRI (e.g., PB03 and Proga et al. 2003). We note that SU05 interpreted INA03 accretion through the barrier, in the form of blobs and streams as a product of interchange instabilities.  SU05 also suggested that the field strength at which these instabilities become effective is most usefully expressed in terms of the degree of support against gravity provided by the magnetic stress $B_R B_Z$. According to SU05, the instabilities become effective when the radial magnetic force, $F_m\\sim 2 B_R B_Z/4\\pi$, is of the order of a few percent of the gravitational force, $F_g=GM\\Sigma/R^2$, where $M$ is the central mass, $R$ is the radius, and $\\Sigma$ is the surface density. For $B_R \\approx B_Z$, there is a range in field strengths between the value at which MRI turbulence is suppressed and the value where dynamical instability of the barrier itself sets in, where no known instability operates (Stehle \\& Spruit 2001). In this range, the disk material cannot mix or penetrate the magnetic field accumulated at the center (e.g., the middle panel of Fig. 1). Instead, mass builds up outside a region with such field strengths until the magnetic field at the center is compressed enough for instability to set in.  Thus, both numerical work and theoretical models of magnetized accretion flows show that the inner most part of the flow and accretor can respond {\\em actively} to changes of the accretion flow at larger radii. In particular, the inner most accretion flow can be halted for a very long time as shown by INA03 or it can be repeatedly halted and reactivated as shown in PB03.  \\subsection{Analytic estimates}  We finish this section with order-of-magnitude estimates of a few key features of our X-ray flare model. We start by estimating the strength and flux of magnetic field required to support the gas. The gas of the surface density, $\\Sigma$ can be supported against gravity by the magnetic tension if $F_g\\sim F_m$. The surface density can be estimated from $\\Sigma=\\MDOT /2\\pi R \\epsilon v_{ff}~{\\rm g~cm^{-2}}$, where  $\\epsilon v_{ff}$ is the flow radial velocity assumed to be a fraction $\\epsilon$ of the free fall velocity, $v_{ff}$. Assuming $B_r \\approx B_z=B$, the force balance yields the field strength $B \\sim 2\\times 10^{16}~\\epsilon_{-3}^{-1/2} r^{-5/4} \\MDOT_1^{1/2} M_3^{-1} $~G, where $\\epsilon_{-3} \\equiv 10^{3} \\epsilon$, $r \\equiv R/R_S=R/(2GM_{BH}/c^2)$, $\\MDOT_1=\\MDOT/1~\\MSUN~{\\rm s^{-1}}$, and $M_3=M/3 \\MSUN$. We estimate the magnetic flux as $\\Phi\\sim\\pi r^2 R_S^2 B(r)= 5\\times10^{28}~\\epsilon_{-3}^{-1/2}r^{3/4}\\MDOT_ 1^{1/2} M_3~{\\rm cm^2~G}$ from which we obtain an estimate to the magnetospheric radius $r_m \\approx 60~\\epsilon_{-3}^{2/3} \\MDOT_1^{-2/3} M_3^{-4/3} \\Phi_{30}^{4/3}$, where $\\Phi_{30}\\equiv \\Phi/(10^{30}~{\\rm cm^2 G})$. Substituting the expression for $B$ into the expression for the surface density, one finds that a given magnetic flux can support the gas with the surface density of $\\Sigma_B=5\\times10^{19}~\\Phi_{30}^2 M_3^{-3}r^{-2}~{\\rm g~cm^{-2}}$.  To stop accretion with the hyper rate of $1\\MSUN~{\\rm s^{-1}}$ onto a 3$\\MSUN$ black hole at r=3 (i.e., for $r_m$ to be 3), the magnetic flux of order $\\Phi_{30} \\sim 0.11$ is required. We now assume that such a magnetic flux is accumulated during hyperaccretion and that it does not change with time. Under these assumptions, $r_m=300$ for the mass supply rate of $10^{-3}~\\MSUNYR$ representative of the late time evolution . This relatively large radius demonstrates one of our key points that the innermost part an accreting system can actively respond, via magnetic fields, to changes in the inflow at large radii.  To estimate the conditions needed to restart accretion, the accretion energetics and related time scales, we ask what is the mass of a disc with $\\Sigma$ high enough to reduce $r_m$ from 300 to 3 or so. To answer this question, we adopt Popham et al.' (1999) model of neutrino-dominated discs. Popham et al. assumed that neutrino cooling produces a thin disc (Shakura \\& Sunyaev 1973) for accretion rates require to power GRBs. Using  the disc solution for the density and height (eqs. 5.4 and 5.5 in Popham et al. 1999), we can express the disc surface density as $\\Sigma_\\alpha=1.8\\times10^{19}~\\alpha^{-1.2} M_3^{-0.8}\\MDOT_1r^{-1.25}$~g, where $\\alpha$ is the dimensionless parameter scaling the stress tensor and the gas pressure (Shakura \\& Sunyaev 1973). Equating $\\Sigma_B$ with $\\Sigma_\\alpha$, one can estimate the mass accretion rate of an $\\alpha$ disc and compute $M_D$ by integrating $\\Sigma_\\alpha$ over radius. For $\\Phi_{30}=0.11$ and $\\alpha=10^{-2}$ the accretion rate through the $\\alpha$ disc is  $0.03~\\MSUN~s^{-1}$ and $M_D$  for $r$ between 3 and  300  is 0.32~$\\MSUN$. This mass accretion rate is more than one order of magnitude  lower than the rate of  $\\sim 1~\\MSUN~s^{-1}$ typical for the early time evolution. Thus, our estimates are consistent with the fact that the X-ray flare luminosity is at least one or two orders of magnitude lower the prompt gamma-ray emission (see section 3). If this disc mass is a result of slow mass accumulation during the late evolutionary stage, then it will take about 400 s to rebuild the disc for the mass supply rate of $10^{-3}~\\MSUN~s^{-1}$ and 12 s to accrete all this mass at the disc accretion rate of $0.03~\\MSUN~s^{-1}$. The latter is a lower estimate for the flare duration because, for simplicity, we assumed a relatively high, {\\em constant} disc accretion rate. It is very likely that the rate changes with time as the shape of the light curve during the flares indicates. In our model, the mass supply rate controls the epochs when the flares happen: the disc is rebuilt on the time scale which increases with time because the mass supply slowdowns. Additionally, the flare duration is coupled to the epoch through  the mass of the rebuilt disc. Thus our model is capable of accounting for the observed duration - time scale correlation. ",
        "conclusions": "The detailed analysis of the X-ray flares revealed that they generally have lower luminosities (by at least one or two orders of magnitude) than the prompt gamma-ray emission. Additionally, the total energy of the flare is also typically smaller than that of the prompt emission, although in some cases both could be comparable (e.g. for GRB 050502B, Falcone et al. 2006). Moreover multiple flares are observed in some GRBs and the durations of these flares seem to be positively correlated with the epochs when the flares happen, i.e. the later the epoch, the longer the duration (O'Brien et al. 2005; Falcone et al. 2006; Barthelmy et al. 2005). The flare analysis also showed that the later the epoch the lower the flare luminosity. The above qualitative properties of the flares provide important constraints on models of them.  Perna et al.'s (2006) disc fragmentation model promises to account for the duration - time scale correlation and the duration - peak luminosity anticorrelation. However, the physical process or processes causing fragmentation are uncertain. It is also uncertain that the conditions for the disc fragmentation are met in GRB progenitors. This seems to be the case especially for the collapsar model as a relatively high rotation of the progenitor is required. We also note that magnetic fields can suppress or even prevent disc fragmentation (e.g., Banerjee \\& Pudritz 2006).  Here, we propose that the X-ray flares in GRBs are consequences of the fact that during the late time evolution of a hyperaccretion system the mass supply rate should decrease with time while the magnetic flux accumulating around a BH should increase. In particular, we point out that the flux accumulated during the main GRB event can change the dynamics of the inner accretion flow. We argue that the accumulated flux is capable of halting intermittently the accretion flow. In our model, the episode or episodes when the accretion resumes correspond to X-ray flares. A comparison of our analytic estimates from Section 2.3 with the observed X-ray flare characteristics, shows that our model is not only physically based but also can both qualitatively and quantitatively account for some aspects of the flares -- such as the peak times. In general, our model fits under the general label of the magnetic jet model for GRBs as we appeal to the magnetic effects to play the key role not only during the main event but also during the late evolution. The importance of the magnetic effects for the X-ray flares can be argued based on energy budget of the accretion model (Fan et al. 2005).  The X-ray flares discovered in GRBs are relatively new and unexpected phenomena. They give a strong incentive to apply the existing models of hyperaccretion systems to circumstances where the mass supply is reduced. Studies of this kind should reveal whether  one  needs to introduce additional physics in order to explain the flares. If so one should explore the effects of this on the early evolution of GRBs and check whether they are consistent with GRBs observations. Our X-ray flare model has the advantage that it is essentially the same as the MHD collapsar model for GRBs, with only one justifiable change in a key physical property of the collapsar model: a decrease of the mass supply rate with time."
    },
    "0601/astro-ph0601044_arXiv.txt": {
        "abstract": "We study the bending of light for static spherically symmetric (SSS) space-times which include a dark energy contribution. Geometric dark energy models generically predict a correction to the Einstein angle written in terms of the distance to the closest approach, whereas a cosmological constant $\\Lambda$ does not. While dark energy is associated with a repulsive force in cosmological context, its effect on null geodesics in SSS space-times can be attractive as for the Newtonian term. This dark energy contribution may be not negligible with respect to the Einstein prediction in lensing involving clusters of galaxies. Strong lensing may therefore be useful to distinguish $\\Lambda$ from other dark energy models. ",
        "introduction": "It is still unclear what drives the universe into acceleration recently. While a cosmological constant $\\Lambda$ is the simplest explanation, its value seems completely at odd with the naive estimate of the vacuum energy due to quantum effects.  An alternative to $\\Lambda$ is obtained considering a dynamical degree of freedom added to the primordial soup. Since dynamical dark energy (dDE) varies in time and space, its fluctuations are potentially important in order to distinguish it from $\\Lambda$ \\cite{CDS}. Another possibility for the explanation of the present acceleration of the universe is given by the geometry itself, through a modification of Einstein gravity at large distances \\cite{DDG}: these are known as geometric dark energy (gDE) models. A non vanishing mass for the graviton is among these possibilities \\cite{GG}.  {\\em Cosmological} observations, such as those coming from Supernovae, cosmic microwave background (CMB) anisotropies and large scale structure (LSS), have not been able to discriminate among dDE models yet (see \\cite{UIS} for updated constraints and forecasts). It is therefore important to explore observational tests at {\\em astrophysical} level for objects which are detached from the cosmological expansion. ",
        "conclusions": "We have discussed light bending in SSS with DE. The importance of general relativistic tests, such as the perihelion precession, has been already emphasized for the DGP model \\cite{DGZ,LS}, while little attention was previously paid to light bending. These tests are complementary to the observational signatures of dark energy in cosmological context, mainly based on the behaviour of perturbations. In cosmology, an important difference between $\\Lambda$ and dDE or gDE is the presence of DE perturbations in the latter case, which are at least gravitationally coupled to the other types of matter. Such DE perturbations are therefore a key point to distinguish $\\Lambda$ from dDE or gDE in CMB and LSS, and sometimes may become so important to strongly constrain models \\cite{CDS,CF,koyama} with respect to what Supernovae data can do.  In this article we have shown that in objects which have detached from the expansion of the universe, $\\Lambda$ may be distinguishable from other DE models through the bending of light. In order to link our findings with observations, we should insert $\\varphi$ in the lens equation, e.g. \\cite{schneider}: \\begin{equation} \\theta-\\beta=\\frac{d_{SL}}{d_{OS}} \\varphi \\end{equation} where $\\theta$ and $\\beta$ are the angular positions of the image and of the source measured respect to the line from the observer to the lens; $d_{LS}$ and $d_{OS}$ are the angular diameter distances between the lens and the source and between the observer and the source, respectively. On considering for simplicity alignment between the lens and the source, an Einstein ring forms with angle $\\theta_{E} = \\theta (\\beta = 0)$. From our results, $\\theta_E$ is affected by both the non-perturbative SSS potential around the lens ($ \\varphi \\ne 4 G M/r_0$ if DE $\\ne \\Lambda$) and the cosmology of a given model. The gDE corrections to the Einstein deflection angle for clusters in Eq. (\\ref{dephi}) are as important as the cosmology for an observable as $\\theta_E$. The differential of $\\theta_E$ is \\be \\frac{\\Delta \\theta_E}{\\theta_E} = \\frac{\\Delta \\varphi}{\\varphi} + \\Delta \\ln \\frac{d_{SL}}{d_{OS}} \\,, \\ee which reveals how cosmological information is encoded just in the second term to the right. By considering the cosmology of the DGP model for instance \\cite{LUE}, one finds that the second term is $\\sim -0.06 (\\Delta \\Omega_{\\rm M}/\\Omega_{\\rm M}) + \\Delta H_0/H_0$ for a source and a lens located at $z = 1$ and $z =0.3$, respectively ($\\Omega_{\\rm M} \\sim 0.3$ and $H_0$ are the present matter density and Hubble parameter, respectively). On considering the uncertainties on the cosmological parameters of the order of percent, this simple quantitative example shows how the corrections to the Einstein deflection angle we have found in Table I should be taken into account in the study of strong lensing by clusters.  We believe that results similar to what we have found here for gDE models, might occur for dDE scenarios as well, in which the non-asymptotically flat term is due to the non-perturbative clumping of DE into objects detached from the cosmological expansion. However, dDE models may be less predictive than gDE models: gDE contain the same number of parameters of $\\Lambda$CDM, while dDE may need more.  Let us end on noting that some of the gDE models considered here may have serious theoretical issues \\cite{BD,koyama2}, whose resolution clearly go beyond the present project. However, the main result in Eq. (\\ref{dephi}) of this paper remains valid: models alternative to general relativity with a cosmological constant predict a correction to the Einstein angle, which can be used to distinguish $\\Lambda$ from other DE models.  \\vspace{.5cm}  {\\bf Acknowledgements} We wish to thank Lauro Moscardini for many discussions and suggestions on galaxies, groups and clusters. We are grateful to Robert R. Caldwell and one of the anonymous referees for valuable comments on this project. FF and MG are partially supported by INFN IS PD51; FF and and AG are partially supported by INFN IS BO11. The authors thank the Galileo Galilei Institute for Theoretical Physics for the hospitality and the INFN for partial support during the developments of this project."
    },
    "0601/astro-ph0601334_arXiv.txt": {
        "abstract": "We present a {\\it Hubble Space Telescope} Advanced Camera for Surveys (ACS) weak-lensing study of RX J0849+4452 and RX J0848+4453, the two most distant (at $z=1.26$ and $z=1.27$, respectively) clusters yet measured with weak-lensing. The two clusters are separated by $\\sim4\\arcmin$ from each other and appear to form a supercluster in the Lynx field. Using our deep ACS $i_{775}$ and $z_{850}$ imaging, we detected weak-lensing signals around both clusters at $\\sim4\\sigma$ levels. The mass distribution indicated by the reconstruction map is in good spatial agreement with the cluster galaxies. From the singular isothermal sphere (SIS) fitting, we determined that RX J0849+4452 and RX J0848+4453 have similar projected masses of $(2.0\\pm0.6)\\times10^{14} M_{\\sun}$ and $(2.1\\pm0.7)\\times10^{14} M_{\\sun}$, respectively, within a 0.5 Mpc ($\\sim60\\arcsec$) aperture radius. In order to compare the weak-lensing measurements with the X-ray results calibrated by the most recent low-energy quantum efficiency determination and time-dependent gain correction, we also re-analyzed the archival $Chandra$ data and obtained $T=3.8_{-0.7}^{+1.3}$ and $1.7_{-0.4}^{+0.7}$~keV for RX J0849+4452 and RX J0848+4453, respectively. Combined with the X-ray surface brightness profile measurement under the assumption of isothermal $\\beta$ model, the temperature of RX J0849+4452 predicts that the projected mass of the cluster within $r=0.5~$Mpc is $2.3_{-0.4}^{+0.8}\\times10^{14} M_{\\sun}$, consistent with the weak-lensing analysis. On the other hand, for RX J0848+4453 we find that the mass derived from this X-ray analysis is much smaller ($6.3_{-1.5}^{+2.6}\\times10^{13} M_{\\sun}$) than the weak-lensing measurement. One possibility for this observed discrepancy is that the intracluster medium (ICM) of RX J0848+4453 has not yet fully thermalized. Although this interpretation is rather simplistic, the relatively loose distribution of the cluster galaxies in part supports this possibility of low degree of virialization. We also discuss other scenarios that might give rise to the discrepancy. ",
        "introduction": "It has become clear that massive clusters are not extremely rare at high redshifts ($z>0.8$) and the presence of these large collapsed structures when the age of the Universe is less than half its present value is no longer in conflict with our current understanding of the structure formation, especially in a $\\Lambda$-dominated flat cosmology. Pursuit of galaxy clusters to higher and higher redshift is important in the extension of the evolutionary sequences to earlier epochs, when the effect of the different cosmological frameworks becomes more discriminating.  A great deal of observational efforts have been made in the last decade in enlarging the sample of high-redshift clusters. X-ray surveys have provided an efficient method of cluster identification and probe of cluster properties because a hot intracluster medium (ICM) within the cluster generates strong diffuse X-ray emission and is believed to be in quasi-equilibrium with gravity. However, it is questionable how well the clusters selected by their X-ray excess can provide the unbiased representation of the typical large scale structure at the cluster redshift. If the degree of the virialization decreases significantly with redshift and is strongly correlated with X-ray temperature, the cosmological dimming $\\sim (1+z)^{-4}$ can bias our selection progressively towards higher and higher mass, relaxed structures.  Among other important approaches to detect high-redshift clusters is a red-cluster-sequence (RCS) survey using the distinctive spectral feature in cluster ellipticals. This so-called 4000\\AA~break feature is well-captured by a careful combination of two passbands, and Gladders \\& Yee (2005) recently reported 67 candidate clusters at a photometric redshift of $0.9 < z < 1.4$ from the $\\sim10$\\%  subregion of the total $\\sim100~\\mbox{deg}^2$ RCS survey field. A related method but giving a higher contrast of cluster members with respect to the background sources is to use deep near-infrared (NIR) imaging (e.g., Stanford et al. 1997) for the selection of cluster candidates. High-redshift clusters identified in these color selection methods are expected to serve as less biased samples encompassing the lower mass regime at high redshifts.  In the current paper, we study two $z\\sim1.3$ clusters, namely RX J0849+4452 and RX J0848+4453 (hereafter Lynx-E and Lynx-W for brevity), using the deep F775W and F850LP (hereafter $i_{775}$ and $z_{850}$, respectively) images obtained with the Advanced Camera for Surveys (ACS) on the $Hubble$ $Space$ $Telescope$ ($HST$). Interestingly, although these two clusters are separated by only $\\sim4\\arcmin$ from each other, they were discovered by different methods.  Standford et al. (1997) discovered Lynx-W in a NIR survey as an overdense region of the $J-K > 1.9$ galaxies and spectroscopically confirmed 8 cluster members. They also analyzed the archival ROSAT-PSPC observation of the region and found diffuse X-ray emission near the confirmed cluster galaxies. However, they could not rule out the possibility that the X-ray flux might be coming from the foreground point sources because of the PSPC PSF is too broad to identify such objects. The subsequent study of the field using the $Chandra$ observations showed that, although the previous ROSAT-PSPC observation is severely contaminated by the X-ray point sources adjacent to the cluster, the cluster is still responsible for some diffuse X-ray emission. Both the X-ray temperature and luminosity of the cluster appear to be low ($T_X\\sim1.6$~keV and $L_{bol}\\sim0.69\\times10^{44} ~ \\mbox{ergs}~\\mbox{s}^{-1}$; Stanford et al. 2001).  Lynx-E was, on the other hand, first discovered in the ROSAT Deep Cluster Survey (RDCS) as a cluster candidate and follow-up near-infrared imaging showed an excess of red ($1.8<J-K<2.1$) galaxies around the peak of the X-ray emission (Rosati et al. 1999). They also showed that five galaxies around the X-ray centroid have redshifts that are consistent with the cluster redshift at $z=1.26$ using the Keck spectroscopic observations. From the $Chandra$ data analysis, Stanford et al. (2001) determined the cluster temperature and luminosity to be $T_X=5.8_{-1.7}^{+2.8}$ keV and $L_{bol}=3.3_{-0.5}^{+0.9}\\times10^{44} \\mbox{ergs}~\\mbox{s}^{-1}$, respectively.  The rather large difference in the X-ray properties of these two clusters may be viewed as representing the characteristics of the sample obtained from different survey methods. Lynx-E, the X-ray selected cluster, has much higher X-ray temperature and luminosity than Lynx-W, the NIR-selected cluster. If the stronger X-ray emission means higher dynamical maturity, the more compact distribution of the Lynx-E galaxies provides an alternate support of this hypothesis. For dynamically relaxed systems the observed X-ray properties can be easily translated into the mass properties under the assumption of hydrostatic equilibrium. However, as we probe into the higher and higher-redshift regime, it is natural to expect that there will be more frequent occasions when the equilibrium assumption loses its validity in deriving the mass properties of the system. In addition, at $z>1$ we expect to have a growing list of low-mass clusters that are also X-ray dark because of the evasively low-temperature, as well as the substantial cosmological dimming. Therefore, it is plausible to suspect that these two $z\\sim1.3$ clusters (especially Lynx-W, the poorer X-ray system) might lie on a border where the X-ray observations alone start to become insufficient to infer the mass properties.  Weak-lensing provides an alternative approach to deriving the mass of a gravitationally bound system without relying on assumptions about the dynamical state. This can help us to probe the properties of the high-redshift clusters in lower mass regimes, where the X-ray measurements alone may not provide useful physical quantities. In our particular case, weak-lensing is an important tool to test how the masses of the two Lynx clusters at $z\\sim1.3$ compare with their X-ray measurements. Especially for Lynx-W, weak-lensing seems to be the unique route for probing the cluster mass, considering the poor and amorphous X-ray emission. Another interesting question is whether the low X-ray temperature of Lynx-W arises simply from a low mass or from a yet poor thermalization of the ICM.  However, the detection of weak-lensing signal at $z\\sim1.3$ is difficult and much more so if the lens is not very massive. In our previous investigation of the two $z\\sim0.83$ high-redshift clusters (Jee et al. 2005a, hereafter Paper I; Jee et al. 2005b, hereafter Paper II), we were able to detect clear lensing signals. They revealed the complicated dark matter substructure of the clusters in great detail. The effective source plane (defined by the effective mean redshift of the background galaxies) in Paper I and II is at $z_{eff}\\sim1.3$, corresponding to the redshift of the lenses targeted in the current paper! Therefore, the number density of background galaxies decreases substantially compared to our $z\\sim0.8$ studies and, in addition, the higher fraction of non-background population in our source sample inevitably dilutes the resulting lensing signal quite severely. Furthermore, the accurate removal of instrumental artifacts becomes more critical as stronger signals come from more distant, and thus fainter and smaller galaxies. They are more severely affected by the point-spread-function (PSF). Nevertheless, our analyses of RDCS 1252-2927 at $z=1.24$ (Lombardi et al. 2005; Jee et al. in preparation) demonstrate that weak-lensing can still be applied to clusters even at these redshifts and reveals the cluster mass distribution with high significance.  Returning to the X-ray properties, the low-energy quantum efficiency (QE) degradation of the $Chandra$ instrument can cause noticeable biases in cluster temperature measurements. Although there have been many suggestions regarding this issue, it was not until recently that a convergent prescription to remedy the situation has become available from the $Chandra$ X-ray Center\\footnote{see http://cxc.harvard.edu/ciao3.0/threads/apply\\_acisabs/ or http://cxc.harvard.edu/ciao3.2/releasenotes/}. Because we suspect that the previous X-ray analyses of the Lynx clusters suffered from the relatively insufficient understanding of this problem, we have also re-analyzed the archival $Chandra$ data to enable a fairer comparison between the weak-lensing and X-ray measurements.  Throughout the paper, we assume a $\\Lambda$CDM cosmology favored by the Wilkinson Microwave Anisotropy Probe (WMAP), where $\\Omega_M$, $\\Omega_{\\Lambda}$, and $H_0$ are 0.27, 0.73, and 71 $\\mbox{km}~\\mbox{s}^{-1}~\\mbox{Mpc}^{-1}$, respectively. All the quoted uncertainties are at the 1 $\\sigma$ ($\\sim68$\\%) level. ",
        "conclusions": "} We have presented a weak-lensing analysis of the two Lynx clusters at $\\bar{z}=1.265$ using the deep ACS $i_{775}$ and $z_{850}$ images. Our mass reconstruction clearly detects the dark matter clumps associated with the two high-redshift clusters and other intervening objects within the ACS field, including the known foreground cluster at $z=0.57$. In order to verify the significance of the cluster detection and to separate the high-redshift signal from the low-redshift contributions, we performed a weak-lensing tomography by selecting an alternate lower-redshift source plane. This second mass reconstruction does not show the mass clumps around the high-redshift clusters, while maintaining most of the other structures seen in the first mass map. This experiment strongly confirms that the weak-lensing signals observed in the first mass reconstruction are real and come from the high-redshift Lynx clusters.  Interestingly, both clusters are found to have similar weak-lensing masses of $\\sim 2.0\\times 10^{14} M_{\\sun}$ within 0.5 Mpc ($\\sim60\\arcsec$) aperture radius despite their discrepant X-ray properties. Our re-analysis of the Chandra archival data with the use of the latest calibration of the low-energy QE degradation shows that Lynx-E and W have temperatures of $T=3.8_{-0.7}^{+1.3}$ and $1.7_{-0.4}^{+0.7}$~keV, respectively. Combined with the X-ray surface brightness profile measurements, the X-ray temperature of Lynx-E gives a mass estimate in good agreement with the weak-lensing result. On the other hand, the X-ray mass of Lynx-W is much smaller than the weak-lensing estimation nearly by a factor of three. According to our experiment in \\textsection\\ref{section_mass_estimate}, it is unlikely that any foreground contamination or cosmic shear effect in weak-lensing measurement causes this large discrepancy.  Apart from a simplistic, but valid possibility that Lynx-W might have a filamentary structure extended along the line of sight, yielding a substantial, projected mass but with yet only low-temperature thermal emission, we can also consider the self-similarity breaking (e.g., Ponman et al. 1999; Tozzi \\& Norman 2001; Rosati, Stefano, \\& Norman et al. 2002) typically observed for low-temperature X-ray systems. There have been quite a few suggestions that a non-gravitational heating (thus extra entropy) might prevent the ICM from further collapsing at the cluster core. The effect is supposed to be more pronounced in colder systems whose virial temperature is comparable to the temperature created by this non-gravitational heating, leading to shallower gas profiles than those of high-temperature systems (e.g., Balogh et al. 1999; Tozzi \\& Norman 2001). Interestingly, our determination of the surface brightness profile of Lynx-W is much shallower ($\\beta=0.42\\pm0.07$) than that of Lynx-E ($\\beta=0.71\\pm0.12$) (however, Ettori et al. (2004) obtained $\\beta=0.97\\pm0.43$ for Lynx-W).  The relatively loose distribution of the cluster galaxies in Lynx-W without any apparent BCG defining the cluster center leads us to consider another possibility that the system might be dynamically young and the ICM has not fully thermalized within the potential well. If we imagine that the ICM is not primordial, but has been ejected from the cluster galaxies at some recent epoch, it is plausible to expect that the X-ray temperature of the ICM might yet under-represent the depth of the cluster potential well. Tozzi et al. (2003) investigated the iron abundance in the ICM at $0.3<z<1.3$ and argued that the result was consistent with no evolution of the mean iron abundance out to $z\\simeq1.2$. If we assume that, as they suggested, Type Ia SNe are the dominant sources of this iron enrichment and have already injected their metals into the ICM by $z\\sim1.2$, a significant fraction of clusters at $z\\gtrsim 1.2$ may possess dynamically young ICM.  Recently, Nakata et al. (2005) reported with a photometric redshift technique the discovery of seven other cluster candidates around these two Lynx clusters possibly forming a $z\\sim1.3$ supercluster. Although further evidence is needed that the individual clumps are dynamically bound, the clear enhancement of the red galaxies consistent with the color at the redshift of the two known Lynx clusters is worthy of our attention. If they are indeed found to be forming groups/clusters at $z\\sim1.3$, but missed by X-ray observations because of their low X-ray contrast, the detailed studies of these young high-redshift structures will provide a critical benchmark in testing our understanding of the structure formation as well as the individual galaxy evolution in the context of different environments.  Deep two band ($i_{775}$ and $z_{850}$ $HST$/ACS imaging of the five out of the seven group/cluster candidates of Nakata el al. (2005) are scheduled in $HST$ Cycle 14 (Prop. 10574, PI. Mei). Studies similar to the current investigation will not only test whether there exist dark matter clumps around the candidate galaxies, but also quantify the environments for the investigation of the cluster galaxy color/morphology evolution.  ACS was developed under NASA contract NAS5-32865, and this research was supported by NASA grant NAG5-7697. We are grateful for an equipment grant from Sun Microsystems, Inc.  Some of the data presented herein were obtained at the W.M. Keck Observatory, which is operated as a scientific partnership among the California Institute of Technology, the University of California and the National Aeronautics and Space Administration. The Observatory was made possible by the generous financial support of the W.M. Keck Foundation.  The authors wish to recognize and acknowledge the very significant cultural role and reverence that the summit of Mauna Kea has always had within the indigenous Hawaiian community.  We are most fortunate to have the opportunity to conduct observations from this mountain."
    },
    "0601/astro-ph0601102_arXiv.txt": {
        "abstract": "We present spectroscopy and discuss the  photometric history of a previously obscure star in M31.  The spectrum of the star is an extremely close match to that of P~Cygni, one of the archetypes of  Luminous Blue Variables (LBVs). The star has not shown much variability over the past 40 years ($<0.2$~mag), although small-scale (0.05~mag) variations over a year appear to be real. Nevertheless, the presence of a sub-arcsecond extension around the star is indicative of a past outburst, and from the nebula's size (0.5~pc diameter) we estimate the outburst took place roughly 2000 yrs ago.  P Cygni itself exhibits a similar photometric behavior, and has a similar nebula (0.2~pc diameter).  We argue that this may be more typical behavior for LBVs than commonly assumed.  The star's location in the HR diagram offers substantial support for stellar evolutionary models that include the effects of rotation, as the star is just at a juncture in the evolutionary track of a 85$M_\\odot$ star.  The star is likely in a transition from an O star to a late-type WN Wolf-Rayet. ",
        "introduction": "\\label{Sec-intro}  Luminous Blue Variables (LBVs) are a rare class of luminous stars which undergo episodic mass-loss, and probably represent a transition between the most massive O stars and the red supergiant and/or Wolf-Rayet stage. (See the recent conference proceedings edited by Davidson et al.\\ 1988, Nota \\& Lamers 1997, and de Groot \\& Sterken 2001.) All of the ``accepted\" LBVs were discovered on the basis photometric and/or spectroscopic variability, but the number of LBV ``candidates\" (based upon spectroscopic similarities to known LBVs) exceeds this number (Parker 1997).  Understanding the number of bona-fide LBVs relative to that of other evolved massive stars (Wolf-Rayets, red supergiants) in nearby galaxies (i.e., at differing metallicities) provides a crucial test of the success of modern stellar evolutionary theory (Maeder \\& Meynet 2000).  Massey (2002) has argued that if the archetypal LBVs P~Cygni or $\\eta$ Carina were located in M~31 or M~33 we would not know of them today, since their spectacular photometric outbursts were hundreds of years ago.  In this {\\it Letter} we report the discovery of a star in M31 whose spectrum is uniquely similar to that of P Cygni itself.  A comparison of the star's location with stellar evolutionary tracks provides a strong substantiation of the latest generation of stellar models that include the effects of rotation (Meynet \\& Maeder 2005), while also offering a caution in interpreting evolutionary stages of these hydrostatic models.  The star has been relatively photometrically stable for the past 40 years, yet it shows compelling evidence of an outburst 2 millennia ago.  This provides further support for the argument large photometric variations on the time-scale of decades may be more a selection effect than a consequence of the physics of LBVs. ",
        "conclusions": "The LGGS project has obtained photometry of $>$350,000 stars seen towards M31.  One of the brightest of these, LGGS 004341.84+411112.0, turns out to have a spectrum that is remarkably similar to that of the P Cygni, one of the best studied LBVs in the Milky Way. The star has been relatively constant in $B$ over the past 40 years, with variations $<0.2$~mag, but with evidence of variations of 0.05~mag during a year.  Much the same can be said for P Cygni (Israelian \\& de Groot 1999). However, the star has likely had an outburst two millennia ago, as judged by the fact that the star is slightly extended {\\it HST} images, indicative of a 0.5~pc nebula, similar to the 0.2~pc nebula seen around P~Cygni.  This has consequences for interpreting the nature of objects found with spectroscopic similarity to other known LBVs; these stars are considered LBV ``candidates\" until variability is established, but the present study emphasizes that this may require more than a lifetime.  If we assume that the spectral similarity of our star to that of P Cygni is indicative of a similar effective temperature (and hence intrinsic color and bolometric correction), we can place the star on the H-R diagram using our photometry.  The star falls on the evolutionary track for a 85$M_\\odot$ star at the extreme end its evolution to cooler temperature, just  as the evolutionary tracks turns to lower luminosity (Meynet \\& Maeder 2005).  The star is likely in a transition between an O star and a WNL.  This provides strong vindication of the current generation of stellar evolutionary models that include rotation, but it also notes the difficulties in interpreting the evolution phases implied by the models. An atmospheric analysis of the star is planned, and long-term follow-up (both photometric and spectroscopic) is warranted."
    },
    "0601/astro-ph0601428_arXiv.txt": {
        "abstract": "Among the light elements created in the Big Bang, deuterium is one of the most difficult to detect but is also the one whose abundance depends most sensitively on the density of baryons. Thus, although we still have only a few positive identifications of D at high redshifts---when the D/H ratio was close to its primordial value---they give us the most reliable determination of the baryon density, in excellent agreement with measures obtained from entirely different probes, such as the anisotropy of the cosmic microwave background temperature and the average absorption of the UV light of quasars by the intergalactic medium. In this review, I shall relate observations of D/H in distant gas clouds to the large body of data on the local abundance of D obtained in the last few years with {\\it FUSE\\/}. I shall also discuss some of the outstanding problems in light element abundances and consider future prospects for advances in this area. ",
        "introduction": "The measurement of the interstellar abundance of deuterium was one of the main science drivers of the \\emph{FUSE} mission right from its inception. Five years on, this promise has been amply fulfilled, as demonstrated by the numerous talks and posters at this meeting devoted to \\emph{FUSE} results on D/H.  The importance of deuterium stems from the fact that, among the light elements created in Big Bang nucleosynthesis (BBNS), it is the one whose primordial abundance responds most sensitively to cosmological density of baryons, $\\Omega_{\\rm b}$. While $^4$He is the most abundant, because it soaks up essentially all the available neutrons, this property also makes it a rather insensitive `baryometer'. The quantity $\\Omega_{\\rm b}$, or more precisely the baryon to photon ratio $\\eta$, only affects $Y_{\\rm p}$ (the primordial mass fraction in  $^4$He) by determining the time delay before BBNS can set in---and thus the time available for neutrons to decay---in the first few minutes of the universe history. The more fragile deuterium, on the other hand, is easily destroyed by two-body reactions with protons, neutrons and other nuclei so that its abundance relative to hydrogen when BBNS is over, (D/H)$_0$ or D$_0$ for short, shows a steep, inverse, dependence on $\\Omega_{\\rm b}$. $^7$Li is less useful than D in this respect because it is far less abundant, by about five orders of magnitude, and its dependence on $\\Omega_{\\rm b}$ is double-valued because it can be synthesised via different nuclear reactions in the high and low baryon density regimes.  The detection of interstellar D was among the first discoveries made by \\emph{FUSE}'s predecessor, the \\emph{Copernicus} satellite. Rogerson \\& York (1973) resolved the isotope shift in the Ly$\\gamma$ line seen towards the bright B1 III star $\\beta$~Centauri, and deduced $N$(\\ion{D}{1})/$N$(\\ion{H}{1})$\\,= (1.4 \\pm 0.2) \\times 10^{-5}$. Three decades later, the mean of 21 measurements of D/H in the `Local Bubble' (the nearby region of the Milky Way disk) is in excellent agreement with \\emph{Copernicus}' first detection: $\\langle$D/H$\\rangle = (1.56 \\pm 0.04) \\times 10^{-5}$ (Wood et al. 2004).\\footnote{An interesting observation is that the distance to $\\beta$~Cen has `doubled' since 1973. The \\emph{Hipparcos} parallax to this star implies a distance $d = 161\\,$pc, while the parallactic distance available to Rogerson \\& York (1973) was $d = 81$\\,pc. This is a clear demonstration that it is easier for astronomers to measure chemical abundances than distances, even to the brightest stars.}  An important property of deuterium is that it is only destroyed whenever interstellar gas is cycled through stars (a process termed astration), so that its abundance relative to H steadily decreases with the progress of galactic chemical evolution. Analytically, this reduces to a simple expression for the time evolution of D: \\begin{equation} {\\rm (D/D_0)}~ = ~f^{(1/\\alpha~ -1)}~ = ~e^{-Z(1/\\alpha~ -1)} \\end{equation} where $f$ is the gas fraction, $Z$ the metallicity (in units of the yield of a primary element such as oxygen) and $\\alpha$ is the mass fraction which is locked up in long-lived stars and stellar remnants whenever a quantity $M$ of interstellar matter is turned into stars (Ostriker \\& Tinsley 1975; Pagel 1997). Equation (1) is valid in the simplest case of a `closed-box' model of chemical evolution. More realistic models which include inflow and/or outflow generally result in lower reductions of the primordial D/H as a function of either $f$ or $Z$ (Edmunds 1994). We cannot measure the lock-up fraction $\\alpha$ directly, but only deduce it theoretically by assuming a distribution of stellar masses (the IMF) and guessing at what fraction of its initial mass each star returns to the interstellar medium (ISM).  In principle, the degree of astration suffered by D in the Milky Way, where $f = 0.1 - 0.2$, could be anywhere between 20\\% and 90\\%, depending on the uncertain value of $\\alpha$. Consequently, Rogerson and York could only use their measurement of D/H in the ISM today to place a lower limit on (D/H)$_0$ and a corresponding upper limit $\\Omega_{\\rm b} \\leq 0.0675$ (for $h = 0.7$ where, as usual, $h$ is the Hubble constant in units of 100\\,km~s$^{-1}$~Mpc$^{-1}$. In this article I shall use $h = 0.7$ throughout and dispose of the $h^2$ term in the value of $\\Omega_{\\rm b}$). Note that the above upper limit on $\\Omega_{\\rm b}$, even without any correction for D astration, implies that most of the matter in the universe is non-baryonic. ",
        "conclusions": "We have come a long way since that pioneering measurement by Rogerson \\& York of the interstellar abundance of deuterium with {\\em Copernicus\\/} more than thirty years ago. The number of such measurements now approaches fifty, thanks in particular to the capabilities of {\\em FUSE\\/} and the GHRS and STIS instruments on the {\\em Hubble Space Telescope\\/}. With large ground-based telescopes we have been able to probe high redshift clouds where D is still close to its primordial abundance. New methods have been exploited to determine the density of baryons, the most impressive of which is the mapping of the temperature anisotropies in the cosmic microwave background over a wide range of angular scales. Bringing all of these developments together we find that many aspects of the overall picture fit together remarkably well, giving us confidence in the validity of the whole cosmological framework. Others still provide challenging puzzles, particularly the unexplained dispersion in the local determinations of D/H and the very low abundance of $^7$Li in some of the oldest stars of our Galaxy. But I am optimistic that we will not have to wait another three decades to iron out these remaining wrinkles in our understanding of the abundances of the light elements."
    },
    "0601/astro-ph0601664_arXiv.txt": {
        "abstract": "{We investigated abundance ratios along the profiles of six high-redshift Damped Lyman-$\\alpha$ systems, three of them associated with H$_2$ absorption, and derived optical depths in each velocity pixel. The variations of the pixel abundance ratios were found to be remarkably small and usually smaller than a factor of two within a profile. This result holds even when considering independent sub-clumps in the same system. The depletion factor is significantly enhanced only in those components where H$_2$ is detected. There is a strong correlation between [Fe/S] and [Si/S] abundances ratios, showing that the abundance ratio patterns are definitely related to the presence of dust. The depletion pattern is usually close to the one seen in the warm halo gas of our Galaxy. ",
        "introduction": "\\label{sec:introduction}  Damped Lyman-$\\alpha$ systems (hereafter DLAs) observed in QSO spectra are characterized by strong H~{\\sc i}$\\lambda$1215 absorption lines with broad damping wings. Although the definition has been restricted for historical reasons to absorptions with log~$N$(H~{\\sc i})~$>$~20.3 (Wolfe et al. 1986), damping wings are easily detected in present-day, high quality data for much lower column densities (down to log~$N$(H~{\\sc i})~$\\sim$~18.5). A more appropriate definition should be related to the physical state of the gas. If we impose the condition that the gas must be neutral, then the definition should be limited to systems with log $N$(H~{\\sc i})~$>$~19.5 (e.g. Viegas 1995).   Since their discovery twenty years ago \\citep{Wolfe86}, DLAs clearly have something to do with galaxy formation. What kind of galaxy DLAs are associated to is, however, still a matter of debate. Some authors identify these systems with large rotating discs \\citep{Prochaska97,Hou01}, while others think that DLAs arise mostly either in dwarf galaxies \\citep{Centurion} or galactic building blobs \\citep{Haehnelt98,Ledoux98}. The answer is probably not unique. In any case, DLAs represent the major reservoir of neutral hydrogen at any redshift \\citep{Storrie2000}, and they probe the chemical enrichment and evolution of the neutral Universe  (see Pettini et al. 1994; Lu et al. 1996; Prochaska et al. 1999; Ledoux et al. 2002a and references therein). Since abundances can be measured very accurately in DLAs, we can both discuss the connection between observed abundance ratios and dust content and to trace the nucleosynthesis history of the dense gas in the universe.  In this context, it is helpful to compare these results with measurements in the ISM of our Galaxy. Refractory elements that condense easily into dust grains - namely, Cr, Fe, Ni- are strongly depleted (up to a factor hundred) in the ISM, while non-refractory elements remain in its gaseous phase - S, Zn, and partially Si-. The amount of depletion depends on the physical condition of the gas. Thus, different depletion patterns are observed depending on whether the gas is cold or warm and/or whether the gas is located in the disc or the halo of the Galaxy \\citep{Savage96}. The LMC and SMC also exhibit different gas-phase abundance ratios \\citep{Welty99}.  However, a particular nucleosynthesis history can give rise to peculiar metallicity patterns and mimic the presence of dust. \\citet{Tinsley} suggested that type Ia supernova are the major producers of Fe. An enhancement in [$\\alpha$/Fe] ratios ($\\alpha$-elements are mostly O, S, Si) could reflect an IMF skewed to high masses and therefore a predominant role of type II supernova. For very low metallicity stars ([Fe/H]~$<$~$-$3) in the Galaxy,  large variations in several abundance ratios have been reported \\citep{william}, which suggests that peculiar nucleosynthesis processes and inhomogeneous chemical enrichment are probably taking place.  As mentioned above, DLAs trace the chemical evolution of galaxies at early epochs in the universe. Many  detailed studies have been performed so far, revealing that their metallicities range between 1/300 $Z_{\\odot}$ and solar values. The abundance pattern is fairly uniform and compatible with low dust content (see Pettini et al. 1994, Lu et al. 1996, Prochaska et al. 1999, Ledoux et al. 2002a). This uniformity in the relative abundance patterns observed from one DLA to the other has been emphasized by Prochaska \\& Wolfe (2002) and suggests that protogalaxies have common enrichment histories.  Few studies have adressed the question of the homogeneity inside each particular system. Prochaska \\& Wolfe (1996) first studied chemical abundance variations in a single DLA, and showed that the chemical abundances were uniform to within statistical uncertainties. Lopez et al. (2002) confirmed this finding from analysis of another DLA using Voigt profile decomposition. Petitjean et al. (2002) and Ledoux et al. (2002b) showed that the depletion patterns in subcomponents were very similar along DLA profiles except in the components where molecular hydrogen is detected and where depletion is larger. More recently, \\citet{Prochaska03}, performed a study of 13 systems concluding that the majority of DLAs have very uniform relative abundances. This contrasts in particular with the dispersion in nucleosynthetic enrichment of the Milky Way as traced by stellar abundances. Here, we use the best data from our survey of DLAs (Ledoux et al. 2003) to investigate  this issue further using an inversion method to derive the velocity profiles in different abundance ratios. In particular, we investigate the consequence of the presence of molecular hydrogen in some of the DLAs. The paper is structured  as follows: we describe the data in Sect. 2; in Sect. \\ref{sec:method} we briefly introduce the method used for the analysis; and results are presented and discussed in Sects. \\ref{sec:results} and \\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{Relative abundance ratios }\\label{subsec:histo}  \\begin{figure*} \\centering \\includegraphics[height=5.5cm,width=7.5cm]{2504f13a.ps} \\includegraphics[height=5.5cm,width=7.5cm]{2504f13b.ps} \\includegraphics[height=5.5cm,width=7.5cm]{2504f13c.ps} \\includegraphics[height=5.5cm,width=7.5cm]{2504f13d.ps} \\includegraphics[height=5.5cm,width=7.5cm]{2504f13e.ps} \\includegraphics[height=5.5cm,width=7.5cm]{2504f13f.ps} \\caption{Depletion of heavy elements relative to Sulfur (except for Q~1157+014, where Zinc was used instead) in the different subclumps of the six DLA systems studied here. Filled symbols represent the different sub-systems considered (see  Table~\\ref{tab:res}). Error bars correspond to typical scatter for each sub-system. The histograms show the observed values in the cold (solid line) or warm (dashed line) disc clouds and in halo clouds (dotted line) of the Galaxy.} \\label{fig:depletion_pattern} \\end{figure*}    \\begin{figure*} \\centering \\includegraphics[height=10.5cm,width=13.5cm]{2504f14.ps} \\caption{[Fe/S] vs [Si/S] for all the subclumps analyzed in this paper. Different symbols represent different DLA: squares for Q~0013$-$004, diamonds for Q~0528$-$250, triangles for Q~0405$-$443, and circles for Q~1037$-$270. Open symbols are used to distinguish subclumps  where H$_2$ is detected. Otherwise symbols are filled. We also indicate the typical [Fe/S] vs [Si/S] values observed in the cold, warm ISM, and halo of our Galaxy from \\citet{Welty99}.} \\label{fig:correl} \\end{figure*}    \\begin{figure*} \\centering \\includegraphics[height=5.5cm,width=8.0cm]{2504f15a.ps} \\includegraphics[height=5.5cm,width=8.0cm]{2504f15b.ps} \\caption{Summary of the inhomogeneity amplitude observed in the subsystems. The y-axis represents the maximum deviation of $[Fe/S]\\pm2\\sigma$ in each subclump. Subsystems with molecules are respresented with filled symbols.} \\label{fig:inhomog} \\end{figure*}  Relative abundance ratios for each sub-clump considered in the six systems analyzed in this work are given in  Table~\\ref{tab:res}. In the same way, results are summarized in Fig.~\\ref{fig:depletion_pattern}. It is apparent that, within each subsystem, large departures from the mean ratio are rare: the scatter is small if we take the observational and fitting uncertainties into account .  Moreover, when we compare depletion values from one sub-clump to another in the same system, differences are small. A distinct depletion pattern is observed only for some molecular components .This is remarkably summarized by Fig.~\\ref{fig:correl}, where we plot the [Fe/S] ratio versus the [Si/S] ratio in all subclumps. First, as already emphasized by Petitjean et al. (2002) and Ledoux et al. (2002b), the sequence seen in this figure is a dust-depletion sequence. Indeed, there is a correlation between the two quantities which is expected if the depletion is due to the presence of dust. Secondly, the values measured in different clumps of the same system are gathered at the same place in the figure. The only exception is the H$_2$ component at $-$480~km~s$^{-1}$ in Q~0013$-$004 (see above). Thirdly, most of the depletion pattern is similar to that of the gas observed in the Galactic halo. Finally, it seems that silicon is overabundant by about 0.2~dex even relative to sulfur. In all this, however, it must be recalled that we do not have access to the absolute metallicity in the subclumps, because we are not able to disentangle the H~{\\sc i} absorptions of the different subclumps.  In the left hand panel of Fig.~\\ref{fig:inhomog}, we plot the scatter, measured as $\\sigma$, of the ratio [Fe/S] in all the subclumps considered. In this case, we considered the subclumps as a whole, not isolating the molecular component, as our aim was to see if there is a relation between the inhomogeneity of a system and its molecular content. The mean value of $\\sigma$ over the subclumps is 0.3, which means that inhomogeneities are less than a factor of 2. Only a few subclumps where H$_2$ is detected have larger $\\sigma$. This is expected because we have seen that depletion is larger over the specific small velocity ranges over which H$_2$ is detected.  In the right hand panel of Fig.~\\ref{fig:inhomog}, we plot the different scatter values for each subclump as a function of the total [Fe/S] ratio. This figure confirms that (i) larger [Fe/S] ratios are observed in subclumps where H$_2$ is detected with one exception in a subclump of Q~0013$-$0004, and that (ii) the scatter is larger for subclumps where H$_2$ is detected.    \\subsection{The presence of \\hdos~ molecules}\\label{subsec:molecules}   \\cite{Ledouxsurvey} have systematically  searched for molecular hydrogen in high redshift DLAs, with a $\\sim$20\\% detection rate over the whole sample. The observed molecular fraction is often much smaller than in the ISM of the Galactic disc (Rachford et al. 2002) and is closer to what is observed in the magellanic clouds (Tumlinson et al. 2002). Here, we confirm what was already noticed by Ledoux et al. (2002b) and \\citet{ppjq0013} that, although the presence of molecules sometimes reveals  gas with larger depletion into dust grains than average, this is not always the case. In most of the systems, the depletion factor is only a factor of two larger in the components with H$_2$ compared to the overall system. There are a few exceptions, the most noticeable being the molecular component at $-$480~km~s$^{-1}$ toward Q~0013$-$014, in which depletion is as large as in the cold gas of the Galactic disc.   \\subsection{Consequences}\\label{subsec:overall} Variations in the relative abundance pattern within DLAs are expected from different nucleosynthesis histories and from its depletion onto dust. The magnitude of variation in the nucleosynthesis pattern may depend on the history of star formation and the level of enrichment. For example, stars in the Milky Way, Magellanic clouds, and local dwarf galaxies with metallicities varying from solar to 1/100 of solar show a dispersion in [Si/Fe] of the order of 0.3~dex (e.g. Shetrone et al. 2001, Venn et al. 2001).  Peculiar nucleosynthesis histories may be reflected in the variation of abundance ratios from one subclump to the other. Although we observed differences in the relative metal abundances of different sub-clumps, they are not large. This may indicate that sub-clumps in DLAs have the same origin and history and could be part of the same object. This  contrasts with the large differences in absolute metallicities that have been observed in LLS with similar velocity differences (e.g. D'Odorico \\& Petitjean 2001).  That depletion onto dust depends on the local physical conditions should induce a large scatter in the observed pixel-to-pixel relative abundance ratios. The fact that only small scatter was observed may reveal that the gas in DLAs is neither very dense nor cold but rather diffuse and warm. At least the filling factor of highly depleted gas is small.  All this implies uniform physical conditions and homogeneous and efficient mixing. One can speculate that this is only possible if DLAs are small objects with dimensions on the order of one kilo-parsec. This is difficult to ascertain as direct detection of high-redshift DLAs have not been very successful till now (e.g. Kulkarni et al. 2000, M\\o ller et al. 2004). It is, however, very important to pursue these observations in order to better constrain the nature and physical properties of these objects."
    },
    "0601/astro-ph0601387_arXiv.txt": {
        "abstract": "We present the on-going observational program of a VIMOS Integral Field Unit survey of the central regions of massive, gravitational lensing galaxy clusters at redshift $z \\simeq 0.2$. We have observed six clusters using the low-resolution blue grism ($R \\simeq 200$), and the spectroscopic survey is complemented by a wealth of photometric data, including {\\it Hubble Space Telescope} optical data  and near infrared VLT data. The principal scientific aims of this project are: the study of the high-$z$ lensed galaxies, the transformation and evolution of galaxies in cluster cores and the use of multiple images to constrain cosmography. We briefly report here on the first results from this project on the clusters Abell 2667 and Abell 68. ",
        "introduction": "\\label{sec:1}  Because of their intense gravitational field, massive (i.e., $M > 10^{14} M_{\\odot}$) clusters of galaxies act as Gravitational Telescopes (GTs) and are therefore an important tool to investigate the high-redshift Universe (see, e.g., \\cite{gt}). In order to fully exploit the scientific potential of the GTs we have started an extensive integral field spectroscopy (IFS) survey of massive galaxy clusters. Targets have been selected among well-known gravitational lensing clusters between redshift $\\sim 0.2$ and 0.3, for which complementary {\\it Hubble Space Telescope} (HST) data are available. All the clusters are X-ray bright sources. Our sample partially overlaps with the one analyzed in \\cite{smith05}.  IFS is the ideal tool in order to obtain spatially complete spectroscopic information of compact sky regions such as cluster cores. Moreover, cluster cores are the regions where strong lensing phenomena are observed (i.e., giant arcs and multiply imaged sources): for clusters in our sample, strongly lensed galaxies are within $\\theta \\simeq 1$ arcmin from the cluster center.  VIMOS-IFU \\cite{lefevre03} is thus the natural choice for such observational program, since, at present, it provides the largest f.o.v. for among integral field spectrographs mounted on the 8-10m telescopes.  All the clusters in our sample have been observed using the low-resolution blue (LR-B) grism, with a spectral resolution $R \\sim 200$. Taking into account the lower efficiency at the end of the spectra and the zero order contaminations, the final useful spectral range is limited between $\\simeq 3900$ \\AA \\, and 6800 \\AA. This spectral range is suitable both for detecting high-redshift source (e.g., Ly$\\alpha$ emitters in the redshift range $2.2 < z < 5.5$ or [OII] emitters out to $z\\sim 0.8$) and to sample the rest-frame 4000 \\AA \\, break for the cluster galaxy population. With a fiber size 0.66 arcsec, the IFU f.o.v. covers a contigous region of $54\\times 54$ arcsec$^2$, sampled by 6400 fibers.  A subset of the clusters in the sample has also been observed using a higher resolution grism ($R\\simeq3000$) covering the $\\lambda= 6300-8600 \\, $ \\AA\\ range. These observations are useful to probe higher redshift Ly$\\alpha$ emitters at $4.2<z<6.1$ or [OII] emitters at 0.7$<z<$1.3.  Observations have been completed (see Table 1 in \\cite{soucail06}), and data reduction is now in progress. The data reduction process has been described in \\cite{covone06}, and \\cite{zanichelli05} gives details about VIPGI, the VIMOS dedicated pipeline. Furthermore, we have developed a Sextractor-based tool to help in object detection and spectra extraction from the fully reduced datacube \\cite{jullo05}.  Altogether, our IFS survey covers a region of about 9 square arcmin in the central regions of six massive clusters.  Hereafter, we briefly report on the first results from this project: the mass distribution in Abell 2667 in Sect. \\ref{sec:2}, the properties of a magnified high-$z$ source in Abell 68 in Sect. \\ref{sec:3} and an investigation of cluster galaxies in Sect. \\ref{sec:4}. We refer to the presentation by G. Soucail \\cite{soucail06} for a detailed discussion of the cosmography aspect of the project. ",
        "conclusions": ""
    },
    "0601/astro-ph0601178_arXiv.txt": {
        "abstract": "\\noindent The presentation of new results from an \\oiii\\ survey in a search for planetary nebulae (PNe) in the galactic bulge is continued. A total of 60 objects, including 19 new PNe, have been detected in the remaining 34 per cent of the survey area, while 41 objects are already known. Deep \\ha $+$\\nitrogen~CCD images as well as low resolution spectra have been acquired for these objects. Their spectral signatures suggest that the detected emission originates from photoionized nebulae. In addition, absolute line fluxes have been measured and the electron densities are given. Accurate optical positions and optical diameters are  also determined. ",
        "introduction": "This is the second of a series of papers presenting optical imaging and spectrophotometric results on newly discovered planetary nebulae. In Paper I (\\citealt{Bo03}), the discovery method and first results on the \\oiii\\ survey, in the northern Galactic bulge, were presented.  PNe are powerful tracers of our Galaxy's star formation history (Paper I; \\citealt{Be00} and references therein). The study of PNe can provide insight to the late stages of stellar evolution, the nucleosynthesis in low and intermediate mass stars (1 M$_{\\odot}$ to 8 M$_{\\odot}$) and the chemical evolution of galaxies. When a low or intermediate mass star is closing to its end, it goes through the red--giant (RG) phase, followed by a heavy mass--loss period known as the asymptotic--giant--branch (AGB) stage. Finally, it becomes a white--dwarf (WD) surrounded by a planetary nebula (PN) made up of ionized gas from earlier stages of mass--loss.  PNe near the Galactic center can be considered to lie, roughly, at the same distance. Approximately 500 PNe have been discovered up to now, while $\\sim$ 3500 have been discovered in our Galaxy (Paper I and references therein; \\citealt{Pa03,Pa05a,Pa05b}; \\citealt{Ja04}). These numbers are considered small when compared with that generally expected (15000 to 30000; \\citealt{Ac92a}; \\citealt{Zi91}). Therefore, our Galaxy and especially, the Bulge area is well suited to search for new PNe.  Information concerning the \\oxygen~survey (observations and analysis) are given in sect. 2, while in sect. 3 the follow-up observations are presented (CCD imaging and spectroscopy) of the newly discovered PNe. In addition, flux measurements, diameter determination, accurate positions and other physical properties for the new PNe are given is sect. 4. In an forthcoming paper (Paper III -- in preparation), we will present photoionization modelling applied to the new PNe, in order to gain more insight into their physical parameters. In particular, flux calibrated images for most of the new PNe together with their spectroscopic results will be used in a PNe photoionization model (CLOUDY; e.g. \\citealt{Va99}). ",
        "conclusions": "We presented results from the narrow band \\oiii\\ survey of PNe in the Galactic bulge ($l>0^{o}$). Covering the remaining 34 percent of our selected region, we detected 60 objects, including 41 known PNe and 19 new PNe. \\ha $+$\\nitrogen\\ images as well as low resolution spectra of the new PNe were taken which confirmed the photoionized nature of the emission. About 84 percent of the detected objects have angular sizes $\\leq$ 20--25 arcsec, while all show \\ha/\\oiii\\ ratios less than 1. Four (4) of our new PNe are found to be associated with IRAS sources. Radio sources for fifteen (15) of our new PNe (including those of paper I) were identified. All new PNe display low N/O and He abundances implying old progenitor stars, which is one of the characteristics of the Galactic bulge PNe. In a forthcoming paper (paper III), the use of the photoionization model CLOYDY will provide a deeper insight to the physical parameters of the new PNe. The information presented in papers I and II, along with the results that will be obtained from the CLOUDY model will offer a valuable tool to studies of the dynamics and kinematics of the Galactic bulge.   \\begin{figure*} \\centering \\scalebox{0.90}{\\includegraphics{Fig1.eps}} \\caption[]{Optical imaging survey grid in equatorial coordinates. Galactic coordinates are also included (dash lines) to permit an accurate drawing of the selected Bulge region (bold solid lines). The dark filled and light filled rectangles represent the remaining observed field and the observed fields in the year 2000, respectively.} \\label{fig01} \\end{figure*}  \\begin{figure*} \\centering \\scalebox{0.90}{\\includegraphics{Fig2a.eps}} \\caption[]{\\ha $+$\\nitrogen~images of all new PNe taken with the 1.3 m telescope. The arrows point to their position. The images have a size of 150\\arcsec on both sides. North is at the top, East to the left.} \\label{fig02a} \\end{figure*}  \\begin{figure*} \\centering \\addtocounter{figure}{-1} \\scalebox{0.90}{\\includegraphics{Fig2b.eps}} \\caption[]{continued} \\label{fig02b} \\end{figure*}   \\begin{figure*} \\centering \\scalebox{0.90}{\\includegraphics{Fig3a.eps}} \\caption[]{Observed spectra of our new PNe taken with the 1.3 m telescope. They cover the range of 4750\\AA\\ to 6815\\AA\\ and the emission line fluxes (in units of $10^{-16}$ erg s$^{-1}$ cm$^{-2}$ arcsec$^{-2}$ \\AA$^{-1}$) are corrected for atmospheric extinction. Line fluxes corrected for interstellar extinction are given in Table 2.} \\label{fig03a} \\end{figure*}  \\begin{figure*} \\centering \\addtocounter{figure}{-1} \\scalebox{0.90}{\\includegraphics{Fig3b.eps}} \\caption[]{continued} \\label{fig03b} \\end{figure*}  \\begin{figure*} \\centering \\addtocounter{figure}{-1} \\scalebox{0.65}{\\includegraphics{Fig3c.eps}} \\caption[]{continued} \\label{fig03c} \\end{figure*}  \\begin{figure*} \\centering \\scalebox{0.65}{\\includegraphics{Fig4.eps}} \\caption[]{(a) IRAS colour--colour diagram of 13 of our objects (stars), overlaid on the colour--colour diagram of Acker's catalogued PNe with good (circles) and not good (rectangles) quality fluxes and (b) Diagnostic diagram (Garcia et al. 1991), where the positions of the new PNe are shown with a triangle ($\\bigtriangleup$). For comparison, the position of (i) the new PNe presented in paper I are also shown with a cross (X) and (ii) other objects (supernova remnants - SNR, H {\\sc ii} regions and Herbig Haro objects - HH) are shown, too.} \\label{fig04} \\end{figure*}  \\begin{figure*} \\centering \\scalebox{0.80}{\\includegraphics{Fig5.eps}} \\caption[]{(a) Enlargement of the \\ha $+$\\nitrogen~image of PTB34 in high contrast in order to show its outer halo (b) Image of the same PN overlaid with the different scale contours as derived from the method described in the text. North is in the top, east to the left in both figures.} \\label{fig05} \\end{figure*}"
    },
    "0601/astro-ph0601452_arXiv.txt": {
        "abstract": "{Deep X--ray surveys are providing crucial information on the evolution of AGN and galaxies.  We review some of the latest results based on the X--ray spectral analysis of the sources detected in the Chandra Deep Field South, namely: i) constraints on obscured accretion; ii) constraints on the missing fraction of the X--ray background; iii) the redshift distribution of Compton--thick sources and TypeII QSO; iv) the detection of star formation activity in high--z galaxies through stacking techniques; v) the detection of large scale structure in the AGN distribution and its effect on nuclear activity.  Such observational findings are consistent with a scenario where nuclear activity and star formation processes develop together in an anti--hierarchical fashion. ",
        "introduction": "In the last years deep X--ray surveys with the {\\sl Chandra} and {\\sl XMM--Newton} satellites (Brandt et al. 2001; Rosati et al. 2002; Alexander et al. 2003; Hasinger et al. 2001), paralleled by multiwavelength campaigns (see, e.g., GOODS, Giavalisco et al. 2004), provided several crucial information on the evolution of the AGN and galaxy populations.  The bold result from the two deepest X--ray fields, the Chandra Deep Field North (CDFN, observed for 2 Ms) and the Chandra Deep Field South (CDFS, observed for about 1Ms) is constituted by the resolution of the X--ray background (XRB) into single sources, mostly AGN, at a level between 80\\% and 90\\% (see Bauer et al. 2004 and the recently revised estimate by Hickox \\& Markevitch 2005), providing an almost complete census of the accretion history of matter onto supermassive black holes through the cosmic epochs.  However, the most interesting outcomes go well beyond the demographic characterization of the extragalactic X--ray sky.  Indeed, the physical and evolutionary properties of the AGN population are now revealing how they formed and how they are linked to their host galaxies.  For the first time, the luminosity function of AGN has been measured up to high redshift.  A striking feature is the {\\sl downsizing}, or {\\sl anti--hierarchical} behaviour, of the nuclear activity: the space density of the brightest Seyfert I and QSO is peaking at $z\\geq 2$, while the less luminous Seyfert II and I peak at $z\\leq 1$ (Ueda et al. 2003; Hasinger et al. 2005; La Franca et al. 2005).  An analogous behaviour is presently observed in the cosmic star formation history: at low redshift star formation is mostly observed in small objects (see, e.g., Kauffmann et al. 2004), while at redshift 2 or higher, star formation activity is observed also in massive galaxies (with $M_* \\sim 10^{11} M_\\odot$, see, Daddi et al. 2004a).  The global picture, as outlined by the present data, requires a tight link between the formation of the massive spheroids and the central black holes, as witnessed by the relation between black hole and stellar masses or between black hole mass and the velocity dispersion of the bulge (Kormendy \\& Richstone 1995; Magorrian et al. 1998; Ferrarese \\& Merritt 2000).  The anti--hierarchical behaviour in both star formation and AGN activity (which reflects in an anti--hierarchical supermassive black holes growth, see Merloni 2004; Marconi et al. 2004; for an alternative view see Hopkins et al. 2005), is envisaged by theoretical models where energy feedback is invoked to self--regulate both processes (see Fabian 1999; Granato et al. 2004).  In these Proceedings, we will describe a few observational results obtained from the latest analysis of the X--ray and optical data in the Chandra Deep Field South and North and which, in our view, are consistent with this picture.  In detail, these results concern the following issues:  \\begin{itemize} \\item the physical properties of AGN from the X--ray spectral analysis of faint X--ray sources;  \\item the missing fraction of the XRB;  \\item the distribution of obscured QSO and Compton--thick sources and their relation with the cosmic mass accretion history;  \\item star formation in high--z galaxies measured in the X--ray band thanks to stacking techniques;  \\item the effects of large scale structure onto nuclear activity.  \\end{itemize} ",
        "conclusions": "We presented few selected topics which we find particularly relevant among the latest results from deep X--ray surveys.  We can summarize our conclusions as follows:  \\begin{itemize}  \\item a population of strongly absorbed, possibly Compton--thick AGN at $z\\sim 1$ is still missing to the census of the X--ray sky; the detailed X--ray spectral analysis of faint sources shows that we are detecting some of them, and help us in obtaining a complete reconstruction of the cosmic accretion history onto supermassive black holes;  \\item we find several absorbed sources (the so--called TypeII QSO) among the population of bright AGN, possibily witnessing the rapid growth of the super massive black holes associated to strong star formation events;  \\item thanks to stacking techniques, we detected the X--ray emission associated to massive star forming galaxies at redshift as high as $z\\sim 2$, therefore peering with X--rays in the epoch of massive galaxy formation;  \\item investigation of large scale structure in the X--ray detected AGN distribution provides tantalizing hints of its effect on nuclear activity.  Studies of spatial correlation of X--ray sources require larger fields of view to kill the cosmic variance and therefore evaluate properly the evolution of the AGN clustering properties.  \\end{itemize}  Such observational findings are providing crucial information on the evolution of AGN and galaxies.  Present--day data are consistent with a scenario where nuclear activity and star formation processes develop together in an anti--hierarchical fashion."
    },
    "0601/astro-ph0601208_arXiv.txt": {
        "abstract": "\\begin{center} \\end{center}  Born-Infeld electromagnetic waves interacting with a static magnetic background are studied in an expanding universe. The non-linear character of Born-Infeld electrodynamics modifies the relation between the energy flux and the distance to the source, which gains a new dependence on the redshift that is governed by the background field. We compute the luminosity distance as a function of the redshift and compare with Maxwellian curves for supernovae type Ia. ",
        "introduction": "The discovery of an unexpected diminution in the observed energy fluxes coming from supernovae type Ia \\cite{Riess-Perlmutter}, which are thought of as standard candles, has been interpreted in the context of the standard cosmological model as evidence for an accelerating universe dominated by something called dark energy. This fact is one of the most puzzling and deepest problems in cosmology and fundamental physics today. Although the cosmological constant seems to be the simplest explanation for the phenomenon, several dynamical scenarios have been tried out (see, for instance, \\cite{padmanabhan} and references therein). It is worthwhile to emphasize that the evidence for an accelerating universe mainly relies on energy flux measurements for type Ia supernovae at different values of cosmological redshifts. They provide the most direct and consistent way to determine the recent expansion history of the universe. Nevertheless the relation between the cosmological redshift and the energy flux for a point-like source involves not only the evolution of the universe during the light journey but some assumptions about the nature of the light itself. Customarily, one accepts the linear Maxwell theory to describe the light propagation, where light propagates without interacting with other electric or magnetic fields. However, in the context of non-linear electrodynamics the interaction between the light emitted from such distant sources and cosmological magnetic backgrounds modifies the relation between the redshift and the flux of energy. If this kind of effect were not correctly interpreted it could lead to an erroneous conclusion about the expansion history of the universe. Concretely, an effect coming from non-linear electrodynamics could explain the curves of luminosity distance vs redshift for type Ia supernovae without invoking dark energy. This remark drives us to study the propagation of non-linear electromagnetic waves in an expanding universe with a magnetic background. We will benefit from recently obtained results for Born-Infeld electromagnetic plane waves propagating in a magnetic uniform background in Minkowski space-time \\cite{9abf}.  Born-Infeld electrodynamics \\cite{1born,2born} is a non-linear theory for the electromagnetic field $F_{\\mu\\nu}$ governed by the Lagrangian \\begin{equation} {\\cal L}=\\sqrt{-g}\\ \\frac{b^2}{4\\:\\pi}\\,\\left(1-\\sqrt{1+\\frac{2S}{b^2}- \\frac{P^2}{b^4}}\\right) \\end{equation} where $b$ is a new fundamental constant and $S$ and $P$ are the scalar and pseudoscalar field invariants \\begin{equation} S=\\frac{1}{4}F_{\\mu\\nu}F^{\\mu\\nu}\\ \\ \\ \\ \\ \\ \\ \\ \\ \\ P=\\frac{1}{4}\\:^{*}F_{\\mu\\nu}F^{\\mu\\nu} \\end{equation} (in Minkowski space-time it is $2S=|{\\bf{\\mathbb B}}|^2-|{\\bf{\\mathbb E}}|^2$ and $P= {\\bf{\\mathbb E}}\\cdot{\\bf{\\mathbb B}}$, $\\bf{\\mathbb E}$ and $\\bf{\\mathbb B}$ being the electric and magnetic fields respectively). Born-Infeld electrodynamics goes to Maxwell electromagnetism when $b\\rightarrow \\infty$. In particular, the Born-Infeld field of a point-like charge reaches the finite value $b$ at the charge position and becomes Coulombian far from the charge. Born-Infeld theory is the only non-linear spin 1 field theory displaying causal propagation and absence of birefringence \\cite{deser,5boi}.  Nowadays, Born-Infeld theory is reborn in the context of superstrings because Born-Infeld-like Lagrangians emerge in the low energy limit of string theories \\cite{strings}. Born-Infeld-like Lagrangians have been also proposed to describe a matter dynamics able to drive the universe to an accelerated expansion \\cite{varios}. In spite of this revival of Born-Infeld's ideas, there is no experimental evidence for Born-Infeld effects in electrodynamics: the value of $b$ remains unknown (see an upper bound for $b$ in \\cite{Jack}).  In the next section we will summarize the recently obtained results on Born-Infeld waves propagating in Minkowski space-time in the presence of a uniform magnetic background \\cite{9abf}. These results --properly adapted-- will be used in section III to understand the Born-Infeld energy flux coming from a point-like source in a spatially flat Friedman-Robertson-Walker (FRW) expanding universe. In section IV we will reformulate the relation between the luminosity distance $d_L$ and the redshift $z$ within the framework of Born-Infeld electrodynamics. We will show that the presence of magnetic backgrounds modify the curves $d_L$ vs $z$, and we will analyze the consequences for the measurements of the luminosity distance of supernovae type Ia. ",
        "conclusions": "In this paper we have solved the Born-Infeld equations for electromagnetic plane waves propagating in a background magnetic field. In the absence of a background field, the Born-Infeld plane waves are equal to the Maxwell ones. On the contrary, in the presence of a background magnetic field ${\\bf B}$ the non-linear effects modify both the phase and the amplitude of the wave with corrections that depend on the combination $\\vert{\\bf B}\\vert^2\\, a^{-4}\\, b^{-2}$, where $a$ is the scale factor of the universe. It is remarkable that Born-Infeld electrodynamics depends on $a$ and $b$ only through the combination $a^4 b^2$. This means that the Maxwellian approximation ($b\\rightarrow\\infty$) also corresponds to the limit $a\\rightarrow\\infty$. So, although the electromagnetic field is presently well described by Maxwell equations for a wide range of phenomena, the non-linear Born-Infeld electrodynamics could have an influence in the past when the scale factor was smaller. Therefore the expanding universe is a good laboratory to test Born-Infeld electrodynamics; many non-linear aspects of its equations could be relevant when highly redshifted objects are observed.  In this work we have begun the search for this kind of effects. We found that the influence of Born-Infeld electrodynamics on the luminosity distance (\\ref{Hod5}) exhibits interesting features that could be experimentally established by means of more precise supernova observations and a better knowledge of the cosmological background fields. Firstly, the experimental data for $d_L$ vs $z$ could be fitted without invoking dark energy, although there is no observational evidence of the background field that would be required. Secondly, the shape of the curve $d_L$ vs $z$ predicted by the standard cosmology ($\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$, $b\\rightarrow \\infty$) for high redshifts differs appreciably from the one predicted by Born-Infeld electrodynamics, which opens the possibility of detecting non-linear electrodynamics effects in a future."
    },
    "0601/astro-ph0601514_arXiv.txt": {
        "abstract": "Name{ABSTRACT}% \\def\\abstract{% \\endmode \\bigskip\\bigskip \\centerline{\\AbstractName}% \\medskip \\bgroup \\let\\endmode=\\endabstract \\narrower\\narrower \\singlespaced \\everyabstract}% \\def\\everyabstract{}% \\def\\endabstract{\\smallskip\\egroup} \\def\\pacs#1{\\medskip\\centerline{PACS numbers: #1}\\smallskip} \\def\\submit#1{\\bigskip\\centerline{Submitted to {\\sl #1}}} \\def\\submitted#1{\\submit{#1}}% \\def\\toappear#1{\\bigskip\\raggedcenter To appear in {\\sl #1} \\endraggedcenter} \\def\\disclaimer#1{\\footnote{}\\bgroup\\tenrm\\singlespaced This manuscript has been authored under contract number #1 \\@disclaimer\\par} \\def\\disclaimers#1{\\footnote{}\\bgroup\\tenrm\\singlespaced This manuscript has been authored under contract numbers #1 \\@disclaimer\\par} \\def\\@disclaimer{% with the U.S. Department of Energy.  Accordingly, the U.S. Government retains a non-exclusive, royalty-free license to publish or reproduce the published form of this contribution, or allow others to do so, for U.S. Government purposes. \\egroup}  \\catcode`@=11 \\chardef\\other=12 \\def\\center{% \\flushenv \\advance\\leftskip \\z@ plus 1fil \\advance\\rightskip \\z@ plus 1fil \\obeylines\\@eatpar}% \\def\\flushright{% \\flushenv \\advance\\leftskip \\z@ plus 1fil \\obeylines\\@eatpar}% \\def\\flushleft{% \\flushenv \\advance\\rightskip \\z@ plus 1fil \\obeylines\\@eatpar}% \\def\\flushenv{% \\vskip \\z@ \\bgroup \\def\\flushhmode{F}% \\parindent=\\z@  \\parfillskip=\\z@}% \\def\\endcenter{\\endflushenv} \\def\\endflushleft{\\endflushenv} \\def\\endflushright{\\endflushenv} \\def\\@eatpar{\\futurelet\\next\\@testpar} \\def\\@testpar{\\ifx\\next\\par\\let\\@next=\\@@eatpar\\else\\let\\@next=\\relax\\fi\\@next} \\long\\def\\@@eatpar#1{\\relax} \\def\\raggedcenter{% \\flushenv \\advance\\leftskip\\z@ plus4em \\advance\\rightskip\\z@ plus 4em \\spaceskip=.3333em \\xspaceskip=.5em \\pretolerance=9999 \\tolerance=9999 \\hyphenpenalty=9999 \\exhyphenpenalty=9999 \\@eatpar}% \\def\\endraggedcenter{\\endflushenv}% \\def\\hcenter{\\hflushenv \\advance\\leftskip \\z@ plus 1fil \\advance\\rightskip \\z@ plus 1fil \\obeylines\\@eatpar}% \\def\\hflushright{\\hflushenv \\advance\\leftskip \\z@ plus 1fil \\obeylines\\@eatpar}% \\def\\hflushleft{\\hflushenv \\advance\\rightskip \\z@ plus 1fil \\obeylines\\@eatpar}% \\def\\hflushenv{% \\def\\par{\\endgraf\\indent}% \\hbox to \\z@ \\bgroup\\hss\\vtop \\flushenv\\def\\flushhmode{T}}% \\def\\endflushenv{% \\ifhmode\\endgraf\\fi \\if T\\flushhmode \\egroup\\hss\\fi \\egroup}% \\def\\flushhmode{U} \\def\\endhcenter{\\endflushenv} \\def\\endhflushleft{\\endflushenv} \\def\\endhflushright{\\endflushenv} \\newskip\\EnvTopskip     \\EnvTopskip=\\medskipamount \\newskip\\EnvBottomskip  \\EnvBottomskip=\\medskipamount \\newskip\\EnvLeftskip    \\EnvLeftskip=2\\parindent \\newskip\\EnvRightskip   \\EnvRightskip=\\parindent \\newskip\\EnvDelt@skip   \\EnvDelt@skip=0pt \\newcount\\@envDepth     \\@envDepth=\\z@ \\def\\beginEnv#1{% \\begingroup \\def\\@envname{#1}% \\ifvmode\\def\\@isVmode{T}% \\else\\def\\@isVmode{F}\\vskip 0pt\\fi \\ifnum\\@envDepth=\\@ne\\parindent=\\z@\\fi \\advance\\@envDepth by \\@ne \\EnvDelt@skip=\\baselineskip \\advance\\EnvDelt@skip by-\\normalbaselineskip \\@setenvmargins\\EnvLeftskip\\EnvRightskip \\setenvskip{\\EnvTopskip}% \\vskip\\skip@\\penalty-500 } \\def\\endEnv#1{% \\ifnum\\@envDepth<1 \\emsg{> Tried to close ``#1'' environment, but no environment open!}% \\begingroup \\else \\def\\test{#1}% \\ifx\\test\\@envname\\else \\emsg{> Miss-matched environments!}% \\emsg{> Should be closing ``\\@envname'' instead of ``\\test''}% \\fi \\fi \\vskip 0pt \\setenvskip\\EnvBottomskip \\vskip\\skip@\\penalty-500 \\xdef\\@envtemp{\\@isVmode}% \\endgroup \\if F\\@envtemp\\vskip-\\parskip\\par\\noindent\\fi } \\def\\setenvskip#1{\\skip@=#1 \\divide\\skip@ by \\@envDepth} \\def\\@setenvmargins#1#2{% \\advance \\leftskip  by #1    \\advance \\displaywidth by -#1 \\advance \\rightskip by #2    \\advance \\displaywidth by -#2 \\advance \\displayindent by #1}% \\def\\itemize{\\beginEnv{itemize}% \\let\\itm=\\itemizeitem \\vskip-\\parskip } \\def\\itemizeitem{% \\par\\noindent \\hbox to 0pt{\\hss\\itemmark\\space}}% \\def\\enditemize{\\endEnv{itemize}}% \\def\\itemmark{$\\bullet$}% \\newcount\\enumDepth     \\enumDepth=\\z@ \\newcount\\enumcnt \\def\\enumerate{\\beginEnv{enumerate}% \\global\\advance\\enumDepth by \\@ne \\setenumlead \\enumcnt=\\z@ \\let\\itm=\\enumerateitem \\if F\\@isVmode\\vskip-\\parskip\\fi } \\def\\enumerateitem{% \\par\\noindent \\advance\\enumcnt by \\@ne \\edef\\lab@l{\\enumlead \\enumcur}% \\hbox to \\z@{\\hss \\lab@l \\enummark \\hskip .5em\\relax}% \\ignorespaces}% \\def\\endenumerate{% \\global\\advance\\enumDepth by -\\@ne \\endEnv{enumerate}}% \\def\\enumPoints{% \\def\\setenumlead{\\ifnum\\enumDepth>1 \\edef\\enumlead{\\enumlead\\enumcur.}% \\else\\def\\enumlead{}\\fi}% \\def\\enumcur{\\number\\enumcnt}% } \\def\\enumpoints{\\enumPoints}% \\def\\enumOutline{% \\def\\setenumlead{\\def\\enumlead{}}% \\def\\enumcur{\\ifcase\\enumDepth \\or\\uppercase{\\XA\\romannumeral\\number\\enumcnt}% \\or\\LetterN{\\the\\enumcnt}% \\or\\XA\\romannumeral\\number\\enumcnt \\or\\letterN{\\the\\enumcnt}% \\or{\\the\\enumcnt}% \\else $\\bullet$\\space\\fi}% } \\def\\enumoutline{\\enumOutline}% \\def\\enumNumOutline{% \\def\\setenumlead{\\def\\enumlead{}}% \\def\\enumcur{\\ifcase\\enumDepth \\or{\\XA\\number\\enumcnt}% \\or\\letterN{\\the\\enumcnt}% \\or{\\XA\\romannumeral\\number\\enumcnt}% \\else $\\bullet$\\space\\fi}% } \\def\\enumnumoutline{\\enumNumOutline}% \\def\\LetterN#1{\\count@=#1 \\advance\\count@ 64 \\XA\\char\\count@} \\def\\letterN#1{\\count@=#1 \\advance\\count@ 96 \\XA\\char\\count@} \\def\\enummark{.}% \\def\\enumlead{}% \\enumpoints \\newbox\\@desbox \\newbox\\@desline \\newdimen\\@glodeswd \\newcount\\@deslines \\newif\\ifsingleline \\singlelinefalse \\def\\description#1{\\beginEnv{description}% \\setbox\\@desbox=\\hbox{#1}% \\@glodeswd=\\wd\\@desbox \\@setenvmargins{\\@glodeswd}{0pt}% \\let\\itm=\\descriptionitem \\if F\\@isVmode\\vskip-\\parskip\\fi }% \\def\\descriptionitem#1{% \\goodbreak\\noindent \\setbox\\@desline=\\vtop\\bgroup \\hfuzz=100cm\\hsize=\\@glodeswd \\rightskip=\\z@ \\leftskip=\\z@ \\raggedright \\noindent{#1}\\par \\global\\@deslines=\\prevgraf \\egroup \\ifsingleline \\ifnum\\@deslines>1 \\@deslineitm{#1}% \\else \\setbox\\@desline=\\hbox{#1}% \\ifdim \\wd\\@desline>\\wd\\@desbox \\@deslineitm{#1}% \\else\\@desitm\\fi \\fi \\else \\@desitm \\fi \\ignorespaces} \\def\\@desitm{% \\noindent \\hbox to \\z@{\\hskip-\\@glodeswd \\hbox to \\@glodeswd{\\vtop to \\z@{\\box\\@desline\\vss}% \\hss}\\hss}}% \\def\\@deslineitm#1{% \\hbox{\\hskip-\\@glodeswd {#1}\\hss}% \\vskip-\\parskip\\nobreak\\noindent } \\def\\enddescription{\\ifhmode\\par\\fi \\@setenvmargins{-\\wd\\@desbox}{0pt}% \\endEnv{description}} \\def\\example{\\beginEnv{example}% \\parskip=\\z@ \\parindent=\\z@ \\baselineskip=\\normalbaselineskip }% \\def\\endexample{\\endEnv{example}% \\noindent}% \\let\\blockquote=\\example \\let\\endblockquote=\\endexample \\def\\Listing{% \\beginEnv{Listing}% \\vskip\\EnvDelt@skip \\baselineskip=\\normalbaselineskip \\parskip=\\z@ \\parindent=\\z@ \\def\\\\##1{\\char92##1}% \\catcode`\\{=\\other \\catcode`\\}=\\other \\catcode`\\(=\\other \\catcode`\\)=\\other \\catcode`\\\"=\\other \\catcode`\\|=\\other \\catcode`\\%=\\other \\catcode`\\&=\\other \\catcode`\\-=\\other \\catcode`\\==\\other \\catcode`\\$=\\other \\catcode`\\#=\\other \\catcode`\\_=\\other \\catcode`\\^=\\other \\catcode`\\~=\\other \\obeylines \\tt\\Listingtabs \\everyListing}% \\def\\endListing{\\endEnv{Listing}}% \\def\\everyListing{\\relax} \\def\\ListCodeFile#1{% \\Listing \\rightskip=\\z@ plus 5cm \\catcode`\\\\=\\other \\input #1\\relax \\endListing} {\\catcode`\\^^I=\\active\\catcode`\\ =\\active \\gdef\\Listingtabs{\\catcode`\\^^I=\\active\\let^^I\\@listingtab \\catcode`\\ =\\active\\let \\@listingspace}% }% \\def\\@listingspace{\\hskip 0.5em\\relax}% \\def\\@listingtab{\\hskip 4em\\relax}% \\def\\TeXexample{\\beginEnv{TeXexample}% \\vskip\\EnvDelt@skip \\parskip=\\z@ \\parindent=\\z@ \\baselineskip=\\normalbaselineskip \\def\\par{\\leavevmode\\endgraf}% \\obeylines \\catcode`|=\\z@ \\ttverbatim \\@eatpar}% \\def\\endTeXexample{% \\vskip 0pt \\endgroup \\endEnv{TeXexample}}% \\def\\ttverbatim{\\begingroup \\catcode`\\(=\\other \\catcode`\\)=\\other \\catcode`\\\"=\\other \\catcode`\\[=\\other \\catcode`\\]=\\other \\catcode`\\~=\\other \\let\\do=\\uncatcode \\dospecials \\obeyspaces\\obeylines \\def\\n{\\vskip\\baselineskip}% \\tt}% \\def\\uncatcode#1{\\catcode`#1=\\other}% {\\obeyspaces\\gdef {\\ }}% \\def\\TeXquoteon{\\catcode`\\|=\\active}% \\let\\TeXquoteson=\\TeXquoteon \\def\\TeXquoteoff{\\catcode`\\|=\\other}% \\let\\TeXquotesoff=\\TeXquoteoff {\\TeXquoteon\\obeylines% \\gdef|{\\ifmmode\\vert\\else% \\ttverbatim\\spaceskip=\\ttglue% \\let^^M=\\ \\relax% \\let|=\\endgroup\\fi}% } \\def\\ttvert{\\hbox{\\tt\\char`\\|}} \\outer\\def\\begintt{$$\\let\\par=\\endgraf \\ttverbatim \\parskip=0pt \\catcode`\\|=0 \\rightskip=-5pc \\ttfinish} {\\catcode`\\|=0 |catcode`|\\=\\other% |obeylines% |gdef|ttfinish#1^^M#2\\endtt{#1|vbox{#2}|endgroup$$}% } \\def\\beginlines{\\par\\begingroup\\nobreak\\medskip\\parindent=0pt \\hrule\\kern1pt\\nobreak \\obeylines \\everypar{\\strut}} \\def\\endlines{\\kern1pt\\hrule\\endgroup\\medbreak\\noindent} \\def\\beginproclaim#1#2#3#4#5{\\medbreak\\vskip-\\parskip \\global\\XA\\advance\\csname #2\\endcsname by \\@ne \\edef\\lab@l{\\@chaptID\\@sectID \\number\\csname #2\\endcsname}% \\tag{#4#5}{\\lab@l}% \\noindent{\\bf #1 \\lab@l.\\space}% \\begingroup #3}% \\def\\endproclaim{% \\par\\endgroup\\ifdim\\lastskip<\\medskipamount \\removelastskip\\penalty55\\medskip\\fi}% \\newcount\\theoremnum           \\theoremnum=\\z@ \\def\\theorem#1{\\beginproclaim{Theorem}{theoremnum}{\\sl}{Thm.}{#1}} \\let\\endtheorem=\\endproclaim \\def\\Theorem#1{Theorem~\\use{Thm.#1}} \\newcount\\lemmanum             \\lemmanum=\\z@ \\def\\lemma#1{\\beginproclaim{Lemma}{lemmanum}{\\sl}{Lem.}{#1}} \\let\\endlemma=\\endproclaim \\def\\Lemma#1{Lemma~\\use{Lem.#1}} \\newcount\\corollarynum         \\corollarynum=\\z@ \\def\\corollary#1{\\beginproclaim{Corollary}{corollarynum}{\\sl}{Cor.}{#1}} \\let\\endcorollary=\\endproclaim \\def\\Corollary#1{Corollary~\\use{Cor.#1}} \\newcount\\definitionnum        \\definitionnum=\\z@ \\def\\definition#1{\\beginproclaim{Definition}{definitionnum}{\\rm}{Def.}{#1}} \\let\\enddefinition=\\endproclaim \\def\\Definition#1{Definition~\\use{Def.#1}} \\def\\proof{\\medbreak\\vskip-\\parskip\\noindent{\\it Proof. }} \\def\\blackslug{% \\setbox0\\hbox{(}% \\vrule width.5em height\\ht0 depth\\dp0}% \\def\\QED{\\blackslug}% \\def\\endproof{\\quad\\blackslug\\par\\medskip}  \\catcode`@=11 \\def\\paper{% \\auxswitchtrue \\refswitchtrue \\texsis \\def\\titlepage{% \\bgroup \\let\\title=\\Title \\let\\endmode=\\relax \\pageno=1}% \\def\\endtitlepage{% \\endmode \\goodbreak\\bigskip \\egroup}% \\autoparens \\quoteon } \\def\\Tbf{\\fourteenpoint\\bf}% \\def\\tbf{\\twelvepoint\\bf}% \\def\\preprint{% \\auxswitchtrue \\refswitchtrue \\texsis \\def\\titlepage{% \\bgroup \\pageno=1 \\let\\title=\\Title \\let\\endmode=\\relax \\banner}% \\def\\endtitlepage{% \\endmode \\vfil\\eject \\egroup}% \\autoparens \\quoteon } \\def\\Manuscript{% \\preprint \\showsectIDfalse \\showchaptIDfalse \\def{% \\endmode \\bigskip\\bigskip\\medskip \\bgroup\\singlespaced \\let\\endmode=\\endabstract \\narrower\\narrower \\everyabstract}% \\def\\endabstract{% \\medskip\\egroup\\bigskip}% \\def\\FootText{--\\ \\tenrm\\folio\\ --}% \\def\\Tbf{\\sixteenpoint\\bf}% \\def\\tbf{\\fourteenpoint\\bf}% \\twelvepoint \\doublespaced \\autoparens \\quoteon }% \\autoload\\thesis{thesis.txs} \\autoload\\UTthesis{thesis.txs} \\autoload\\YaleThesis{thesis.txs} \\def\\Letter{% \\ContentsSwitchfalse \\refswitchfalse \\auxswitchfalse \\texsis \\singlespaced \\LetterFormat}% \\def\\letter{\\Letter}% \\def\\Memo{% \\ContentsSwitchfalse \\refswitchfalse \\auxswitchfalse \\texsis \\singlespaced \\MemoFormat}% \\def\\memo{\\Memo}% \\def\\Referee{% \\ContentsSwitchfalse \\auxswitchfalse \\refswitchfalse \\texsis \\RefReptFormat}% \\def\\referee{\\Referee}% \\def\\Landscape{% \\texsis \\hsize=9in \\vsize=6.5in \\voffset=.5in \\nopagenumbers \\LandscapeSpecial } \\def\\landscape{\\Landscape}% \\def\\LandscapeSpecial{\\special{papersize=11in,8.5in}} \\def\\slides{% \\quoteon \\autoparens \\ATlock \\pageno=1 \\twentyfourpoint \\doublespaced \\raggedright\\tolerance=2000 \\hyphenpenalty=500 \\raggedbottom \\nopagenumbers \\hoffset=-.25in \\hsize=7.0in \\voffset=-.25in \\vsize=9.0in \\parindent=30pt \\def\\bl{\\vskip\\normalbaselineskip}% \\def\\np{\\vfill\\eject}% \\def\\nospace{\\nulldelimiterspace=0pt \\mathsurround=0pt}% \\def\\big##1{{\\hbox{$\\left##1 \\vbox to2ex{}\\right.\\nospace$}}}% \\def\\Big##1{{\\hbox{$\\left##1 \\vbox to2.5ex{}\\right.\\nospace$}}}% \\def\\bigg##1{{\\hbox{$\\left##1 \\vbox to3ex{}\\right.\\nospace$}}}% \\def\\Bigg##1{{\\hbox{$\\left##1 \\vbox to4ex{}\\right.\\nospace$}}}% } \\autoload\\twinout{twin.txs}% \\def\\twinprint{% \\preprint \\let\\t@tl@=\\title \\def\\title{\\vskip-1.5in\\t@tl@}% \\let\\endt@tlep@ge=\\endtitlepage \\def\\endtitlepage{\\endt@tlep@ge \\twinformat}% } \\def\\twinformat{% \\tenpoint\\doublespaced \\def\\Tbf{\\twelvebf}\\def\\tbf{\\tenbf}% \\headlineoffset=0pt \\twinout}%  \\catcode`\\@=11 \\let\\NX=\\noexpand\\let\\XA=\\expandafter \\offparens \\newcount\\tabnum        \\tabnum=\\z@ \\newcount\\fignum        \\fignum=\\z@ \\newif\\ifRomanTables    \\RomanTablesfalse \\newif\\ifCaptionList    \\CaptionListfalse \\newif\\ifFigsLast       \\FigsLastfalse \\newif\\ifTabsLast       \\TabsLastfalse \\def\\FiguresLast{\\FigsLasttrue}\\def\\FiguresNow{\\FigsLastfalse} \\def\\TablesLast{\\TabsLasttrue}\\def\\TablesNow{\\TabsLastfalse} \\long\\def\\figure{\\@figure\\topinsert} \\long\\def\\topfigure{\\@figure\\topinsert}% \\long\\def\\midfigure{\\@figure\\midinsert} \\long\\def\\fullfigure{\\@figure\\pageinsert} \\long\\def\\bottomfigure{\\@figure\\bottominsert} \\long\\def\\heavyfigure{\\@figure\\heavyinsert} \\long\\def\\widefigure{\\@figure\\widetopinsert} \\long\\def\\widetopfigure{\\@figure\\widetopinsert} \\long\\def\\widefullfigure{\\@figure\\widepageinsert} \\def\\FigureName{Figure}% \\def\\TableName{Table}% \\def\\@figure#1#2{% \\vskip 0pt \\begingroup \\def\\CaptionName{\\FigureName}% \\def\\@prefix{Fg.}% \\let\\@count=\\fignum \\let\\@FigInsert=#1\\relax \\def\\@arg{#2}\\ifx\\@arg\\empty\\def\\@ID{}% \\else\\LabelParse #2;;\\endlist\\fi \\ifFigsLast \\emsg{\\CaptionName\\space\\@ID. {#2} [storing in \\jobname.fg]}% \\@fgwrite{\\@comment> \\CaptionName\\space\\@ID.\\space{#2}}% \\@fgwrite{\\string\\@FigureItem{\\CaptionName}{\\@ID}{\\NX#1}}% \\seeCR\\let\\@next=\\@copyfig \\else \\emsg{\\CaptionName\\ \\@ID.\\ {#2}}% \\let\\endfigure=\\@endfigure \\setbox\\@capbox\\vbox to 0pt{}% \\def\\@whereCap{N}% \\let\\@next=\\@findcap \\ifx\\@FigInsert\\midinsert\\goodbreak\\fi \\@FigInsert \\fi \\@next} \\def\\@endfigure{\\relax \\if B\\@whereCap\\relax \\vskip\\normalbaselineskip \\centerline{\\box\\@capbox}% \\fi \\endinsert \\endgroup}% \\def\\endfigure{\\emsg{> \\string\\endfigure before \\string\\figure!}} \\def\\figuresize#1{\\vglue #1}% \\def\\@copyfig#1#2\\endfigure{\\endgroup \\ifx#1\\par\\@fgNXwrite{#2\\@endfigure}\\else\\@fgNXwrite{#1#2\\@endfigure}\\fi} \\def\\@FGinit{\\@FileInit\\fgout=\\jobname.fg[Figures]\\gdef\\@FGinit{\\relax}} \\def\\@fgwrite#1{\\@FGinit\\immediate\\write\\fgout{#1}} \\long\\def\\@fgNXwrite#1{\\@FGinit\\unexpandedwrite\\fgout{#1}} \\def\\PrintFigures{\\ifFigsLast\\@PrintFigures\\fi} \\def\\@PrintFigures{% \\@fgwrite{\\@comment>>> EOF \\jobname.fg <<<}% \\immediate\\closeout\\fgout \\begingroup \\FigsLastfalse \\vbox to 0pt{\\hbox to 0pt{\\ \\hss}\\vss}% \\offparens \\catcode`@=11 \\emsg{[Getting figures from file \\jobname.fg]}% \\Input\\jobname.fg \\relax \\endgroup}% \\def\\@FigureItem#1#2#3{% \\begingroup #3\\relax \\def\\@ID{#2}% \\def\\CaptionName{#1}% \\setbox\\@capbox\\vbox to 0pt{}\\def\\@whereCap{N}% \\@findcap}% \\long\\def\\table{\\@table\\topinsert} \\long\\def\\toptable{\\@table\\topinsert}% \\long\\def\\midtable{\\@table\\midinsert} \\long\\def\\fulltable{\\@table\\pageinsert} \\long\\def\\bottomtable{\\@table\\bottominsert} \\long\\def\\heavytable{\\@table\\heavyinsert} \\long\\def\\widetable{\\@table\\widetopinsert} \\long\\def\\widetoptable{\\@table\\widetopinsert} \\long\\def\\widefulltable{\\@table\\widepageinsert} \\def\\@table#1#2{% \\vskip 0pt \\begingroup \\def\\CaptionName{\\TableName}% \\def\\@prefix{Tb.}% \\let\\@count=\\tabnum \\let\\@FigInsert=#1\\relax \\def\\@arg{#2}\\ifx\\@arg\\empty\\def\\@ID{}% \\else\\ifRomanTables \\global\\advance\\@count by\\@ne \\edef\\@ID{\\uppercase\\expandafter {\\romannumeral\\the\\@count}}% \\tag{\\@prefix#2}{\\@ID}% \\else \\LabelParse #2;;\\endlist\\fi \\fi \\ifTabsLast \\emsg{\\CaptionName\\space\\@ID. {#2} [storing in \\jobname.tb]}% \\@tbwrite{\\@comment> \\CaptionName\\space\\@ID.\\space{#2}}% \\@tbwrite{\\string\\@FigureItem{\\CaptionName}{\\@ID}{\\NX#1}}% \\seeCR\\let\\@next=\\@copytab \\else \\emsg{\\CaptionName\\ \\@ID.\\ {#2}}% \\let\\endtable=\\@endfigure \\setbox\\@capbox\\vbox to 0pt{}% \\def\\@whereCap{N}% \\let\\@next=\\@findcap \\ifx\\@FigInsert\\midinsert\\goodbreak\\fi \\@FigInsert \\fi \\@next} \\def\\endtable{\\emsg{> \\string\\endtable before \\string\\table!}} \\def\\@copytab#1#2\\endtable{\\endgroup \\ifx#1\\par\\@tbNXwrite{#2\\@endfigure}\\else\\@tbNXwrite{#1#2\\@endfigure}\\fi} \\def\\@TBinit{\\@FileInit\\tbout=\\jobname.tb[Tables]\\gdef\\@TBinit{\\relax}} \\def\\@tbwrite#1{\\@TBinit\\immediate\\write\\tbout{#1}} \\long\\def\\@tbNXwrite#1{\\@TBinit\\unexpandedwrite\\tbout{#1}} \\def\\PrintTables{\\ifTabsLast\\@PrintTables\\fi} \\def\\@PrintTables{% \\@tbwrite{\\@comment>>> EOF \\jobname.tb <<<}% \\immediate\\closeout\\tbout \\begingroup \\TabsLastfalse \\catcode`@=11 \\offparens \\emsg{[Getting tables from file \\jobname.tb]}% \\Input\\jobname.tb \\relax \\endgroup}% \\newbox\\@capbox \\newcount\\@caplines \\def\\CaptionName{}% \\def\\@ID{}% \\def\\captionspacing{\\normalbaselines}% \\def\\@inCaption{F}% \\long\\def\\caption#1{% \\def\\lab@l{\\@ID}% \\global\\setbox\\@capbox=\\vbox\\bgroup \\def\\@inCaption{T}% \\captionspacing\\seeCR \\dimen@=20\\parindent \\ifdim\\colwidth>\\dimen@\\narrower\\narrower\\fi \\noindent{\\bf \\linkname{\\@TagName}{\\CaptionName~\\@ID}:\\space}% #1\\relax \\vskip 0pt \\global\\@caplines=\\prevgraf \\egroup \\ifnum\\@caplines=\\@ne \\global\\setbox\\@capbox=\\vbox{\\noindent\\seeCR \\hfil{\\bf \\linkname{\\@TagName}{\\CaptionName~\\@ID}:\\space}% #1\\hfil}\\fi \\if N\\@whereCap\\def\\@whereCap{B}\\fi \\if T\\@whereCap \\centerline{\\box\\@capbox}% \\vskip\\baselineskip\\medskip \\fi}% \\def\\Caption{\\begingroup\\seeCR\\@Caption}% \\long\\def\\@Caption#1\\endCaption{\\endgroup \\ifCaptionList \\incaplist{#1}\\fi \\caption{#1}}% \\def\\endCaption{\\emsg{> \\string\\endCaption\\ called before \\string\\Caption.}} \\long\\def\\@findcap#1{% \\ifx #1\\Caption \\def\\@whereCap{T}\\fi \\ifx #1\\caption \\def\\@whereCap{T}\\fi #1}% \\def\\@whereCap{N}% \\def\\ListCaptions{\\@ListCaps\\caplist=\\jobname.cap[List of Captions] {\\let\\FIGLitem=\\CAPLitem}} \\def\\ListFigureCaptions{% \\@ListCaps\\figlist=\\jobname.fgl[List of Figure Captions] {\\let\\FIGLitem=\\CAPLitem}} \\def\\ListTableCaptions{% \\@ListCaps\\tablelist=\\jobname.tbl[List of Table Captions] {\\let\\FIGLitem=\\CAPLitem}} \\def\\CAPLitem#1#2#3\\@endFIGLitem#4{% \\bigskip \\begingroup \\raggedright\\tolerance=1700 \\hangindent=1.41\\parindent\\hangafter=1 \\noindent #1\\ #2 #3 \\hskip 0pt plus 10pt \\vskip 0pt \\endgroup}% \\def\\infiglist{\\begingroup\\seeCR \\@infiglist\\figlist} \\def\\intablelist{\\begingroup\\seeCR \\@infiglist\\tablelist} \\def\\incaplist{\\begingroup\\seeCR \\@infiglist\\caplist} \\def\\FigListWrite#1#2{% \\ifx#1\\figlist\\relax   \\FigListInit\\fi \\ifx#1\\tablelist\\relax \\TabListInit\\fi \\ifx#1\\caplist\\relax   \\CapListInit\\fi \\edef\\@line@{{#2}}% \\write#1\\@line@}% \\def\\FigListInit{\\@FileInit\\figlist=\\jobname.fgl[List of Figures]% \\gdef\\FigListInit{\\relax}} \\def\\TabListInit{\\@FileInit\\tablelist=\\jobname.tbl[List of Tables]% \\gdef\\TabListInit{\\relax}} \\def\\CapListInit{\\@FileInit\\caplist=\\jobname.cap[List of Captions]% \\gdef\\CapListInit{\\relax}} \\def\\FigListWriteNX#1#2{\\writeNX#1{#2}} \\def\\@infiglist#1#2{% \\FigListWrite#1{\\@comment > \\CaptionName\\space\\@ID:}% \\FigListWrite#1{\\string\\FIGLitem{\\CaptionName} {\\@ID.\\space}}% \\@copycap#1#2\\endlist \\FigListWrite#1{{\\NX\\folio}}% \\endgroup}% \\def\\@copycap#1#2#3\\endlist{% \\ifx#2\\space\\writeNX#1{#3\\@endFIGLitem}% \\else\\writeNX#1{#2#3\\@endFIGLitem}\\fi} \\def\\ListFigures{\\@ListCaps\\figlist=\\jobname.fgl[List of Figures]{}} \\def\\ListTables{\\@ListCaps\\tablelist=\\jobname.tbl[List of Tables]{}} \\def\\@ListCaps#1=#2[#3]#4{% \\immediate\\closeout#1 \\openin#1=#2 \\relax \\ifeof#1\\closein#1 \\else\\closein#1\\emsg{[Getting #3]}% \\begingroup \\showsectIDtrue \\ATunlock\\quoteoff\\offparens #4 \\input #2 \\relax \\endgroup \\fi} \\long\\def\\FIGLitem#1#2#3\\@endFIGLitem#4{% \\medskip \\begingroup \\raggedright\\tolerance=1700 \\ifx\\TOCmargin\\undefined\\skip0=\\parindent \\else\\skip0=\\TOCmargin\\fi \\advance\\rightskip by \\skip0 \\parfillskip=-\\skip0 \\hangindent=1.41\\parindent\\hangafter=1 \\noindent \\ifshowsectID #1\\ \\fi #2 #3 \\hskip 0pt plus 10pt \\leaddots \\hbox to 2em{\\hss\\linkto{page.#4}{#4}}% \\vskip 0pt \\endgroup} \\def\\Fig#1{\\linkto{Fg.#1}{Fig.~\\use{Fg.#1}}} \\def\\Figs#1{\\linkto{Fg.#1}{Figs.~\\use{Fg.#1}}} \\def\\Fg#1{\\linkto{Fg.#1}{\\use{Fg.#1}}} \\def\\Tab#1{\\linkto{Tb.#1}{Table~\\use{Tb.#1}}} \\def\\Tbl#1{\\linkto{Tb.#1}{Table~\\use{Tb.#1}}} \\def\\Tb#1{\\linkto{Tb.#1}{\\use{Tb.#1}}} \\autoload\\Tablebody{Tablebod.txs}\\autoload\\Tablebodyleft{Tablebod.txs} \\autoload\\tablebody{Tablebod.txs} \\autoload\\epsffile{epsf.tex}    \\autoload\\epsfbox{epsf.tex} \\autoload\\epsfxsize{epsf.tex}   \\autoload\\epsfysize{epsf.tex} \\autoload\\epsfverbosetrue{epsf.tex}\\autoload\\epsfverbosefalse{epsf.tex} \\obsolete\\topFigure\\figure \\obsolete\\midFigure\\midfigure \\obsolete\\fullFigure\\fullfigure \\obsolete\\TOPFIGURE\\figure \\obsolete\\MIDFIGURE\\midfigure \\obsolete\\FULLFIGURE\\fullfigure \\obsolete\\endFigure\\endfigure \\obsolete\\ENDFIGURE\\endfigure \\obsolete\\topTable\\toptable \\obsolete\\midTable\\midtable \\obsolete\\fullTable\\fulltable \\obsolete\\TOPTABLE\\toptable \\obsolete\\MIDTABLE\\midtable \\obsolete\\FULLTABLE\\fulltable \\obsolete\\endTable\\endtable \\obsolete\\ENDTABLE\\endtable \\def\\FIG{\\@obsolete\\FIG\\Fig\\Fig}% \\def\\TBL{\\@obsolete\\TBL\\Tbl\\Tbl}%  \\catcode`@=11 \\catcode`\\|=12 \\catcode`\\&=4 \\newcount\\ncols         \\ncols=\\z@ \\newcount\\nrows         \\nrows=\\z@ \\newcount\\curcol        \\curcol=\\z@ \\let\\currow=\\nrows \\newdimen\\thinsize      \\thinsize=0.6pt \\newdimen\\thicksize     \\thicksize=1.5pt \\newdimen\\tablewidth    \\tablewidth=-\\maxdimen \\newdimen\\parasize      \\parasize=4in \\newif\\iftableinfo      \\tableinfotrue \\newif\\ifcentertables   \\centertablestrue \\def\\centeredtables{\\centertablestrue}% \\def\\noncenteredtables{\\centertablesfalse}% \\def\\nocenteredtables{\\centertablesfalse}% \\let\\plaincr=\\cr \\let\\plainspan=\\span \\let\\plaintab=& \\def\\ampersand{\\char`\\&}% \\let\\lparen=( \\let\\NX=\\noexpand \\def\\ruledtable{\\relax \\@BeginRuledTable \\@RuledTable}% \\def\\@BeginRuledTable{% \\ncols=0\\nrows=0 \\begingroup \\offinterlineskip \\def~{\\phantom{0}}% \\def\\span{\\plainspan\\omit\\relax\\colcount\\plainspan}% \\let\\cr=\\crrule \\let\\CR=\\crthick \\let\\nr=\\crnorule \\let\\|=\\Vb \\def\\hfill{\\hskip0pt plus1fill\\hbox{}}% \\ifx\\tablestrut\\undefined\\relax \\else\\let\\tstrut=\\tablestrut\\fi \\catcode`\\|=13 \\catcode`\\&=13\\relax \\TableActive \\curcol=1 \\ifdim\\tablewidth>-\\maxdimen\\relax \\edef\\@Halign{\\NX\\halign to \\NX\\tablewidth\\NX\\bgroup\\TablePreamble}% \\tabskip=0pt plus 1fil \\else \\edef\\@Halign{\\NX\\halign\\NX\\bgroup\\TablePreamble}% \\tabskip=0pt \\fi \\ifcentertables \\ifhmode\\vskip 0pt\\fi \\line\\bgroup\\hss \\else\\hbox\\bgroup \\fi}% \\long\\def\\@RuledTable#1\\endruledtable{% \\vrule width\\thicksize \\vbox{\\@Halign \\thickrule #1\\killspace \\tstrut \\linecount \\plaincr\\thickrule \\egroup}% \\vrule width\\thicksize \\ifcentertables\\hss\\fi\\egroup \\endgroup \\global\\tablewidth=-\\maxdimen \\iftableinfo \\immediate\\write16{[Nrows=\\the\\nrows, Ncols=\\the\\ncols]}% \\fi}% \\def\\TablePreamble{% \\TableItem{####}% \\plaintab\\plaintab \\TableItem{####}% \\plaincr}% \\def\\@TableItem#1{% \\hfil\\tablespace #1\\killspace \\tablespace\\hfil }% \\def\\@tableright#1{% \\hfil\\tablespace\\relax #1\\killspace \\tablespace\\relax}% \\def\\@tableleft#1{% \\tablespace\\relax #1\\killspace \\tablespace\\hfil}% \\let\\TableItem=\\@TableItem \\def\\RightJustifyTables{\\let\\TableItem=\\@tableright}% \\def\\LeftJustifyTables{\\let\\TableItem=\\@tableleft}% \\def\\NoJustifyTables{\\let\\TableItem=\\@TableItem}% \\def\\LooseTables{\\let\\tablespace=\\quad}% \\def\\TightTables{\\let\\tablespace=\\space}% \\LooseTables \\def\\TrailingSpaces{\\let\\killspace=\\relax}% \\def\\NoTrailingSpaces{\\let\\killspace=\\unskip}% \\TrailingSpaces \\def\\setRuledStrut{% \\dimen@=\\baselineskip \\advance\\dimen@ by-\\normalbaselineskip \\ifdim\\dimen@<.5ex \\dimen@=.5ex\\fi \\setbox0=\\hbox{\\lparen}% \\dimen1=\\dimen@ \\advance\\dimen1 by \\ht0% \\dimen2=\\dimen@ \\advance\\dimen2 by \\dp0% \\def\\tstrut{\\vrule height\\dimen1 depth\\dimen2 width\\z@}% }% \\def\\tstrut{\\vrule height 3.1ex depth 1.2ex width 0pt}% \\def\\bigitem#1{% \\dimen@=\\baselineskip \\advance\\dimen@ by-\\normalbaselineskip \\ifdim\\dimen@<.5ex \\dimen@=.5ex\\fi \\setbox0=\\hbox{#1}% \\dimen1=\\dimen@ \\advance\\dimen1 by \\ht0 \\dimen2=\\dimen@ \\advance\\dimen2 by \\dp0 \\vrule height\\dimen1 depth\\dimen2 width\\z@ \\copy0}% \\def\\vctr#1{\\hfil\\vbox to 0pt{\\vss\\hbox{#1}\\vss}\\hfil}% \\def\\nextcolumn#1{% \\plaintab\\omit#1\\relax\\colcount \\plaintab}% \\def\\tab{% \\nextcolumn{\\relax}}% \\let\\novb=\\tab \\def\\vb{% \\nextcolumn{\\vrule width\\thinsize}}% \\def\\Vb{% \\nextcolumn{\\vrule width\\thicksize}}% \\def\\dbl{% \\nextcolumn{\\vrule width\\thinsize \\hskip 2\\thinsize \\vrule width\\thinsize}}% {\\catcode`\\|=13 \\let|0 \\catcode`\\&=13 \\let&0 \\gdef\\TableActive{\\let|=\\vb \\let&=\\tab}% }% \\def\\crrule{\\killspace \\tstrut \\linecount \\plaincr\\tablerule }% \\def\\crthick{\\killspace \\tstrut \\linecount \\plaincr\\thickrule }% \\def\\crnorule{\\killspace \\tstrut \\linecount \\plaincr }% \\def\\crpart{\\killspace \\linecount \\plaincr}% \\def\\tablerule{\\noalign{\\hrule height\\thinsize depth 0pt}}% \\def\\thickrule{\\noalign{\\hrule height\\thicksize depth 0pt}}% \\def\\cskip{\\omit\\relax}% \\def\\crule{\\omit\\leaders\\hrule height\\thinsize depth0pt\\hfill}% \\def\\Crule{\\omit\\leaders\\hrule height\\thicksize depth0pt\\hfill}% \\def\\linecount{% \\global\\advance\\nrows by1 \\ifnum\\ncols>0 \\ifnum\\curcol=\\ncols\\relax\\else \\immediate\\write16 {\\NX\\ruledtable warning: Ncols=\\the\\curcol\\space for Nrow=\\the\\nrows}% \\fi\\fi \\global\\ncols=\\curcol \\global\\curcol=1}% \\def\\colcount{\\relax \\global\\advance\\curcol by 1\\relax}% \\long\\def\\para#1{% \\vtop{\\hsize=\\parasize \\normalbaselines \\noindent #1\\relax \\vrule width 0pt depth 1.1ex}% }% \\def\\begintable{\\relax \\@BeginRuledTable \\@begintable}% \\long\\def\\@begintable#1\\endtable{% \\@RuledTable#1\\endruledtable}%  \\def\\E#1{\\hbox{$\\times 10^{#1}$}} \\def\\square{\\hbox{{$\\sqcup$}\\llap{$\\sqcap$}}}% \\def\\grad{\\nabla}% \\def\\del{\\partial}% \\def\\frac#1#2{{#1\\over#2}} \\def\\smallfrac#1#2{{\\scriptstyle {#1 \\over #2}}} \\def\\half{\\ifinner {\\scriptstyle {1 \\over 2}}% \\else {\\textstyle {1 \\over 2}}\\fi} \\def\\bra#1{\\langle#1\\vert}% \\def\\ket#1{\\vert#1\\/\\rangle}% \\def\\vev#1{\\langle{#1}\\rangle}% \\def\\simge{% \\mathrel{\\rlap{\\raise 0.511ex \\hbox{$>$}}{\\lower 0.511ex \\hbox{$\\sim$}}}} \\def\\simle{% \\mathrel{\\rlap{\\raise 0.511ex \\hbox{$<$}}{\\lower 0.511ex \\hbox{$\\sim$}}}} \\def\\gtsim{\\simge}% \\def\\ltsim{\\simle}% \\def\\therefore{% \\setbox0=\\hbox{$.\\kern.2em.$}\\dimen0=\\wd0 \\mathrel{\\rlap{\\raise.25ex\\hbox to\\dimen0{\\hfil$\\cdotp$\\hfil}}% \\copy0}} \\def\\|{\\ifmmode\\Vert\\else \\char`\\|\\fi} \\def\\sterling{{\\hbox{\\it\\char'44}}} \\def\\degrees{\\hbox{$^\\circ$}}% \\def\\degree{\\degrees}% \\def\\real{\\mathop{\\rm Re}\\nolimits}% \\def\\imag{\\mathop{\\rm Im}\\nolimits}% \\def\\tr{\\mathop{\\rm tr}\\nolimits}% \\def\\Tr{\\mathop{\\rm Tr}\\nolimits}% \\def\\Det{\\mathop{\\rm Det}\\nolimits}% \\def\\mod{\\mathop{\\rm mod}\\nolimits}% \\def\\wrt{\\mathop{\\rm wrt}\\nolimits}% \\def\\diag{\\mathop{\\rm diag}\\nolimits}% \\def\\TeV{{\\rm TeV}}% \\def\\GeV{{\\rm GeV}}% \\def\\MeV{{\\rm MeV}}% \\def\\keV{{\\rm keV}}% \\def\\eV{{\\rm eV}}% \\def\\Ry{{\\rm Ry}}% \\def\\mb{{\\rm mb}}% \\def\\mub{\\hbox{\\rm $\\mu$b}}% \\def\\nb{{\\rm nb}}% \\def\\pb{{\\rm pb}}% \\def\\fb{{\\rm fb}}% \\def\\cmsec{{\\rm cm^{-2}s^{-1}}}% \\def\\units#1{\\hbox{\\rm #1}} \\let\\unit=\\units \\def\\dimensions#1#2{\\hbox{$[\\hbox{\\rm #1}]^{#2}$}} \\def\\parenbar#1{{\\null\\! \\mathop{\\smash#1}\\limits ^{\\hbox{\\fiverm(--)}}% \\!\\null}}% \\def\\nunubar{\\parenbar{\\nu}} \\def\\ppbar{\\parenbar{p}} \\def\\buildchar#1#2#3{{\\null\\! \\mathop{\\vphantom{#1}\\smash#1}\\limits ^{#2}_{#3}% \\!\\null}}% \\def\\overcirc#1{\\buildchar{#1}{\\circ}{}} \\def\\sun{{\\hbox{$\\odot$}}}\\def\\earth{{\\hbox{$\\oplus$}}} \\def\\slashchar#1{\\setbox0=\\hbox{$#1$}% \\dimen0=\\wd0 \\setbox1=\\hbox{/} \\dimen1=\\wd1 \\ifdim\\dimen0>\\dimen1 \\rlap{\\hbox to \\dimen0{\\hfil/\\hfil}}% #1 \\else \\rlap{\\hbox to \\dimen1{\\hfil$#1$\\hfil}}% / \\fi}% \\def\\subrightarrow#1{% \\setbox0=\\hbox{% $\\displaystyle\\mathop{}% \\limits_{#1}$}% \\dimen0=\\wd0 \\advance \\dimen0 by .5em \\mathrel{% \\mathop{\\hbox to \\dimen0{\\rightarrowfill}}% \\limits_{#1}}}% \\newdimen\\vbigd@men \\def\\vbigl{\\mathopen\\vbig} \\def\\vbigm{\\mathrel\\vbig} \\def\\vbigr{\\mathclose\\vbig} \\def\\vbig#1#2{{\\vbigd@men=#2\\divide\\vbigd@men by 2 \\hbox{$\\left#1\\vbox to \\vbigd@men{}\\right.\\n@space$}}} \\def\\Leftcases#1{\\smash{\\vbigl\\{{#1}}} \\def\\Rightcases#1{\\smash{\\vbigr\\}{#1}}} \\def\\doublecolumns{\\relax}\\def\\enddoublecolumns{\\relax} \\def\\leftcolrule{\\relax}\\def\\rightcolrule{\\relax} \\def\\longequation{\\relax}\\def\\endlongequation{\\relax} \\def\\newcolumn{\\relax} \\def\\widetopinsert{\\topinsert}\\def\\widepageinsert{\\pageinsert} \\def\\forceleft{\\relax}\\def\\forceright{\\relax} \\def\\SetDoubleColumns#1{% \\imsg{The double column macros are not a part of mTeXsis.} \\imsg{If you want to use double column mode, get TXSdcol.tex} \\imsg{and add \\string\\input\\space TXSdcol.tex to your .tex file.} }   \\def\\addTOC#1#2#3{\\relax}\\def\\Contents{\\relax}  % \\newif\\ifContents                               % \\def\\ContentsSwitchtrue{\\Contentstrue}\\def\\ContentsSwitchfalse{\\Contentsfalse}  \\def\\obsolete#1#2{\\let#1=#2\\relax #2}           %  \\let\\Input=\\input                               % \\newdimen\\colwidth      \\colwidth=\\hsize        % \\def\\ORGANIZATION{}%   \\newhelp\\@utohelp{% loadstyle: The definition of the macro named above is actually contained^^J% in a style file, and so it cannot be used with mTeXsis.  If you really^^J% need to load the definition from that file, you should do so explicitly^^J% at the begining of your manuscript file, with % '\\string\\input\\space stylefilename.txs'^^J}  \\Ignore \\def\\loadstyle#1#2{% \\newlinechar=10                              % \\errhelp=\\@utohelp                           % \\emsg{> Whoops! Trying to load \\string#1\\space from style file #2.}% \\errmessage{You cannot use macro definitions from style files in mTeXsis}} \\endIgnore   \\hbadness=10000         % \\overfullrule=0pt       % \\vbadness=10000         %   \\ATunlock \\SetDate                                % \\ReadAUX                                % \\def\\fmtname{TeXsis}\\def\\fmtversion{2.17}% \\def\\revdate{1 January 1998}% \\def\\imsg#1{\\emsg{\\@comment #1}}% \\imsg{=========================================================== \\@comment} \\imsg{This is mTeXsis, the core macros from TeXsis.} \\imsg{You can get the complete TeXsis package (and avoid this annoying} \\imsg{advertisement) from ftp://lifshitz.ph.utexas.edu/texsis, } \\imsg{or from a CTAN server near you (in macros/texsis).} \\imsg{See the README and INSTALL files there for more information.} \\imsg{============================================================ \\@comment} \\emsg{m\\fmtname\\space version \\fmtversion\\space (\\revdate)  loaded.}% \\ATlock                                 % \\texsis                                 % \\quoteoff \\global\\mathchardef\\DELTA=\"7001 \\global\\mathchardef\\LAMBDA=\"7003 \\global\\mathchardef\\XI=\"7004 \\global\\mathchardef\\SIGMA=\"7006 \\global\\mathchardef\\UPSILON=\"7007 \\global\\mathchardef\\OMEGA=\"700A \\global\\mathchardef\\Delta=\"7101 \\global\\mathchardef\\Lambda=\"7103 \\global\\mathchardef\\Xi=\"7104 \\global\\mathchardef\\Sigma=\"7106 \\global\\mathchardef\\Upsilon=\"7107 \\global\\mathchardef\\Omega=\"710A   \\catcode`\\@=11  \\font\\tenmsa=msam10 \\font\\sevenmsa=msam7 \\font\\fivemsa=msam5 \\font\\tenmsb=msbm10 \\font\\sevenmsb=msbm7 \\font\\fivemsb=msbm5 \\newfam\\msafam \\newfam\\msbfam \\textfont\\msafam=\\tenmsa  \\scriptfont\\msafam=\\sevenmsa \\scriptscriptfont\\msafam=\\fivemsa \\textfont\\msbfam=\\tenmsb  \\scriptfont\\msbfam=\\sevenmsb \\scriptscriptfont\\msbfam=\\fivemsb  \\def\\hexnumber@#1{\\ifnum#1<10 \\number#1\\else \\ifnum#1=10 A\\else\\ifnum#1=11 B\\else\\ifnum#1=12 C\\else \\ifnum#1=13 D\\else\\ifnum#1=14 E\\else\\ifnum#1=15 F\\fi\\fi\\fi\\fi\\fi\\fi\\fi}  \\def\\msa@{\\hexnumber@\\msafam} \\def\\msb@{\\hexnumber@\\msbfam} \\global\\mathchardef\\boxdot=\"2\\msa@00 \\global\\mathchardef\\boxplus=\"2\\msa@01 \\global\\mathchardef\\boxtimes=\"2\\msa@02 \\global\\mathchardef\\square=\"0\\msa@03 \\global\\mathchardef\\blacksquare=\"0\\msa@04 \\global\\mathchardef\\centerdot=\"2\\msa@05 \\global\\mathchardef\\lozenge=\"0\\msa@06 \\global\\mathchardef\\blacklozenge=\"0\\msa@07 \\global\\mathchardef\\circlearrowright=\"3\\msa@08 \\global\\mathchardef\\circlearrowleft=\"3\\msa@09 \\global\\mathchardef\\rightleftharpoons=\"3\\msa@0A \\global\\mathchardef\\leftrightharpoons=\"3\\msa@0B \\global\\mathchardef\\boxminus=\"2\\msa@0C \\global\\mathchardef\\Vdash=\"3\\msa@0D \\global\\mathchardef\\Vvdash=\"3\\msa@0E \\global\\mathchardef\\vDash=\"3\\msa@0F \\global\\mathchardef\\twoheadrightarrow=\"3\\msa@10 \\global\\mathchardef\\twoheadleftarrow=\"3\\msa@11 \\global\\mathchardef\\leftleftarrows=\"3\\msa@12 \\global\\mathchardef\\rightrightarrows=\"3\\msa@13 \\global\\mathchardef\\upuparrows=\"3\\msa@14 \\global\\mathchardef\\downdownarrows=\"3\\msa@15 \\global\\mathchardef\\upharpoonright=\"3\\msa@16 \\let\\restriction=\\upharpoonright \\global\\mathchardef\\downharpoonright=\"3\\msa@17 \\global\\mathchardef\\upharpoonleft=\"3\\msa@18 \\global\\mathchardef\\downharpoonleft=\"3\\msa@19 \\global\\mathchardef\\rightarrowtail=\"3\\msa@1A \\global\\mathchardef\\leftarrowtail=\"3\\msa@1B \\global\\mathchardef\\leftrightarrows=\"3\\msa@1C \\global\\mathchardef\\rightleftarrows=\"3\\msa@1D \\global\\mathchardef\\Lsh=\"3\\msa@1E \\global\\mathchardef\\Rsh=\"3\\msa@1F \\global\\mathchardef\\rightsquigarrow=\"3\\msa@20 \\global\\mathchardef\\leftrightsquigarrow=\"3\\msa@21 \\global\\mathchardef\\looparrowleft=\"3\\msa@22 \\global\\mathchardef\\looparrowright=\"3\\msa@23 \\global\\mathchardef\\circeq=\"3\\msa@24 \\global\\mathchardef\\succsim=\"3\\msa@25 \\global\\mathchardef\\gtrsim=\"3\\msa@26 \\global\\mathchardef\\gtrapprox=\"3\\msa@27 \\global\\mathchardef\\multimap=\"3\\msa@28 \\global\\mathchardef\\therefore=\"3\\msa@29 \\global\\mathchardef\\because=\"3\\msa@2A \\global\\mathchardef\\doteqdot=\"3\\msa@2B \\let\\Doteq=\\doteqdot \\global\\mathchardef\\triangleq=\"3\\msa@2C \\global\\mathchardef\\precsim=\"3\\msa@2D \\global\\mathchardef\\lesssim=\"3\\msa@2E \\global\\mathchardef\\lessapprox=\"3\\msa@2F \\global\\mathchardef\\eqslantless=\"3\\msa@30 \\global\\mathchardef\\eqslantgtr=\"3\\msa@31 \\global\\mathchardef\\curlyeqprec=\"3\\msa@32 \\global\\mathchardef\\curlyeqsucc=\"3\\msa@33 \\global\\mathchardef\\preccurlyeq=\"3\\msa@34 \\global\\mathchardef\\leqq=\"3\\msa@35 \\global\\mathchardef\\leqslant=\"3\\msa@36 \\global\\mathchardef\\lessgtr=\"3\\msa@37 \\global\\mathchardef\\backprime=\"0\\msa@38 \\global\\mathchardef\\risingdotseq=\"3\\msa@3A \\global\\mathchardef\\fallingdotseq=\"3\\msa@3B \\global\\mathchardef\\succcurlyeq=\"3\\msa@3C \\global\\mathchardef\\geqq=\"3\\msa@3D \\global\\mathchardef\\geqslant=\"3\\msa@3E \\global\\mathchardef\\gtrless=\"3\\msa@3F \\global\\mathchardef\\sqsubset=\"3\\msa@40 \\global\\mathchardef\\sqsupset=\"3\\msa@41 \\global\\mathchardef\\trianglerighteq=\"3\\msa@44 \\global\\mathchardef\\trianglelefteq=\"3\\msa@45 \\global\\mathchardef\\bigstar=\"0\\msa@46 \\global\\mathchardef\\between=\"3\\msa@47 \\global\\mathchardef\\blacktriangledown=\"0\\msa@48 \\global\\mathchardef\\blacktriangleright=\"3\\msa@49 \\global\\mathchardef\\blacktriangleleft=\"3\\msa@4A \\global\\mathchardef\\blacktriangle=\"0\\msa@4E \\global\\mathchardef\\triangledown=\"0\\msa@4F \\global\\mathchardef\\eqcirc=\"3\\msa@50 \\global\\mathchardef\\lesseqgtr=\"3\\msa@51 \\global\\mathchardef\\gtreqless=\"3\\msa@52 \\global\\mathchardef\\lesseqqgtr=\"3\\msa@53 \\global\\mathchardef\\gtreqqless=\"3\\msa@54 \\global\\mathchardef\\Rrightarrow=\"3\\msa@56 \\global\\mathchardef\\Lleftarrow=\"3\\msa@57 \\global\\mathchardef\\veebar=\"2\\msa@59 \\global\\mathchardef\\barwedge=\"2\\msa@5A \\global\\mathchardef\\doublebarwedge=\"2\\msa@5B \\global\\mathchardef\\angle=\"0\\msa@5C \\global\\mathchardef\\measuredangle=\"0\\msa@5D \\global\\mathchardef\\sphericalangle=\"0\\msa@5E \\global\\mathchardef\\varpropto=\"3\\msa@5F \\global\\mathchardef\\smallsmile=\"3\\msa@60 \\global\\mathchardef\\smallfrown=\"3\\msa@61 \\global\\mathchardef\\Subset=\"3\\msa@62 \\global\\mathchardef\\Supset=\"3\\msa@63 \\global\\mathchardef\\Cup=\"2\\msa@64 \\let\\doublecup=\\Cup \\global\\mathchardef\\Cap=\"2\\msa@65 \\let\\doublecap=\\Cap \\global\\mathchardef\\curlywedge=\"2\\msa@66 \\global\\mathchardef\\curlyvee=\"2\\msa@67 \\global\\mathchardef\\leftthreetimes=\"2\\msa@68 \\global\\mathchardef\\rightthreetimes=\"2\\msa@69 \\global\\mathchardef\\subseteqq=\"3\\msa@6A 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\\def\\refFormat{\\catcode`\\^^M=10} \\def\\Refs#1#2{Refs.~\\use{Ref.#1},\\use{Ref.#2}} \\def\\Refsand#1#2{Refs.~\\use{Ref.#1} and~\\use{Ref.#2}} \\def\\Refsrange#1#2{Refs.~\\use{Ref.#1}--\\use{Ref.#2}} \\superrefsfalse % \\def\\Chapter#1{Chapter~\\use{Chap.#1}} \\def ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601500_arXiv.txt": {
        "abstract": "{ The dynamics of the explosive burning process is highly sensitive to the flame speed model in numerical simulations of type Ia supernovae.  Based upon the hypothesis that the effective flame speed is determined by the unresolved turbulent velocity fluctuations, we employ a new subgrid scale model which includes a localised treatment of the energy transfer through the turbulence cascade in combination with semi-statistical closures for the dissipation and non-local transport of turbulence energy. In addition, subgrid scale buoyancy effects are included. In the limit of negligible energy transfer and transport, the dynamical model reduces to the Sharp-Wheeler relation. According to our findings, the Sharp-Wheeler relation is insuffcient to account for the complicated turbulent dynamics of flames in thermonuclear supernovae. The application of a co-moving grid technique enables us to achieve very high spatial resolution in the burning region. Turbulence is produced mostly at the flame surface and in the interior ash regions. Consequently, there is a pronounced anisotropy in the vicinity of the flame fronts. The localised subgrid scale model predicts significantly enhanced energy generation and less unburnt carbon and oxygen at low velocities compared to earlier simulations. ",
        "introduction": "For supernovae of type Ia, \\citet{HoyFow60} proposed a thermonuclear runaway initiated in C+O white dwarfs close to the Chandrasekhar limit as the cause of the explosion. Since the original proposal, there has been vivid controversy of how such an explosion might come about and what the exact physical mechanism could be. Today the computational facilities to process three-dimensional large-eddy simulations (LES) of the explosion event are available. Remarkably, these powerful means have not aided in arriving at a consensus yet. The disagreement stems from some crucial questions. Firstly, what is the appropriate flame speed model?  Secondly, does the explosion completely proceed as a deflagration, or does a transition to a delayed detonation set in at some point?  The deflagration to detonation transition (DDT) proposed by \\citet{Khok91a} and \\citet{WoosWeav94} appears to resolve the drawbacks of the pure deflagration model.  In particular, the energy output obtained from simulations with artificial DDT is closer to the observed one, and less carbon and oxygen is left behind \\citep{GamKhok04,GamKhok05,GolNie05}.  For the theoretical understanding of thermonuclear supernovae, however, the lack of a convincing explanation for the initiation of the transition is unsatisfactory \\citep{KhokOr97,NieWoos97,Nie99,ZingWoos05}. In the aforementioned numerical models, a DDT is artifically triggered at more or less arbitrary instants of time.  As for the flame speed model, the controversy is whether subgrid scale (SGS) turbulence is mostly driven by Rayleigh-Taylor instabilities or dominated by the energy transfer through the turbulence cascade. The former point of view holds that the magnitude of SGS velocity fluctuations $v^{\\prime}$ is basically given by the Sharp-Wheeler relation \\citep{DavTay50,Sharp84} \\begin{equation} \\label{eq:vel_sharp} v_{\\mathrm{RT}}(l)=0.5\\sqrt{l g_{\\mathrm{eff}}} \\end{equation} where $v_{\\mathrm{RT}}(l)$ is the asymptotic rise velocity of a perturbation of size $l$ due to buoyancy. The effective gravity $g_{\\mathrm{eff}}$ is determined by the density contrast at the interface between burned and unburned material. Setting the turbulent flame speed equal to $v_{\\mathrm{RT}}(\\Delta)$, where $\\Delta$ is the resolution of the numerical grid, has been used in some simulations of type Ia supernovae \\citep{GamKhok03,CaldPle03,GamKhok04}. However, simple scaling arguments disfavour this proposition \\citep{NieHille95a,NieKer97}. Assuming that non-linear interactions between turbulent eddies of different size $l$ set up a Kolmogorov spectrum, the root mean square turbulent velocity fluctuations obey the scaling law $v^{\\prime}(l)\\propto l^{1/3}$. Since the Sharp-Wheeler relation implies $v_{\\mathrm{RT}}(l)\\propto l^{1/2}$, we have $v_{\\mathrm{RT}}(l)/v^{\\prime}(l)~\\propto l^{1/6}\\rightarrow 0$ towards decreasing length scales. Consequently, \\citet{NieHille95a} adopted a SGS model based on the dynamical equation for the turbulence energy $k_{\\mathrm{sgs}}$, i.e. the kinetic energy of unresolved velocity fluctuations \\citep{Schumann75}. The major weakness of their approach arises from the fairly tentative closures which were formulated \\emph{ad hoc} for LES of stellar convection \\citep{Clement93}.  Various refutations of the scaling argument have been put forward. To begin with, the spectrum of turbulence energy might be different from the Kolmogorov spectrum. However, recent direct numerical simulations support the hypothesis of a Kolmogorov spectrum in buoyancy-driven turbulent combustion \\citep{ZingWoos05}. A more serious concern is that there might be not enough time to reach the state of developed turbulence with a Kolmogorov spectrum in the transient scenario of a supernova explosion. This question is difficult to settle \\emph{a priori}.  For this reason, we took an unbiased point of view and accommodated buoyancy effects in the form of an Archimedian force term in the SGS turbulence energy model.  In contrast to the previously used SGS turbulence energy model, the new localised model which is thoroughly discussed in paper I neither presumes isotropy nor a certain turbulence energy spectrum function. This is possible by virtue of a dynamical procedure for the determination of the SGS eddy-viscosity $\\nu_{\\mathrm{sgs}}= C_{\\nu}\\Delta k_{\\mathrm{sgs}}^{1/2}$ which was adapted from \\cite{KimMen99}. Furthermore, we apply the co-expanding grid introduced by \\cite{Roep05}. The growth of the cutoff length $\\Delta$ due to the grid expansion poses a challenge for the SGS model because of the partitioning between resolved energy and SGS energy changes in time. We will show that this rescaling effect can be taken into account by utilising the dynamical procedure for the calculation of eddy-viscosity parameter $C_{\\nu}$. The rescaling algorithm as well as the computation of the Archimedian force is explained in Sect.~\\ref{sc:flame_speed_model}, followed by the discussion of results from three-dimension numerical simulations in Sect.~\\ref{sc:num_simul}. It is demonstrated that the newly proposed SGS model substantially alters the predictions of the deflagration model. In particular, we will analyse the significance of SGS buoyancy affects. ",
        "conclusions": "We applied the SGS turbulence energy model to the large eddy simulation of turbulent deflagration in thermonuclear supernova explosions. The novel features of this model are a localised closure for the rate of energy transfer, an additional Archimedian force term which accounts for buoyancy effects on unresolved scales and the rescaling of the SGS turbulence energy due to the shift of the cutoff length in simulations with a co-expanding grid. We found that the production of turbulence is largely confined to the regions near the flame fronts and in the interior ash regions. Consequently, there is pronounced anisotropy at the flame surface which can be tackled by the localised SGS model only.  The Archimedian force contributes noticeably to the turbulent flame speed, particularly once the flame surface has grown substantially. However, the dominating effect is the energy transfer through the turbulence cascade. In the late stage of the explosion, sustained turbulence energy comes from the rescaling, while the major dynamical contribution is SGS dissipation. Furthermore, it appears that numerical grids with more than $N=256^{3}$ cells in one octant are necessary in order to sufficiently resolve the turbulent dynamics in the burning regions and to obtain converged results.  An investigation of probability density functions over the flame fronts (see Fig.~\\ref{fg:pdf_speed}) reveals that the Rayleigh-Taylor velocity scale $v_{\\mathrm{RT}}(\\Delta_{\\mathrm{eff}})$ given by the Sharp-Wheeler relation~(\\ref{eq:vel_sharp}) is not negligible compared to the SGS turbulence velocity $q_{\\mathrm{sgs}}$, once the regime of fully turbulent burning has been entered.  This reflects the slow decrease of the ratio $v_{\\mathrm{RT}}(\\Delta_{\\mathrm{eff}})/q_{\\mathrm{sgs}}$ with the numerical cutoff scale according to the scaling argument discussed in the introduction. The underlying scaling relations do not necessarily apply to transient and inhomogenous flows as in the supernova explosion scenario.  The PDF for $q_{\\mathrm{sgs}}$ shows a considerably wider spread than the sharply peaked PDF for $v_{\\mathrm{RT}}(\\Delta_{\\mathrm{eff}})$. We interpret this observation as a consequence of the additional physics in the localised SGS model, which also encompasses turbulent energy transfer (i.e. interaction between resolved and subgrid scales) and turbulent transport (i.e. non-local interactions among subgrid scales). The relation between the Sharp-Wheeler and the turbulence energy models may be analogous to the relation between the mixing length and Reynolds stress models of convection.  The final kinetic energy in the simulation with the highest resolution is about $6\\cdot 10^{50}\\,\\mathrm{erg}$. The produced mass of iron group elements, $0.58M_{\\sun}$, falls within the range deduced from observations of type Ia supernovae \\citep{Leib00}. However, some observed events are substantially more energetic. Regarding the numerical simulations, the explosion energy is very sensitive to the initial conditions and the localised SGS model appears to increase the sensitivity even further. Using initial conditions that are different to the highly artificial centrally ignited flame is in progress.  A persistent difficulty is the large amount of left over carbon and oxygen (see Fig.~\\ref{fg:densty_150} and~\\ref{fg:mass_distrb_500}). It is not clear yet to what extent this problem can be solved with the aid of the localised SGS model in simulations with more realistic ignition scenarios. It appears more likely that a different mode of burning is required in the late explosion phase. The currently implemented numerical burning extinction at a density threshold of $10^{7}\\,\\mathrm{g\\,cm^{-3}}$ is mostly arbitrary and should be replaced by a physical criterion motivated by the properties of distributed burning at low densities. This might turn out to be an alternative to the DDT scenario."
    },
    "0601/physics0601198_arXiv.txt": {
        "abstract": "This paper provides a detailed study of turbulent statistics and scale-by-scale budgets in turbulent Rayleigh-B\\'enard convection. It aims at testing the applicability of \\cite{k41} and \\cite{bolgiano59} theories in the case of turbulent convection and at improving the understanding of the underlying inertial range scalings, for which a general agreement is still lacking. Particular emphasis is laid on anisotropic and inhomogeneous effects, which are often observed in turbulent convection between two differentially heated plates. For this purpose, the SO(3) decomposition of structure functions \\citep*{arad99a} and a method of description of inhomogeneities  proposed by \\cite{danaila01} are used to derive inhomogeneous and anisotropic generalizations of Kolmogorov and Yaglom equations applying to Rayleigh-B\\'enard convection, which can be extended easily to other types of anisotropic and/or inhomogeneous flows. The various contributions to these equations are computed in and off the central plane of a convection cell using data produced by a Direct Numerical Simulation of turbulent Boussinesq convection at $\\ray=10^6$ and $\\pr=1$ with aspect ratio $A=5$. The analysis of the isotropic part of the Kolmogorov equation demonstrates that the shape of the third-order velocity structure function is significantly influenced by buoyancy forcing and large-scale inhomogeneities, while the isotropic part of the mixed third-order structure function $\\moy{(\\Delta\\theta)^2\\Delta\\vec{u}}$ appearing in Yaglom equation exhibits a clear scaling exponent 1 in a small range of scales. The magnitudes of the various low $\\ell$ degree anisotropic components of the equations are also estimated and  are  shown to be comparable to their isotropic counterparts at moderate to large scales. The analysis of anisotropies notably reveals that computing reduced structure functions (structure functions computed at fixed depth for correlation vectors $\\vec{r}$ lying in specific planes only) in order to reveal scaling exponents predicted by isotropic theories is misleading in the case of fully three-dimensional turbulence in the bulk of a convection cell, since such quantities involve linear combinations of different $\\ell$ components which are not negligible in the flow. This observation also indicates that using single points measurements together with the Taylor hypothesis in the particular direction of a mean flow to test the predictions of asymptotic dimensional isotropic theories of turbulence or to calculate intermittency corrections to these theories may lead to significant biases for mildly anisotropic three-dimensional flows. A qualitative analysis is finally used to show that the influence of buoyancy forcing at scales smaller than the Bolgiano scale is likely to remain important up to $\\ray=10^9$, thus preventing Kolmogorov scalings from showing up in convective flows at lower Rayleigh numbers. ",
        "introduction": "The quest for inertial range scaling laws in turbulent convection has been very active in recent years (\\textit{e.~g.} \\cite{chilla93,benzi94,calzavarini02,verzicco03,ching04}). Their determination is expected to give some important insight into the thermal and mechanical processes at work in the flow. To this end, turbulent thermal convection is investigated in convection cells heated from below using laboratory and numerical experiments, within the framework of the Boussinesq approximation. Such a flow exhibits two essential properties: it is both strongly anisotropic and inhomogeneous. Anisotropy comes from gravity, while inhomogeneity results from the presence of top and bottom horizontal boundaries in convection cells. As a consequence, inertial range scalings of convective turbulence, when they can ever be observed, depend  strongly on the vertical coordinate.  Accordingly, asymptotic theories of turbulence constructed under the assumptions of homogeneity and isotropy may be partially irrelevant to understand the observed properties of turbulent convection. Neglecting intermittency effects, the classical picture regarding scaling laws in this flow is that Bolgiano-Obukhov turbulence (\\cite{bolgiano59,obukhov59}, hereafter BO59 theory) should be present on correlation lengths  $r$ larger than the so-called Bolgiano length. A dimensional estimate of this length \\citep{chilla93} is given by \\begin{equation} \\label{bolglength} L_B=\\frac{\\nuss^{1/2}d}{(\\ray\\pr)^{1/4}}~, \\end{equation} where $d$ is the depth of the convective layer, $\\nuss$ is the Nusselt number, $\\ray$ is the Rayleigh number and $\\pr$ is the Prandtl number. Instead, homogeneous and isotropic Kolmogorov turbulence (\\cite{k41}, hereafter K41 theory) should be observed for $r<L_B$. In the Bolgiano-Obukhov subrange, a dominant balance between buoyancy forcing (for an unstably stratified layer) and the third-order structure function occurs, which modifies the turbulent energy cascade substantially in comparison to K41: while longitudinal velocity increment scalings $\\moy{\\Delta u_r(r)}\\sim r^{1/3}$ are predicted asymptotically in the inertial range for K41, they should follow $\\moy{\\Delta u_r(r)}\\sim r^{3/5}$ in the BO59 regime. As far as temperature increments are concerned, $\\moy{\\Delta \\theta(r)}\\sim r^{1/3}$ results from K41 (passive scalar scalings), while $\\moy{\\Delta \\theta(r)}\\sim r^{1/5}$ is expected from dimensional arguments in BO59 theory. These different regimes should in principle be detected through the measurements of structure functions or spectra when turbulence is sufficiently developed. However, clear evidences of K41 or BO59 scalings are still lacking in both numerical and experimental convection, even at very high Rayleigh numbers. Mixed scalings, which are not compatible with dimensional analysis, have sometimes been detected (\\textit{e. g.} K41 scalings for the velocity and BO59 scalings for temperature in the same wave number range, see for instance \\cite{verzicco03}). BO59 scalings have been reported in the bulk of experimental convection cells \\citep{benzi94,ching04}, and in a numerical experiment \\citep*{calzavarini02}, close to the horizontal walls of the convection cell. Direct scaling laws could not be observed for reduced structure functions (structure functions computed at fixed depth for correlation vectors $\\vec{r}$ lying in horizontal planes only) in the latter study, so that the derivation of scaling exponents by \\cite{calzavarini02} relies on Extended Self Similarity (ESS, see \\cite{benzi93}).  Results obtained for various anisotropic flows, such as the channel flow \\citep{arad99b}, demonstrate that it is sometimes possible to observe scaling laws by plotting structure functions directly, even at modest Reynolds numbers, provided that the complete computation and SO(3) decomposition of structure functions are performed: in the work of \\cite{arad99b}, the difference between reduced structure functions and the isotropic (angular mean) component of the complete structure functions is striking.  This spherical harmonics decomposition of structure functions has also been computed by \\cite{biferale03} for homogeneous Rayleigh-B\\'enard convection (HRB) in a regime where $L_B$ is comparable to the box size, thus preventing any possibility of BO59 scalings. The anisotropic scaling exponents reported for HRB are anomalous, which means that they do not match dimensional analysis predictions either. The understanding of inertial range scalings in convective turbulence and their very existence even at high Rayleigh numbers therefore remains a puzzling problem that still deserves important efforts.  This paper aims at presenting a complete description of turbulent statistics in Boussinesq convection  at $\\ray=10^6$ and $\\pr=1$ in order to test explicitly the assumptions of both K41 and BO59 theories in such a flow and to try to  clarify some of the previously mentioned  problems regarding the occurrence of inertial range scaling laws in convective turbulence. The study shows that inhomogeneity, anisotropy and the presence of buoyancy forcing at all scales all contribute to the absence of scaling laws, and points out that very high Rayleigh numbers should be reached in order to be able to observe the definite signature of spectral scalings in turbulent convection. The presentation of the results is as follows. The SO(3) decomposition of statistical averages is first used in \\S\\ref{theorie} to derive generalized Kolmogorov and Yaglom equations including inhomogeneous and anisotropic terms, which describe respectively the production, transport and dissipation of velocity and temperature fluctuations. A detailed numerical study of these equations  is then proposed using Direct Numerical Simulation (DNS) data. The DNS and data processing algorithms are first described in \\S\\ref{numdetails}. The various contributions to the equations are then presented in \\S\\ref{results} at the center of the convection cell and off the central plane, with particular emphasis laid on the anisotropic and inhomogeneous effects inferred from the analysis. The consequences and perspectives offered by these results are finally discussed in \\S\\ref{discuss}. ",
        "conclusions": "Summary and discussion} New results on inertial range scaling laws, scale-by-scale budgets and correlation functions in turbulent Rayleigh-B\\'enard convection at moderate Rayleigh and Reynolds numbers have been reported in this paper. A derivation of  generalized Kolmogorov and Yaglom equations~(\\ref{eqVfinal})-(\\ref{eqTfinal}), which notably makes use of the SO(3) decomposition of statistical averages, has first been presented. The formalism provides a convenient way of disentangling inhomogeneous and anisotropic effects in convective turbulence, but it can also be used to study other anisotropic and/or inhomogeneous numerically simulated or experimental flows. For this purpose, the buoyancy forcing term in the velocity equation should be replaced by any appropriate anisotropic forcing mechanism.  The analysis of scale-by-scale budgets in a  convection DNS at $\\ray=10^6$ has led to the following conclusions. First of all, a compensation effect between the classical $-4/3\\moy{\\varepsilon}r$ term of the generalized Kolmogorov equation on one hand, inhomogeneous terms, viscous terms and buoyancy forcing on the other hand, has been demonstrated in the isotropic sector of the velocity equation. Among these three terms, buoyancy is the most important one. This compensation effect prevents any clear scaling law from showing up at $\\rey_\\lambda=30$ in the case of the third-order velocity structure function. Also, equation~({\\ref{bolglength}) has been shown to underestimate significantly the effective value of the Bolgiano length, at least for convective flows in moderate aspect ratio containers. The analysis of the mixed temperature structure function $\\moy{\\Theta_R}^0_0$ has revealed a scaling behaviour compatible with both K41 and BO59 theories in the central plane, in spite of the presence of inhomogeneous effects on large scales. Meanwhile, neglecting the term related to advection of the mean temperature profile, which is required in the BO59 theory, has been shown to be justified. Scale-by-scale budgets computed off the central plane have led to similar conclusions. An interesting inhomogeneous effect at $z=1/4$ is the significant contribution of pressure diffusion normal to the bottom wall.  A spherical harmonics spectrum of structure functions has finally revealed that the low $\\ell$ degree contributions to the third-order structure functions are not small in comparison to their isotropic counterparts at moderate to large scales. It has particularly been shown that using reduced structure functions at fixed polar angle in the bulk of a convection cell in order to reveal scaling behaviour predicted by isotropic theories is misleading, since these structure functions involve linear combinations of various $\\ell$ components which scale differently with respect to $r$. This argument also indicates that using single points measurements  in the bulk of three-dimensional flows  together with the Taylor hypothesis in the particular direction of a mean flow to test the predictions of asymptotic dimensional isotropic theories of turbulence or to calculate intermittency corrections to  these theories may lead to significant biases. It has however been outlined that using reduced structure functions may still be appropriate when the direction along which these functions are computed has a special importance, as is the case for instance in layered flows or boundary layers. Also, it has been pointed out that spurious anisotropies related to the use of cartesian domains in simulations of mildly turbulent flows, which are difficult to quantify, certainly interfere with the genuine anisotropic effects due to gravity. Thus, the respective amplitudes of the various $\\ell$ components of structure functions at  moderate Reynolds number has to be considered with some caution. Independently of the problem of the origin of anisotropies, these results confirm the analysis of \\cite{arad99b} for the channel flow: disentangling anisotropic effects can not be avoided if one wants to study the scaling behaviour of anisotropic systems correctly. Besides, even though the study of the $r$-dependence of different $\\ell$ components in equations~(\\ref{eqVfinal})-(\\ref{eqTfinal}) seems to indicate qualitatively small-scale return to isotropy in the case of Rayleigh-B\\'enard convection, it can not be ruled out that anisotropic effects  persist at small scales, even at very high Reynolds numbers. Small-scale anisotropies (called ``ramp-cliff structures'') have for instance been reported in studies of turbulent mixing of passive scalars (\\cite{mestayer76,antonia78}, see also \\cite{warhaft2000}).  An important issue is to which extent the present results can be used to infer scaling behaviour at very high Rayleigh and Reynolds numbers. The still important contribution  of buoyancy forcing to the scale-by-scale budget at scales smaller than $L_B$ naturally raises the question of how many scale decades are required in a turbulent convective flow in order to be able to identify definite scalings. In the soft turbulence regime, where $\\nuss\\sim \\ray^{\\gamma}$ with $1/4<\\gamma<1/3$ \\citep*{grossmann00}, the Bolgiano length should not depend significantly on $\\ray$. For instance, following equation~(\\ref{bolglength}), $L_B\\sim\\ray^{-3/28}$ is obtained for $\\gamma=2/7$ \\citep{procaccia89}. In the hard regime where it has recently been shown that $\\nuss\\sim (\\ray\\pr)^{1/2}$ \\citep{calzavarini05}, the Bolgiano length should remain constant (with respect to both $\\ray$ and $\\pr$ in that case, while the $\\pr$ dependence in the hard turbulence regime depends on whether $\\pr>1$ or $\\pr<1$), thus leading to the same result at all Rayleigh numbers. For the present aspect ratio and Prandtl number, $L_B$ should therefore remain close to 1 at all Rayleigh numbers. On the contrary, the ratio between the dissipation scale and the effective Bolgiano scale is expected to increase with increasing Rayleigh number, which should in principle help to observe K41 scalings if $\\eta\\ll r\\ll L_B$ can be achieved \\citep*{grossmannlvov93}. One should however derive more precise conditions of applicability of K41 in that case. Assuming K41 to be valid for $r\\leq L_B$, the buoyancy term, which equals the $-4/3\\moy{\\varepsilon} r$ term at $L_B$, would scale like $r^{5/3}$ in that range. This approximation looks rather crude for $r$ close to $L_B$, but as the buoyancy term should scale like $r^{9/5}$  above $L_B$ according to BO59, the actual local scaling exponent close to $L_B$ may not be very different from this 5/3 value. The important point here is that this exponent be larger than 1 (which is the exponent of the $-4/3\\moy{\\varepsilon} r$ term). According to these scalings, non-dissipative scales as small as $3\\times10^{-2} L_B$ should be available in order for the buoyancy term to become ten times smaller than the $-4/3\\moy{\\varepsilon} r$ term, which roughly corresponds to the conditions of the present simulation, in which not definite K41 plateau can be observed (figure~\\ref{figS3l=0}). In order to find a range of scales in which the buoyancy term would become at least one hundred times smaller than the $-4/3\\moy{\\varepsilon}r$ term, one should then look for a regime where $\\eta<10^{-3} L_B$. The required Rayleigh number can be evaluated by using $\\eta\\sim \\ray^{-9/28}$, proposed by \\cite{grossmannlohse93}. This relation can be calibrated using the results of the present simulation at $\\ray=10^6$, for which $\\eta=0.016$. One finally finds that $\\ray$ has to exceed several times $10^9$ to obtain a flow with $\\eta< 10^{-3} L_B$. Even in this situation, definite scalings are expected only in the subrange $10\\eta<r<0.1 L_B$, that is over one decade. The main problem is that two scale separations  $r\\ll L_B$ and $r\\gg\\eta$ must simultaneously be satisfied, which can only occur at very high Rayleigh number. This is an important argument to understand why inertial range scalings have been difficult to identify in experimental or numerical turbulent convection studies.  For the preceding reasons, it is probable that various results on scaling exponents obtained in the range $10^6<\\ray<10^9$ (\\textit{e.~g.} \\cite{calzavarini02}, who have computed local scaling exponents from reduced structure functions) are biased simultaneously by anisotropic effects and significant buoyancy forcing at $r<L_B$. The lack of complete K41 or BO59 scalings in the simulations at very high Rayleigh number (up to $\\ray=2\\times 10^{11}$) of \\cite{verzicco03} may also be a hint that inertial range scalings in anisotropic flows such as Rayleigh-B\\'enard convection can not be understood easily using isotropic theories. If inhomogeneity and isotropy were to play an important role in the scaling behaviour of very high Rayleigh number convective turbulence, the present results could serve as  guidelines to study them in detail. More generally, the SO(3) decomposition of structure functions appears to be a very powerful (and necessary) tool to study various anisotropic flows. It should be particularly helpful to use it when possible in  research domains such as astrophysics or atmospheric sciences, where  anisotropic turbulence is ubiquitous. Such an investigation for strongly stratified non-Boussinesq turbulent convection, which is relevant to the Sun and other stars, is currently underway.  \\smallskip  The author acknowledges several fruitful discussions with A.~A.~Schekochihin and thanks F. Anselmet, F. Ligni\\`eres and M. Rieutord for many insightful comments and criticisms on early versions of the manuscript. Numerical simulations have been performed on the IBM SP4 supercomputer of  Institut du D\\'eveloppement et des Ressources en Informatique Scientifique (IDRIS, Orsay, France), which is gratefully acknowledged.  \\medskip  \\medskip  \\medskip"
    },
    "0601/astro-ph0601446_arXiv.txt": {
        "abstract": "Period-colour (PC) and amplitude-colour (AC) relations are studied for the Large Magellanic Cloud (LMC) Cepheids under the theoretical framework of the hydrogen ionization front (HIF) - photosphere interaction. LMC models are constructed with pulsation codes that include turbulent convection, and the properties of these models are studied at maximum, mean and minimum light. As with Galactic models, at maximum light the photosphere is located next to the HIF for the LMC models. However very different behavior is found at minimum light. The long period ($P>10$days) LMC models imply that the photosphere is disengaged from the HIF at minimum light, similar to the Galactic models, but there are some indications that the photosphere is located near the HIF for the short period ($P<10$ days) LMC models. We also use the updated LMC data to derive empirical PC and AC relations at these phases. Our numerical models are broadly consistent with our theory and the observed data, though we discuss some caveats in the paper. We apply the idea of the HIF-photosphere interaction to explain recent suggestions that the LMC period-luminosity (PL) and PC relations are non-linear with a break at a period close to 10 days. Our empirical LMC PC and PL relations are also found to be non-linear with the $F$-test. Our explanation relies on the properties of the Saha ionization equation, the HIF-photosphere interaction and the way this interaction changes with the phase of pulsation and metallicity to produce the observed changes in the LMC PC and PL relations. ",
        "introduction": "\\citet{cod47} found that the Galactic Cepheids follow a spectral type that is independent of their pulsational periods at maximum light and gets later as the periods increase at minimum light. \\citet[][hereafter SKM]{sim93} used radiative hydrodynamical models to explain these observational phenomena as being due to the location of the hydrogen ionization front (HIF) relative to the photosphere. Their results agreed very well with Code's observation. SKM further used the Stefan-Boltzmann law applied at the maximum and minimum light, together with the fact that radial variation is small in the optical \\citep{cox80}, to derive:  \\begin{eqnarray} \\log T_{max} - \\log T_{min} = {1\\over{10}}(V_{min} - V_{max}), \\end{eqnarray}  \\ni where $T_{max/min}$ are the effective temperature at the maximum/minimum light, respectively. If $T_{max}$ is independent of the pulsation period $P$ (in days), then equation (1) predicts there is a relation between the $V$-band amplitude and the temperature (or the colour) at minimum light, and vice versa. In other words, if the period-colour (PC) relation at maximum (or minimum) light is flat, then there is an amplitude-colour (AC) relation at minimum (or maximum) light. Equation (1) has shown to be valid theoretically and observationally for the classical Cepheids and RR Lyrae variables \\citep{kan04,kan05}.  For the RR Lyrae variables, \\citet{kan95} and \\citet{kan96} used linear and non-linear hydrodynamic models of RRab stars in the Galaxy to explain why RRab stars follow a flat PC relation at {\\it minimum} light. Later, \\citet{kan05} used MACHO RRab stars in the LMC to prove that LMC RRab stars follow a relation such that higher amplitude stars are driven to cooler temperatures at maximum light. Similar studies were also carried out for Cepheid variables, as in SKM, \\citet{kan96}, \\citet[][hereafter Paper I]{kan04} and \\citet[][hereafter Paper II]{kan04a}. In contrast to the RR Lyrae variables, Cepheids show a flat PC relation at the {\\it maximum} light, and there is a AC relation at the minimum light. Therefore, the PC relation and the AC relation are intimately connected. All these studies are in accord with the predictions of equation (1).  In Paper I, the Galactic, Large Magellanic Cloud (LMC) and Small Magellanic Cloud (SMC) Cepheids were analyzed in terms of the PC and AC relations at the phase of maximum, mean and minimum light. One of the motivations for this paper originates from recent studies on the non-linear LMC PC relation \\citep[as well as the period-luminosity, PL, relation. See Paper I;][]{tam02a,san04,nge05}: the optical data are more consistent with two lines of differing slopes which are continuous or almost continuous at a period close to 10 days. Paper I also applied the the $F$-test \\citep{wei80} to the PC and AC relations at maximum, mean and minimum $V$-band light for the Galactic, LMC and SMC Cepheids. The $F$-test results implied that the LMC PC relations are broken or non-linear, in the sense described above, across a period of 10 days, at mean and minimum light, but only marginally so at maximum light. The results for the Galactic and SMC Cepheids are similar, in a sense that at mean and minimum light the PC relations do not show any non-linearity and the PC(max) relation exhibited marginal evidence of non-linearity. For the AC relation, Cepheids in all three galaxies supported  the existence of two AC relations at maximum, mean and minimum light. In addition, the Cepheids in these three galaxies also exhibited evidence of the PC-AC connection, as implied by equation (1), which give further evidence of the HIF-photosphere interactions as outlined in SKM.  To further investigate the connection between equation (1) and the HIF-photosphere interaction, and also to explain Code's observations with modern stellar pulsation codes, Galactic Cepheid models were constructed in Paper II. In contrast to SKM's purely radiative models, the stellar pulsation codes used in Paper II included the treatment of turbulent convection as outlined in \\citet{yec98}. One of the results from Paper II was that the general forms of the theoretical PC and AC relation matched the observed relations well. The properties of the PC and AC relations for the Galactic Cepheids with $\\log(P)>0.8$ can be explained with the HIF-photosphere interaction. This interaction, to a large extent, is independent of the pulsation codes used, the adopted ML relations, and the detailed input physics.  The aim of this paper is to extend the investigation of the connections between PC-AC relations and the HIF-photosphere interactions in theoretical pulsation models of LMC Cepheids, in addition to the Galactic models presented in Paper II. In Section 2, we describe the basic physics of the HIF-photosphere interaction. The updated observational data, after applying various selection criteria, that used in this paper are described in Section 3. In Section 4, the new empirical PC and AC relations based on the data used are presented. In Section 5, we outline our methods and model calculations, and the results are presented in Section 6. Examples of the HIF-photosphere interaction in astrophysical applications are given in Section 7. Our conclusions \\&  discussion are presented in Section 8. Throughout the paper, short and long period Cepheid are referred to Cepheids with period less and greater than 10 days, respectively. ",
        "conclusions": "In this paper, we have confronted updated PC and AC relations at maximum, mean and minimum light for LMC Cepheids observed by the OGLE team, and additional Cepheids from the literature, with theoretical, full amplitude pulsation models of LMC Cepheids. The observed PC and AC relations provide compelling evidence of a non-linearity or break at a period of 10 days. We also constructed theoretical Cepheid pulsation models appropriate for the LMC using the Florida pulsation codes \\citep{yec98} to study the HIF-photosphere interaction.  The empirical results presented in this paper, as well as in other papers such as \\citet{san04} and \\citet{nge05}, provide strong empirical evidence that the PC and PL relations for the LMC Cepheids are non-linear, in the sense described in previous sections. Issues such as extinction and a lack of long period Cepheids that may cause the non-linear LMC PL and PC relations have been addressed and argued against in Paper I, \\citet{san04} and \\citet{nge05}, and will not be repeated here. Other arguments against the non-linear LMC PL relation include the results presented in \\citet{per04}, as the authors found no evidence for a non-linear PL relation in the LMC at $JHK$-bands. However, \\citet{nge05} treated the data of \\citet{per04} extensively and found, in a statistically rigorous way, that the reason why \\citet{per04} found linear $JHK$ PL relations, is due to the small number of short period Cepheids ($\\sim18$) in their sample. \\citet{nge05} also reduce the number of OGLE/MACHO LMC Cepheids and show how the $F$-test can produce a non-significant result when the number of short/long period Cepheids become small. Instead, using the 2MASS data that are cross-correlated with MACHO Cepheids, \\citet{nge05} have found that the LMC $JH$-band PL relations are non-linear\\footnote{Non-linear PL relations are relatively easier to see in $V$-band than in $J$-band, as shown in \\citet{nge05}.} and the $K$-band PL relation starts to become linear. \\citet{nge05} also discussed why this is the case. Another argument against the non-linear PL relation is that the PL relation should be universal, as found in \\citet{gie05}. We argue that their results are based on a handful of Cepheids ($\\sim15$) and on short periods Cepheids in a cluster whose membership to the LMC is in question. Their shallower Galactic PL relation based on the revised infra-red surface brightness method also contradicts the steeper Galactic PL relation based on independent methods from open cluster main-sequence fitting \\citep{tam03,nge04,san04}.  It is worthwhile to point out that our sample selection does not affect the detection on non-linear LMC PL relation at mean light. Since the mean magnitudes of a Cepheid light curve is less affected by our constrains on selecting the Cepheids with good light curves, we can use the published (reddening corrected) mean magnitudes to test the non-linear LMC PL relation. The anonymous referee kindly provided a large sample of LMC Cepheids that combined the published mean $V$-band magnitudes from the OGLE, \\citet{seb02} and \\citet{cal91} datasets. There are a total of 115 long period Cepheids in this sample and the $F$-test still return a significant detection of the non-linear LMC PL relation. The OGLE+\\citet{seb02} combined data also give very similar results. Similar tests have also been done in \\citet{nge05} by using the MACHO data alone and the MACHO+\\citet{seb02} combined data. The non-linear LMC PL relation is still present from the $F$-test results on these two datasets. Therefore we believe our sample selection does not affect the detection of the non-linear LMC PL relation. The detection of non-linear LMC PL relation from totally independent OGLE and MACHO data, using totally independent reddening estimates, suggested that this non-linearity is real and our paper is the first attempt to theoretically explain this non-linearity in terms of the HIF-photosphere interaction.  Due to small number of LMC models, it is impossible to derive the theoretical PC and AC relations with a small error on the slope and compare directly to the empirical relations. However, these LMC models can be qualitatively compared to the observations by converting some physical quantities to the observable quantities and vice versa, such as the temperature-colour conversion. Hence we compared our model light curves to the observations in terms of theoretical PC and AC relations at the phases of maximum, mean and minimum light and also in terms of the Fourier parameters from theoretical light curves with observations. The theoretical quantities from the models generally agree with the observations, but it was found out that these models tend to have smaller amplitudes and (hence) the temperature is cooler at maximum light than the real Cepheids. Though our models have some drawbacks in this comparison, our main interest is in comparing the interaction of the photosphere and HIF as a function of phase with similar results presented in Paper II for Galactic Cepheid pulsation models. The aim is {\\it not} to compare our models rigorously with observations but rather to study models which match observations reasonably well in the context of the theoretical framework described in previous sections and in Paper I \\& II. Nevertheless we argued that the qualitative nature of the photosphere-HIF interaction is not seriously affected by these problems.  Our postulate is that at certain phases, this interaction can affect the PC relation due to the properties of the Saha ionization equation: specifically for reasonably low densities in Cepheid envelopes, hydrogen ionizes at a temperature that is almost independent of period. Consequently, when the photosphere is located at the base of the HIF, the photospheric temperature and hence the colour is almost independent of period. However, when this engagement occurs, but the density is greater, then the temperature at which hydrogen ionizes again becomes sensitive to global surroundings and hence on period. When the photosphere is not engaged with the HIF in this way, its temperature is again dependent on period and global stellar parameters.  For Galactic Cepheids, this HIF-photosphere interaction occurs mainly at maximum light for Cepheids with $\\log P > 0.8$ (Paper II). At minimum light, there is a strong correlation between the HIF-photosphere distance and period leading to a definite AC relation at minimum light for Galactic Cepheids (SKM, Paper I \\& II). In this paper, we have found tentative evidence that, for short period LMC models which match observations in the period-color plane, the HIF-photosphere interaction occurs at most phases but at densities which are too high to produce a flat PC relation. Why would these short period LMC Cepheids be different in this regard to short period Galactic Cepheids? One possibility could be that this is partly because these LMC Cepheids are hotter than their Galactic counterparts \\citep{kan04,san04}. The HIF-photosphere are disengaged for most of the pulsation cycle for long period LMC Cepheids. This happens because as the period increases, so does the $L/M$ ratio which pushes the HIF further inside the mass distribution. When the HIF-photosphere are disengaged in this way, the photospheric temperature is more dependent on density and hence on period. The change is sudden because the HIF-photosphere are either engaged or they are not. This can lead to a sudden change in the PC relation at 10 days as shown by the observations \\citep{tam02a,kan04,san04,nge05}. However, at maximum light the HIF-photosphere are engaged at low densities for long period LMC Cepheids leading to the observed flat PC relation for these stars. Taken together with equation (1), this theoretical scenario is consistent with the observed PC-AC behavior described in Paper I and in this study. The anonymous referee has noted that these suggestions about photospheric density can be tested by spectroscopic means.  We now enumerate some caveats to our argument that could be addressed in future papers.  \\begin{enumerate} \\item Since the SMC PC relation at mean light is linear (e.g., Paper I), how do SMC (i.e., metal-poor) models fit into the theoretical scenario outlined in this paper and Paper II, if at all? This is a difficult question and its full answer is beyond the scope of this paper. However, as the metallicity decreases, we do note that the SMC has a different ML relation to the LMC and Galaxy and so does the temperatures associated with the instability strip. These will change the relative location of the HIF and photosphere \\citep{kan95,kan96} and possibly alter the phase at which they interact. Further the amplitudes for SMC Cepheids are smaller due to the lower metallicity \\citep{pac00}. This will also affect the HIF-photosphere interaction. One difference which can be consistent with this is the fact that the PC relation at maximum light in the SMC is not flat (see Paper I) but it is the case for the Galaxy and LMC PC relations. This indicates that at maximum light, there is less interaction between the HIF and photosphere at low densities. This leads to an observed linear PC relation at mean light for the SMC Cepheids. These will be investigated further in a future paper in this series.  \\item Could the well-known Hertzsprung progression play any part in causing the observed changes in the Galactic and LMC PC relations?  \\item It may also be that higher order overtones becoming unstable or stable, though with the fundamental mode still being dominant, may also  have an impact on the PC relation in some as yet unknown way (Paper II).  \\item The behavior of short period LMC Cepheids still needs to be understood, for example, what causes the difference between the bottom left panels of Figures \\ref{c9deltalmc} and \\ref{c9delta}? That is, why is it that for short period Galactic/LMC Cepheids, the HIF-photosphere are disengaged/engaged? Our experience suggests that constructing short period full amplitude fundamental mode Cepheids requires more care than the long period case because the first overtone has a non-negligible growth rate. Because of this we feel a thorough study of these short period Cepheids merits a separate paper.  \\item Would more advanced pulsation codes which, for example, can match the observed amplitudes and which contain a more accurate model of time dependent turbulent convection, yield similar results, especially for Figure \\ref{c9delta}? Could such codes fare better in modeling short period LMC Cepheids?  \\end{enumerate}"
    },
    "0601/astro-ph0601393_arXiv.txt": {
        "abstract": "We present $K_{\\rm s}$-band surface photometry of NGC 7690 (Hubble type Sab) and NGC 4593 (SBb).  We find that, in both galaxies, a major part of the ``bulge'' is as flat as the disk and has approximately the same color as the inner disk.  In other words, the ``bulges'' of these galaxies have disk-like properties. We conclude that these are examples of ``pseudobulges'' -- that is, products of secular dynamical evolution.  Nonaxisymmetries such as bars and oval disks transport disk gas toward the center.  There, star formation builds dense stellar components that look like -- and often are mistaken for -- merger-built bulges but that were constructed slowly out of disk material. These pseudobulges can most easily be recognized when, as in the present galaxies, they retain disk-like properties.  NGC 7690 and NGC 4593 therefore contribute to the growing evidence that secular processes help to shape galaxies.  NGC 4593 contains a nuclear ring of dust that is morphologically similar to nuclear rings of star formation that are seen in many barred and oval galaxies.  The nuclear dust ring is connected to nearly radial dust lanes in the galaxy's bar.  Such dust lanes are a signature of gas inflow. We suggest that gas is currently accumulating in the dust ring and hypothesize that the gas ring will starburst in the future. The observations of NGC 4593 therefore suggest that major starburst events that contribute to pseudobulge growth can be episodic. ",
        "introduction": "\\pretolerance=15000 \\tolerance=15000  Internal secular evolution of galaxies is the dynamical redistribution of energy and angular momentum that causes galaxies to evolve slowly between rapid (collapse-timescale) transformation events that are caused by galaxy mergers.  Driving agents include nonaxisymmetries in the gravitational potential such as bars, oval disks, and global spiral structure.  Kormendy (1993) and Kormendy \\& Kennicutt (2004, hereafter KK) review the growing evidence that secular processes have shaped the structure of many galaxies.  The fundamental way that self-gravitating disks evolve~-- provided that there is an efficient driving agent -- is by spreading (Lynden-Bell \\& Kalnajs 1972; Lynden-Bell \\& Pringle 1974; Tremaine 1989; see Kormendy \\& Fisher 2005 for a review in the present context).  In general, it is energetically favorable to shrink the inner parts by expanding the outer parts.  In barred galaxies, one well known consequence is the production of ``inner rings'' around the end of the bar and ``outer rings'' at about 2.2 bar radii.  The most general consequence of secular evolution, and the one that is of interest in this paper, is that some disk gas is driven to small radii where it reaches high densities, feeds starbursts, and builds ``pseudobulges''.  Because of their high stellar densities and steep density gradients, pseudobulges superficially resemble -- and often are mistaken for -- bulges.  Following Sandage (1961) and Sandage \\& Bedke (1994), Renzini (1999) adopts this definition of a bulge: ``It appears legitimate to look at bulges as ellipticals that happen to have a prominent disk around them [and] ellipticals as bulges that for some reason have missed the opportunity to acquire or maintain a prominent disk.''  We adopt the same definition.  Our paradigm of galaxy formation then is that bulges and ellipticals both formed via galaxy mergers (e.{\\ts}g., Toomre 1977; Steinmetz \\& Navarro 2002, 2003), a conclusion that is well supported by observations (see Schweizer 1990 for a review).  Pseudobulges are therefore fundamentally different from bulges -- they were built slowly out of the disk. Two well developed examples, one in the unbarred galaxy NGC 7690 and one in the barred galaxy NGC 4593, are the subjects of this paper.  Hierarchical clustering and galaxy merging are well known. Secular evolution is less studied and less well known. We are therefore still in the ``proof of concept'' phase in which it is useful to illustrate clearcut examples of the results of secular evolution.  This paper continues a series (see the above reviews and Kormendy \\& Cornell 2004) in which we illustrate the variety of disk-like features that define pseudobulges. ",
        "conclusions": "NGC 7690 (Sab) and NGC 4593 (SBb) provide clean examples of relatively early-type galaxies whose ``bulges'' are more disk-like than any elliptical galaxy.  In particular, elliptical galaxies are never flatter than axial ratio $\\simeq 0.4$ (Sandage et al.~1970; Binney \\& de Vaucouleurs 1981; Tremblay \\& Merritt 1995), whereas part (NGC 7690) or essentially all (NGC 4593) of the bulges of the present galaxies are as flat as their outer disks.  We conclude that both galaxies contain pseudobulges -- that is, high-density, central components that were made out of disk gas by secular evolution.  In NGC 7690, blue colors imply that star formation and hence pseudobulge growth are still in progress. In NGC 4593, gas appears currently to be accumulating in a ring that plausibly will form stars in the future. Our results are examples of the general conclusion (see KK for a review) that secular dynamical evolution occurs naturally and often in disk galaxies, whether (NGC 4593) or not (NGC 7690) an engine for the evolution is readily recognized.  To further investigate secular evolution, a desirable next step would be to quantify bar strengths by measuring~bar~torques, the ratio of the bar-induced, tangential force to the mean radial force as a function of radius (Buta \\& Block 2001; Laurikainen \\& Salo 2002; Block et al.~2001, 2004; Laurikainen et al.~2004).  This ratio can be as high as 0.6 in strong bars, emphasizing how efficient bars can be in redistributing angular momentum.  It would particularly be worthwhile to look for correlations between maximum bar torques and quantifiable consequences of secular evolution (e.{\\ts}g., ring-to-disk and pseudobulge-to-disk mass ratios) in a statistically representative sample of galaxies."
    },
    "0601/astro-ph0601116_arXiv.txt": {
        "abstract": "Variable A in M33 is a member of a rare class of highly luminous, evolved stars near the upper luminosity boundary that show sudden and dramatic shifts in apparent temperature due to the formation of optically thick winds in high mass loss episodes. Recent optical and infrared spectroscopy and imaging reveal that its ``eruption'' begun in $\\sim$1950 has ended, {\\it lasting $\\approx$ 45 yrs}. Our current observations show major changes in its wind from a cool, dense envelope to a much warmer state surrounded by low density gas with rare emission lines of Ca II, [Ca II] and K I. Its spectral energy distribution has unexpectedly changed, especially at the long wavelengths, with a significant decrease in its apparent flux, while the star remains optically obscured. We conclude that much of its radiation is now escaping out of our line of sight. We attribute this to the changing structure and distribution of its circumstellar ejecta corresponding to the altered state of its wind as the star recovers from a high mass loss event. ",
        "introduction": "Variable A in M33 is one of the highly luminous and unstable stars that define the upper luminosity limit in the Hertzsprung-Russell Diagram for evolved cool stars (see Humphreys and Davidson 1994). It was one of the original Hubble -- Sandage variables (Hubble and Sandage 1953) and at its maximum light in 1950 one of the visually brightest stars in M33. Its historical light curve shows a rapid decline from maximum by more than three magnitudes in less than a year followed by a brief recovery and a second decline. It had an intermediate F-type spectrum at maximum consistent with its observed colors. After its second decline to fainter than 18th magnitude (see light curves in Hubble and Sandage 1953 and Rosino and Bianchini 1973), no further observations were obtained until our optical and near infrared photometry beginning in 1977. In 1985-86 Var A had the spectrum of an M-type supergiant (Humphreys, Jones and Gehrz 1987, hereafter HJG), and its large infrared excess at 10$\\mu$m and spectral energy distribution showed that it was still as luminous ($5 \\times 10^{5} L_{\\odot}$) as it was at its maximum light in the visible. Thus its large photometric and spectral variations had occurred at nearly constant bolometric luminosity. We concluded that its cool M-type spectrum was produced in a pseudo-photosphere or optically thick wind formed during a high mass loss episode.  A recent spectrum, twenty years later, reveals another dramatic change. Its spectrum is now that of a much warmer star consistent with the warmer photosphere and colors of 50 years ago. In this paper we discuss its remarkable spectral changes reminiscent of the variability of $\\rho$ Cas (Lobel et al. 2003), but on much longer time scales. Variable A was  an obvious target for spectroscopy and imaging with the  Spitzer Space Telescope. Its resulting spectral energy distribution (.4{\\micron} to 24{\\micron}) also shows unexpected changes, especially at the long wavelengths, most likely corresponding to the changes in its wind and its impact on its circumstellar medium. In this paper we combine all of the available spectroscopy, multi-wavelength photometry and imaging to reconstruct the changes in its wind and circumstellar material over the past twenty years. In the next section we present our multi-wavelength observations. The current spectrum and its light curve and spectral energy distribution are described in \\S 3 and \\S 4.  In \\S 5 we discuss Var A's ``eruption'', its duration, and the changing structure of its wind and circumstellar nebula, and in the last section we review its relationship to other cool hypergiants and the possible origins of their instability. ",
        "conclusions": "Our current observations reveal some remarkable changes in the spectrum and energy distribution of Variable A. Although, we do not have a more complete spectroscopic and photometric record over the past 20 years,  our observations do allow us to put together a picture of a very luminous, unstable star experiencing  major changes in the structure of its wind and circumstellar material at the end of a  high mass loss event.  Variable A  still  had its cool, dense wind and was in a high mass loss phase 35 years after it rapid decline in 1951. Thus it was in ``eruption'' for at least that long. The photometric record also shows that Var A faded two magnitudes between 1954 and 1986 presumably due to the formation of dust, and may have created an additional two magnitudes of circumstellar extinction since then. Our recent spectrum 20 years later shows that the star's F-type photosphere has returned. When did this transition occur? The star has remained faint even though the spectrum has presumably recovered. The variation in its near-infrared photometry may correspond to changes in the wind accompanying the subsidence of its cool, dense wind phase. The 2MASS observations obtained in Dec. 1997 show that Var A  had faded significantly in the near-infrared. However, the JHK photometry from 1992 does not show any change from the earlier data in 1986. Assuming that the near-infrared variation is due to the wind in transition, then Var A's  eruption or dense wind stage may have lasted between 41 and 46 years!  How long the transition back to a warmer temperature may have taken is uncertain, but we note that the initial formation of the dense wind took approximately a year based on its 1950's light curve or at most two years if we include the brief recovery and second decline. Using the record of $\\rho$ Cas, which shows similar outbursts or shell episodes as an example, the onset and recovery timescales are comparable, and in $\\rho$ Cas they occur very rapidly, in only a couple of months. Thus many of the changes we discuss below may have occurred over only one to two years at most for the spectroscopic changes and perhaps up to five years (1992 to 1997) or so for the onset of the changes in the distribution and structure of the circumstellar material. The transition in the wind from an M-type false-photosphere to a warmer F-type star, could reasonably have occurred on even shorter timescales. Although we do not have a direct measurement of Var A's wind speed, the expansion velocities for the winds and ejecta in IRC+10420 and VY CMa are 40 -- 60 km s$^{-1}$ and 35 km s$^{-1}$ for the envelope expansion during $\\rho$ Cas's recent episiode. Assuming 50 km s$^{-1}$ for the wind speed, Var A's envelope could have made these transitions in as short  as  3/4's of a year, consistent with its variations during the early 1950's.  To explain Var A's observed energy distribution in 1986, HJG  proposed a simple model of a cool, dense false-photosphere with an obscuring torus and reflection nebulae at the poles. This model combined extinction by circumstellar dust with the blueing effect of scattering by dust grains and assumed that most of the flux was radiated in the mid-infrared. Most of the visible light was scattered and not viewed through the intervening material. In this model the visible light would very probably have been more highly polarized than the upper limit of 15\\% reported here. There was no optical polarimetry in 1986, so we cannot check this model or verify that there has been a change.   Despite its warmer apparent temperature and the corresponding change in the blue-visual energy distribution, Var A has not brightened in the visual, and  comparison with its maximum light implies $\\approx$ 4 magnitudes of circumstellar extinction currently in the line of sight.  HJG showed that the visual interstellar extinction for Var A was about 0.6 to 0.8 mag which is typical for stars in M33. With its current $\\bv$ color of 0.8 to 0.9 this suggests that its current true color is 0.6 to 0.4, appropriate for a late F-type star, and in comparison with its colors at maximum light ($\\sim$ .4), implies virtually no circumstellar  reddening   in the visual.  As we mentioned previously, the dust grains must be large enough to provide the neutral extinction in the visual which is also consistent with  the scattering requirements for the K I emission. Our polarimetry upper limit of 15\\% suggests that the optical light at present is not due mostly to scattered light. If we adopt 30\\% for the polarization of a star completely blocked by circumstellar dust and visible only in reflected light (Johnson and Jones 1991), then at most half the visible light from Var A is from scattered photospheric radiation. The polarimetry also rules out any current asymmetries such as bipolar lobes although reflection nebulosity in a more spherical distribution could still be present.   The most dramatic change in Var A's energy distribution is the apparent decrease in its total flux. Since it is highly unlikely that the total luminosity of the star would have declined by more than a magnitude, the energy must  now be escaping in some direction other than along our line of sight. Several massive stars are now known to have asymmetrical winds or bipolar outflows. The most notable is $\\eta$ Car with a latitude dependent wind (Smith et al 2002) that is both faster and denser at the poles, although $\\eta$ Car is a very different kind of star. However, given the evidence from other evolved, massive stars including the cool hypergiants like IRC+10420 and VY CMa, for irregularities and density variations in their circumstellar ejecta, it is likely that Var A's dusty shroud is not uniform or completely opaque in all directions. In the remaining discussion of  Var A's wind and circumstellar material we will assume that the missing radiation is escaping through large,  low density regions, even holes, in the obscuring material, and based on the timescales discussed above for the duration of the eruption and the recovery, we'll assume that Var A has been in this warmer state for approximately ten years.   HJG fit  Var A's 1986 mid-infrared flux  by a 370$\\arcdeg$ BB which implies a dusty zone or shell with a radius of $\\sim$ 400 AU from the star. The energy distribution had a significant near-infrared flux due to warmer dust closer to the star, so the dusty zone probably extended from 100 to 400 AU or so. Figure 4  shows that Var A slowly faded over 30 years due to increased circumstellar obscuration presumably as the amount and density of the obscuring material increased. HJG had estimated a mass loss rate of $2 \\times 10^{-4}$ M$_{\\odot}$ yr$^{-1}$ from Elitzur's (1981) formulation  assuming radiation coupling to the grains. But when Var A's optically thick wind ceased and quickly subsided back to a warmer state, then its mass loss rate would very likely have decreased as well. For example, the mass loss rates of the LBV's during their eruptions are typically 10 to 100 times that during their quiescent stage, and  the normal mass loss rates of $\\rho$ Cas and HR 8752, two hypergiants of comparable luminosity and temperature, are $\\sim$ $10^{-5}$ M$_{\\odot}$ yr $^{-1}$.  In addition to a less dense wind and a lower mass loss rate, Var A may also have a higher wind speed now. Wind speeds of 100 to 200 km s$^{-1}$ are typical of normal A to F-type supergiants. Following HJG's procedure, using Var A's apparent luminosity inferred from its current 10$\\micron$ flux, and assuming that the optical depth is near one in our line of sight, we find a current mass loss rate of $\\sim$ $6.7  \\times 10^{-5}$ M$_{\\odot}$ yr$^{-1}$ for a 50 km s$^{-1}$ wind. But if the wind speed is higher, the mass loss rate will be $\\sim$ 2 -- 3 $\\times 10^{-5}$ M$_{\\odot}$ yr$^{-1}$.  In the approximately 10 years since its transition, Var A's lower density and possibly faster wind would have reached 100 to 300 AU, the region of the dusty zone. Consequently, the dusty material may not be replenished as efficiently as in the previous dense wind, higher mass loss stage. So as the dusty zone has continued to expand during this same period, the lower density regions or gaps will also have enlarged allowing more radiation to escape. If the dusty distribution is flattened, as HJG suggested, this is most likely to occur at the poles.  Therefore we propose that large dusty condensations in a flattened distribution, possibly a torus, currently obscure our direct view of Var A. The K I emission lines are produced by resonance scattering in this dusty zone. In addition, a very low density gas or wind responsible for  the hydrogen, Ca II, and peculiar [Ca II] emission fills the region between the star's photosphere and the dusty zone.  More than half  of the radiation is escaping from our line of sight, presumably from low density regions, all or most of which must be out of our line of sight. Furthermore, the polarimetry measurements rule out a strongly bipolar structure with asymmetrically scattered light. This then suggests a unique geometry even if the gaps are restricted to the polar regions.  The dusty zone must either be nearly aligned with our line of sight in such a way to also block most of the escaping radiation from being reflected back into our line of sight or else there is very little reflecting nebulosity."
    },
    "0601/astro-ph0601320_arXiv.txt": {
        "abstract": "N-Body simulations are a very important tool in the study of formation of large scale structures. Much of the progress in understanding the physics of galaxy formation and comparison with observations would not have been possible without N-Body simulations. Given the importance of this tool, it is essential to understand its limitations as ignoring these can easily lead to interesting but unreliable results. In this paper we study the limitations due to the finite size of the simulation volume. We explicitly construct the correction term arising due to a finite box size and study its generic features for clustering of matter and also on mass functions. We show that the correction to mass function is maximum near the scale of non-linearity, as a corollary we show that the correction to the number density of haloes of a given mass changes sign at this scale; the number of haloes at small masses is over estimated in simulations. This over estimate results from a delay in mergers that lead to formation of more massive haloes. The same technique can be used to study corrections to other physical quantities. The corrections are typically small if the scale of non-linearity is much smaller than the box-size. However, there are some cases of physical interest in which the relative correction term is of order unity even though a simulation box much larger than the scale of non-linearity is used. Within the context of the concordance model, our analysis suggests that it is very difficult for present day simulations to resolve mass scales smaller than $10^2$~M$_\\odot$ accurately and the level of difficulty increases as we go to even smaller masses, though this constraint does not apply to multi-scale simulations. ",
        "introduction": "Large scale structures like galaxies and clusters of galaxies are believed to have formed by gravitational amplification of small perturbations. For an overview and original references, see, e.g., \\citet{1980lssu.book.....P,1999coph.book.....P,2002tagc.book.....P,2002PhR...367....1B}. Initial density perturbations were present at all scales that have been observed \\citep{2003ApJS..148..175S,2003MNRAS.346...78H,2004ApJ...607..655P}. Understanding evolution of density perturbations for such initial conditions is essential for the study of formation of galaxies and large scale structures. The equations that describe the evolution of density perturbations in an expanding universe have been known for a long time \\citep{1974A&A....32..391P} and these are easy to solve when the amplitude of perturbations is small. These equations describe the evolution of density contrast defined as $\\delta(\\mathbf r, t) = (\\rho(\\mathbf r, t) - \\bar\\rho(t))/\\bar\\rho(t)$. Here $\\rho(\\mathbf r, t)$ is the density at point $\\mathbf r$ and time $t$, and $\\bar\\rho$ is the average density in the universe at time $t$. These are densities of non-relativistic matter, the component that clusters at all scales and is believed to drive the formation of large scale structures in the universe. Once density contrast at relevant scales becomes large, i.e., $|\\delta| \\geq 1$, the perturbation becomes non-linear and coupling with perturbations at other scales cannot be ignored. The equation for evolution of density perturbations cannot be solved for generic perturbations in this regime. N-Body simulations \\citep{1998ARA&A..36..599B,1997Prama..49..161B,2005CSci...88.1088B} are often used to study the evolution in this regime. Alternative approaches can be used if one requires only a limited amount of information and in such a case either quasi-linear approximation schemes \\citep{1970A&A.....5...84Z,1989MNRAS.236..385G,1992MNRAS.259..437M,1993ApJ...418..570B,1994MNRAS.266..227B,1995PhR...262....1S,1996ApJ...471....1H,2002PhR...367....1B} or scaling relations \\citep{1977ApJS...34..425D,1991ApJ...374L...1H,1995MNRAS.276L..25J,2000ApJ...531...17Ka,1998ApJ...508L...5M,1994MNRAS.271..976N,1996ApJ...466..604P,1994MNRAS.267.1020P,1996MNRAS.278L..29P,1996MNRAS.280L..19P,2003MNRAS.341.1311S} suffice.  In cosmological N-Body simulations, we simulate a representative region of the universe. This is a large but finite volume and periodic boundary conditions are often used. Almost always, the simulation volume is taken to be a cube. Effect of perturbations at scales smaller than the mass resolution of the simulation, and of perturbations at scales larger than the box is ignored. Indeed, even perturbations at scales comparable to the box are under sampled.  It has been shown that perturbations at small scales do not influence collapse of perturbations at much larger scales in a significant manner \\citep{1974A&A....32..391P,1985ApJ...297..350P,1991MNRAS.253..295L,1997MNRAS.286.1023B,1998ApJ...497..499C} if we study the evolution of the correlation function or power spectrum at large scales due to gravitational clustering in an expanding universe. This is certainly true if the scales of interest are in the non-linear regime \\citep{1997MNRAS.286.1023B}. Therefore we may assume that ignoring perturbations at scales much smaller than the scales of interest does not affect results of N-Body simulations. However, there may be other effects that are not completely understood at the quantitative level \\citep{2004astro.ph..8429B} even though these have been seen only in somewhat artificial situations.  Perturbations at scales larger than the simulation volume can affect the results of N-Body simulations. Use of periodic boundary conditions implies that the average density in the simulation box is same as the average density in the universe, in other words we ignore perturbations at the scale of the simulation volume (and at larger scales). Therefore the size of the simulation volume should be chosen so that the amplitude of fluctuations at the box scale (and at larger scales) is ignorable. If the amplitude of perturbations at larger scales is not ignorable then clearly the simulation will not be a faithful representation of the model being studied. It is not obvious as to when fluctuations at larger scales can be considered ignorable, indeed the answer to this question depends on the physical quantity of interest, the model being studied and the specific length/mass scale of interest as well.  The effect of a finite box size has been studied using N-Body simulations and the conclusions in this regard may be summarised as follows. \\begin{itemize} \\item If the amplitude of density perturbations around the box scale is small ($\\delta < 1$) but not much smaller than unity, simulations underestimate the correlation function though the number density of small mass haloes does not change by much \\citep{1994ApJ...436..467G,1994ApJ...436..491G}. In other words, the formation of small haloes is not disturbed but their distribution is affected by non-inclusion of long wave modes. \\item In the same situation, the number density of massive haloes drops significantly \\citep{1994ApJ...436..467G,1994ApJ...436..491G,2005MNRAS.358.1076B}. \\item Effects of a finite box size modify values of physical quantities like the correlation function even at scales much smaller than the simulation volume \\citep{2005MNRAS.358.1076B}. \\item The void spectrum is also affected by finite size of the simulation volume if perturbations at large scales are not ignorable \\citep{1992ApJ...393..415K}. \\item It has been shown that properties of a given halo can change significantly as the contribution of perturbations at large scales is removed to the initial conditions but the distribution of most internal properties remain unchanged \\citep{2005astro.ph.12281P}. \\end{itemize}  In some cases, one may be able to devise a method to ``correct'' for the effects of a finite box-size \\citep{1994A&A...281..301C}, but such methods cannot be generalised to all statistical measures or physical quantities.  The effects of perturbations at scales larger than the box size can be added using MAP (Mode Adding Procedure) after a simulation has been run \\citep{1996ApJ...472...14T}. This method makes use of the fact that if the box size is chosen to be large enough then the contribution of larger scales can be incorporated by adding displacements due to the larger scales independently of the evolution of the system in an N-Body simulation. The motivation for development of such a tool is to enhance the range of scales over which results of an N-Body simulation can be used by improving the description at scales comparable to the box size. Such an approach ignores the coupling of large scale modes with small scale modes and this again brings up the issue of what is a large enough scale for a given model such that the effects of mode coupling can be ignored. Large scales contribute to displacements and velocities, and variations in density due to these scales modify the rate of growth for small scales perturbations \\citep{1997MNRAS.286...38C}.  Effects of a finite box size modify values of physical quantities even at scales much smaller than the simulation volume \\citep{2005MNRAS.358.1076B} (BR05, hereafter). In BR05, we suggested use of the fraction of mass in collapsed haloes as an indicator of the effect of a finite box size. We found that if the simulation volume is not large enough, the fraction of mass in collapsed haloes is underestimated. As the collapsed fraction is less sensitive to box-size as compared to measures of clustering, several other statistical indicators of clustering to deviate significantly from expected values in such simulations. A workaround for this problem was suggested in the form of an ensemble of simulations to take the effect of convergence due to long wave modes into account \\citep{2005ApJ...634..728S}, the effects of shear due to long wave modes are ignored here. It has also been shown that the distribution of most internal properties of haloes, e.g., concentration, triaxiality and angular momentum do not change considerably with the box size even though properties of a given halo may change by a significant amount \\citep{2005astro.ph.12281P}.  There is a clear need to develop a formalism for estimating the effect of perturbations at large scales on a variety of physical quantities. Without such a formalism, we cannot decide in an objective manner whether a simulation box size is sufficiently large or not. In this paper we generalise the approach suggested in BR05 and write down an explicit correction term for a number of statistical indicators of clustering. This approach allows us to study generic properties of the expected correction term in any given case, apart of course from allowing us to evaluate the magnitude of the correction as compared to the expected value of the given statistical indicator. We apply this technique to mass functions in this paper. ",
        "conclusions": "Conclusions of this work may be summarised as follows. \\begin{itemize} \\item We have presented a formalism that can be used to estimate the deviations of cosmological N-Body simulations from the models being simulated due to the use of a finite box size. These deviations/errors are independent of the specific method used for doing simulations. \\item For a given model, the deviations can be expressed as a function of the scale $r$ of interest and $\\lbx$, the box size of simulations. \\item We have applied the formalism to study deviations in {\\it rms}\\/ fluctuations in mass in the initial conditions. \\item We find that the errors are small except for models where the slope of the power spectrum is close to $-3$ at scales of interest. \\item The errors in case of the $\\Lambda$CDM model are significant if the scale of interest is smaller than a kpc even if simulations as large as the Millennium simulation \\citep{2005Natur.435..629S} are used. \\item We have studied errors in mass function in the Press-Schechter model, as well as other models. \\item The main error due to a finite box size is that the number of high mass haloes is under estimated. \\item The number of low mass haloes is over predicted in simulations if the box size effects are important. This happens as low mass haloes do not merge to form the (missing) high mass haloes. \\item We have verified these trends using N-Body simulations. \\end{itemize}  We note that it is extremely important to understand the sources of errors in N-Body simulations and the magnitude of errors in quantities of physical interest. N-Body simulations are used to make predictions for a number of observational projects and also serve as a test bed for methods. In this era of ``precision cosmology'', it will be tragic if simulations prove to be a weak link. We would like to note that our results apply equally to all methods of doing cosmological N-Body simulations, save those where techniques like MAP are used to include the effects of scales larger than the simulation volume.  The method for estimating errors due to a finite box-size described in this paper can be used for several physical quantities. In this paper we have used the method to study errors in clustering properties and mass functions. We are studying the effect of finite box size on velocity fields and related quantities, the results will be presented in a later publication."
    },
    "0601/astro-ph0601099_arXiv.txt": {
        "abstract": "We revisit the constraint on the primordial helium mass fraction $Y_p$ from observations of cosmic microwave background (CMB) alone. By minimizing $\\chi^2$ of recent CMB experiments over 6 other cosmological parameters, we obtained rather weak constraints as $0.17\\le Y_p \\le 0.52$ at 1$\\sigma$ C.L. for a particular data set.  We also study the future constraint on cosmological parameters when we take account of the prediction of the standard big bang nucleosynthesis (BBN) theory as a prior on the helium mass fraction where $Y_p$ can be fixed for a given energy density of baryon. We discuss the implications of the prediction of the standard BBN on the analysis of CMB. \\vspace{1cm} ",
        "introduction": "\\setcounter{equation}{0}  Recent precise cosmological observations such as WMAP \\cite{Bennett:2003bz} push us toward the era of so-called precision cosmology.  In particular, the combination of the data from cosmic microwave background (CMB), large scale structure, type Ia supernovae and so on can severely constrain the cosmological parameters such as the energy density of baryon, cold dark matter and dark energy, the equation of state for dark energy, the Hubble parameter, the amplitude and the scale dependence of primordial fluctuation.  Among the various cosmological parameters, the primordial helium mass fraction $Y_p$ is the one which has been mainly discussed in the context of big bang nucleosynthesis (BBN) but not that of CMB so far. One of the reason is that the primordial helium abundance has not been considered to be well constrained by observations of CMB since its effects on the CMB power spectrum is expected to be too small to be measured.  However, since now we have very precise measurements of CMB,  we may have a chance to constrain the primordial helium mass fraction from CMB observations. Since the primordial helium mass fraction can affect the number density of free electron in the course of the recombination history, the effects of $Y_p$ can be imprinted on the power spectrum of CMB. Recently, some works along this line have been done by two different groups \\cite{Trotta:2003xg,Huey:2003ef}, which have discussed the constraints on $Y_p$ from current observations of CMB.  In fact they claim different bounds on the primordial helium mass fraction, especially in terms of its uncertainty: the author of Ref.~\\cite{Trotta:2003xg} obtained $ 0.160 \\le Y_p \\le 0.501$, on the other hand the authors of Ref.~\\cite{Huey:2003ef} got $ Y_p =0.250^{+0.010}_{-0.014}$ at 1$\\sigma$ confidence level.  It should be noticed that the latter bound is much more severe than that of the former. If the helium mass fraction is severely constrained by CMB data, it means that the CMB power spectrum is sensitive to the values of $Y_p$.  In such a case, the prior on $Y_p$ should be important to constrain other cosmological parameters too and the usual fixing of $Y_p=0.24$ in CMB power spectrum calculations might not be a good assumption. Especially, analyses like Refs.~\\cite{Cyburt:2003fe,Cuoco:2003cu,Coc:2003ce,Cyburt:2004cq,Serpico:2004gx} predict light element abundances including $^4$He from the baryon density which is obtained from the CMB data sets with the analysis fixing the value of $Y_p$. Such procedure is only valid when $Y_p$ is not severely constrained by CMB. Thus it is very important to check the CMB bound on $Y_p$.  One of the main purpose of the present paper is that we revisit the constraint on $Y_p$ from observations of CMB alone with a different analysis method from Markov chain Monte Carlo (MCMC) technique which is widely used for the determination of cosmological parameters and adopted in Refs.~\\cite{Trotta:2003xg,Huey:2003ef}.  In this paper, we calculate $\\chi^2$ minimum as a function of $Y_p$ and derive constraints on $Y_p$. We adopt the Brent method of the successive parabolic interpolation to minimize $\\chi^2$ varying 6 other cosmological parameters of the $\\Lambda$CDM model with the power-law adiabatic primordial fluctuation. We obtain the constraint on $Y_p$ by this method and compare it with previously obtained results.  We also study the constraint on $Y_p$ from future CMB experiment. A particular emphasis is placed on investigating the role of the standard BBN theory. Since the primordial helium is synthesized in BBN, once the baryon-to-photon ratio is given, the value of $Y_p$ is fixed theoretically. Thus, using this relation between the baryon density and helium abundance, we do not have to regard $Y_p$ as an independent free parameter when we analyze CMB data. We study how the standard BBN assumption on $Y_p$ affects the determination of other cosmological parameters in the future Planck experiment using the Fisher matrix analysis.  The structure of this paper is as follows. In the next section, we briefly discuss the effects of the helium mass fraction on the CMB power spectrum, in particular its effects on the change of the structure of the acoustic peaks. Then we study the constraint on the primordial helium mass fraction from current observations of CMB using the data from WMAP, CBI, ACBAR and BOOMERANG. In section 4, we discuss the expected constraint on $Y_p$ from future CMB observation of Planck and also study how the standard BBN assumption on $Y_p$ can affect the constraints on cosmological parameters. The final section is devoted to the summary of this paper. ",
        "conclusions": "We revisited the constraint on the primordial helium mass fraction $Y_p$ from current observations of CMB. Some authors have already studied the constraint \\cite{Trotta:2003xg,Huey:2003ef}, however their results were different especially in terms of the uncertainty.  One of the main purpose of the present paper is to study the constraints on $Y_p$ from current observations adopting a different analysis method.  Instead of MCMC method which was adopted by the authors of Refs.~\\cite{Trotta:2003xg,Huey:2003ef} to obtain the constraint, here we adopted a $\\chi^2$ minimization by a nested grid search. We did not obtain a severe constraint in agreement with Ref.~\\cite{Trotta:2003xg}. Using the data from WMAP, CBI and ACBAR as well as recent BOOMERANG data, we get 1$\\sigma$ constraint as $ 0.25 \\le Y_p \\le 0.54$ and $0.17 \\le Y_p \\le 0.52$  for the cases with and without the data from BOOMERANG. It might be interesting to note that usual assumption of $Y_p=0.24$ is not in the 1$\\sigma$ error range of BOOMERANG combined analysis but it is not of high significance at this stage so we can safely assume $Y_p=0.24$ for current CMB data analysis.  We also studied the future constraint from CMB on $Y_p$ taking account of the standard BBN prediction as a prior on $Y_p$.  Although we cannot obtain a severe constraint at present, observations of CMB can be much more precise in the future. Thus we may have a chance to obtain a precise measurement of $Y_p$ from upcoming CMB experiments. On the other hand, since the primordial helium has been formed during the time of BBN, once the baryon-to-photon ratio is given, the value of $Y_p$ can be evaluated theoretically assuming the standard BBN. Thus, in this case, we do not have to assume $Y_p$ as an independent free parameter when we analyze CMB data.  We studied how such BBN theory prior on $Y_p$ affects the determination of other cosmological parameters in the future Planck experiment. We evaluated the uncertainties for the case with $Y_p$ being an independent free parameter and $Y_p$ being fixed for a given $\\omega_b$ using the BBN relation.  We showed that the BBN prior improves the constraints on other cosmological parameters by a factor of $\\mathcal{O}(1)$ and also it induces some correlations among the parameters which appear in the BBN relation.  As shown in Fig.~\\ref{fig:future}, as far as we consider the standard scenario of cosmology, the helium mass fraction can be fixed for CMB analysis even in the future experiments since we can expect that the constraint from Planck is much weaker than the uncertainty of the theoretical calculation of the standard BBN.  However, it is worthwhile to do CMB analysis treating $Y_p$ as a free parameter and measure the helium mass fraction independently from the baryon density since it provides a consistency test for the standard BBN theory (of course, measurements of primordial light element abundances by astrophysical means provide further consistency tests).  By checking the robustness of the consistency from various observations, the golden age of precision cosmology can push us toward the accurate understanding of the universe.  \\bigskip \\noindent {\\bf Acknowledgment:} We acknowledge the use of CMBFAST \\cite{Seljak:1996is} package for our numerical calculations. The work of T.T. is supported by Grand-in-Aid for JSPS fellows."
    },
    "0601/astro-ph0601266_arXiv.txt": {
        "abstract": "We investigate the possibility that part of the dark matter is not made out of the usual cold dark matter (CDM) dustlike particles, but is under the form of a fluid of strings with barotropic factor $w_s = -1/3$ of cosmic origin. To this aim, we split the dark matter density parameter in two terms and investigate the dynamics of a spatially flat universe filled with baryons, CDM, fluid of strings and dark energy, modeling this latter as a cosmological constant or a negative pressure fluid with a constant equation of state $w < 0$. To test the viability of the models and to constrain their parameters, we use the Type Ia Supernovae Hubble diagram and the data on the gas mass fraction in galaxy clusters. We also discuss the weak field limit of a model comprising a significant fraction of dark matter in the form of a fluid of strings and show that this mechanism makes it possible to reduce the need for the elusive and up to now undetected CDM. We finally find that a model comprising both a cosmological constant and a fluid of strings fits very well the data and eliminates the need of phantom dark energy thus representing a viable candidate to alleviate some of the problems plaguing the dark side of the universe. ",
        "introduction": "Soon after the discovery of cosmic acceleration from the Hubble diagram of the high redshift Type Ia Supernovae (SNeIa) \\cite{SNeIa,Riess04}, a strong debate arose in the scientific community about the origin of this unexpected result. An impressive flow of theoretical proposals have appeared, while the observational results were constantly providing more and more evidences substantiating the emergence of a new cosmological scenario. The anisotropy spectrum of the cosmic microwave background radiation (CMBR) \\cite{CMBR,WMAP,VSA}, the matter power spectrum determined from the clustering properties of the large scale distribution of galaxies \\cite{LSS} and the data on the Ly$\\alpha$ emitting regions \\cite{Lyalpha} all provide indications that the universe have to be described as a spatially flat manifold where matter and its fluctuations are isotropically distributed and represent only about $30\\%$ of the overall content. In order to fill the gap and drive the acceleration, a dominant contribute from a homogeneously distributed negative pressure fluid has been invoked. Usually referred to as {\\it dark energy}, the nature and the nurture of this mysteryous component represent a new and fascinating conundrum for theoreticians.  While the models proposed to explain this puzzle increase day by day, the most simple answer is still the old Einstein cosmological constant $\\Lambda$. Although being the best fit to a wide set of different astrophysical observations \\cite{Teg03,Sel04}, it is nevertheless plagued by two evident shortcomings, namely the {\\it cosmological constant problem} and the {\\it coincidence problem}. A possible way to overcome these problems invokes replacing $\\Lambda$ with a scalar field (dubbed {\\it quintessence}) evolving down a suitably chosen self interaction potential \\cite{QuintFirst}. Although solving the cosmological constant problem, quintessence does not eliminate the coincidence one since too severe constraints on the potential seem to be needed thus leading to the {\\it fine tuning problem} \\cite{QuintRev}.  The ignorance of the fundamental physical properties of both dark energy and dark matter has motivated a completely different approach to the problem of cosmic acceleration relying on modification of the matter equation of state (EoS). Referred to as {\\it unified dark energy} (UDE) models, these proposals resort to a single fluid with exotic EoS as the only candidate to both dark matter and dark energy thus automatically solving the coincidence problem. The EoS is then tuned such that the fluid behaves as dark matter at high energy density and quintessence (or $\\Lambda$) at the low energy limit. Interesting examples are the Chaplygin gas \\cite{Chaplygin}, the tachyonic field \\cite{tachyon} and the Hobbit model \\cite{Hobbit}.  It is worth noting that observations only tell us that the universe is accelerating, but they are not direct evidences for new fluids or modifications of the usual matter properties. It is indeed possible to consider cosmic acceleration as the first signal of the breakdown of the laws of physics as we know them. As a consequence, one has to to give off the standard Friedman equations in favour of a generalized version of them arising from some more fundamental theory. Interesting examples of this kind are the Cardassian expansion \\cite{Cardassian} and the Dvali\\,-\\,Gabadadze\\,-\\,Porrati (DGP) gravity \\cite{DGP} both related to higher dimensional braneworld theories. In the same framework, one should also give off the Einsteinian general relativity and turn to fourth order theories of gravitation replacing the Ricci scalar curvature $R$ in the gravity Lagrangian with a generic function $f(R)$ that have been formulated both in the metric \\cite{capozcurv,MetricRn,cct} and Palatini approach \\cite{PalRn,lnR,ABF04,ACCF} providing a good fit to the data in both cases \\cite{noiijmpd,CCF04}. Actually, it is worth noting that it has been recently demonstrated that, under quite general conditions, it is possible to find a $f(R)$ theory that predicts the same dynamics of a given quintessence model.  Although resorting to modified gravity theories is an interesting and fascinating approach, it is worth exploring other possibilities in the framework of standard general relativity. Indeed, all the approaches we have described are mainly interested in solving the dark energy puzzle, while little is said about the dark matter problem. It is worth remembering that dark matter is usually invoked because of the need of a source of gravitational potential other than the visible matter. Considered from this point of view, it is worth wondering whether dark matter could be replaced by a different mechanism that is able to give the same global effect. Moreover, such a mechanism must not alter the delicate balance between dark matter and dark energy that is needed to explain observations. Indeed, if we abruptly reduce the dark matter content of the universe without altering neither the dark energy term nor the background fundamental properties, we are not able to fit the available astrophysical data. Therefore, it is mandatory to test any proposed mechanism both at galactic and cosmological scales.  In a series of interesting papers \\cite{letelier}, Letelier investigated the consequences of changing the properties of the right hand side of the Einstein equations adopting a fluid of strings as source term rather than the usual dust matter. Since such strings are not observed at the present time, it seems meaningful to extend the concept of dust clouds and perfect fluid referred to point particles to the case of strings. In particular, Letelier was able to find exact solutions for the case of a spherically symmetric fluid of strings. It is worth noting that such strings could be of cosmological origin \\cite{Vil84} and have thus to be included in the energy budget when investigating the dynamics of the universe. It is important to stress, however, that the strings we are referring to have {\\it finite lenght} so that the results obtained for a network of cosmic strings of infinite length cannot be extended to the strings considered by Letelier. In this sense, the fluid of finite length strings we are considering represents a generalization of the dust matter. While in this latter case, the matter particles are considered as pointlike, in the case of a fluid of strings\\footnote{Hereafter, by {\\it fluid of strings} we mean a {\\it fluid of finite length strings}.} the elementary constituents are one dimensional objects with finite length. A fluid of strings has a profound impact at galactic scales. Indeed, assuming that the string transverse pressure was proportional to its energy density, Soleng \\cite{soleng} has demonstrated that the force law is altered thus offering the possibility of solving the problem of the flatness of spiral galaxy rotation curves \\cite{rc} in a way similar to the MOND proposal \\cite{mond}.  Motivated by these considerations, we explore here the possibility that a part (if not all) of the dark matter may be replaced by a fluid of strings whose effective gravitational action may be considered as the source of the gravitational potential needed to flatten the rotation curves. To this aim, we consider cosmic strings as components of such fluid so that its e.o.s. may be simply parametrized by a constant barotropic factor $w_s = -1/3$. Before discussing the impact at galaxy scales, it is preliminarily needed to investigate the effects at cosmological scales. We thus fit different cosmological models, both with and without such a component, to the available astrophysical data in order to test the viability of our proposal and explore if and how the constraints on the model parameters are affected by the presence of a fluid of strings.  The paper is organized as follows. The models we discuss are described in Sect.\\,2, while the matching with observations is presented in Sect.\\,3 where we also compare the different models in terms of the information criteria parameters. Sect.\\,4 is devoted to the weak energy limit of models comprising standard matter embedded in a fluid of strings and show how the corresponding modified gravitational potential could help in reducing the need for CDM. A summary of the results and of their implications are presented in the concluding Sect.\\,5 ",
        "conclusions": "Shedding light on the dark side of the universe is a very difficult, but also very attractive challenge of modern cosmology. The nature and the fundamental properties of the two main ingredients of the cosmic pie, namely the dark energy and the dark matter, are still substantially unknown and it is, indeed, this wide ignorance that justifies and motivates the impressive amount of theoretical models proposed to explain the observed astrophysical evidences. Moving in this framework, we have considered the dark matter as made out not only of massive dustlike CDM particles, but also of a fluid of strings of cosmic origin with an equation of state $w_s = -1/3$. Starting from this idea, we have considered four cosmological models comprising four components, namely dustlike baryons and CDM, fluid of strings and dark energy with constant barotropic factor $w$. Two of these four models ($\\Lambda$SDM and QSDM) have a non vanishing fraction of dark matter in the form of a fluid of strings, while in two models ($\\Lambda$CDM and $\\Lambda$SDM) the energy budget is dominated by the $\\Lambda$ term. Our main results are briefly outlined as follows.  \\begin{enumerate}  \\item{All the models are able to fit the data on the SNeIa Hubble diagram and the gas mass fraction in galaxy clusters with very good accuracy. In particular, it is remarkable that the total dark matter density parameter $\\Omega_{DM} = \\Omega_{CDM} + \\Omega_s$ is very well constrained and turns out to be the same in all models. When the assumption $w = -1$ is relaxed, the dark energy barotropic factor is constrained to be in the region $w < -1$ so that phantom like models are clearly preferred with a disturbing violation of the strong energy condition. It is worth noting that present data do not require phantom dark energy since they can be equally well fit by models with the cosmological constant $\\Lambda$ driving the accelerated expansion. Discriminating among the different possibilities will need a large sample of high redshift SNeIa such as those that should be available with the SNAP satellite mission \\cite{SNAP}.}  \\item{According to both the AIC and BIC, the QCDM model is statistically preferred over the other considered possibilities and this is not an unexpected result. However, this is obtained to the price of admitting phantom dark energy which is affected by serious theoretical difficulties. On the other hand, the $\\Lambda$SDM model is preferred over the popular concordance $\\Lambda$CDM scenario and is only slightly disfavoured with respect to the QCDM one. The good fit to the data and the graceful feature of avoiding to enter the realm of ghosts makes this model a good compromise between observations and theory and we therefore consider it as our final best choice.}  \\item{The $\\Lambda$SDM model predicts that a significant fraction ($\\varepsilon \\simeq 59\\%$) of the dark matter is made out by a fluid of strings so that the standard matter density parameter $\\Omega_b + \\Omega_{CDM}$ is only half of the fiducial value in the concordance scenario (0.15 vs 0.30). Since we may assume that the fluid of strings is massless (or nearly so), we should expect a corresponding decrease of the mass of galactic dark haloes. If the gravitational potential is still Newtonian, decreasing the CDM halo mass should lead to lower values of the circular velocity in the outer dark matter dominated regions of galaxies. This is not the case since, in the weak field limit, the $\\Lambda$SDM model gives rise to a modification of the gravitational potential. As a result, the circular velocity due to a mass $M(r)$ is higher than in the classical case so that less massive haloes are necessary to give the observed values of $v_c(r)$. Moreover, a very qualitative calculation suggests that the typical value of the scalelength over which deviations from Newtonian formulae cannot be neglected is sufficiently high that the inner luminous matter dominated rotation curve is unaltered.}  \\end{enumerate}  These encouraging results motivate further studies of the $\\Lambda$SDM model. To this end, there are two different routes connected to two different features of the model which can be followed.  First, because of its scaling with the redshift as $\\rho_s \\propto (1 + z)^2$, that is intermediate between that of CDM and that of $\\Lambda$, a new era dominated by a fluid of strings is predicted in the expansion history of the universe. It is thus worth investigating how this imprints on the CMBR anisotropies in order to see whether the spectrum measured by WMAP is still accurately reproduced. To this regard, it is worth noting that the attempts recently made to constrain the cosmic strings contribution to the CMBR spectrum \\cite{Fraisse} may not be extended to our case since they refer to a network of cosmic strings rather than clouds of finite length strings. Less theoretically demanding, but more observationally ambitious is the possibility to test the proposed scenario on the basis of the transition redshift $z_T$. As Table 1 shows, for the $\\Lambda$SDM model, $z_T$ is significantly higher than in the other cases so that a model independent estimate of this quantity could be a powerful discriminating tool.  One of the most peculiar features of the $\\Lambda$SDM model is the modified gravitational potential in Eq.(\\ref{eq: phiext}) leading to the corrected circular velocity in Eq.(\\ref{eq: vcstrings}). Having been obtained in the weak field limit, such correction should be tested at the scale of galaxies and clusters of galaxies thus offering the possibility to test the model at a very different level. To this aim, one should try fitting the rotation curve of spiral galaxies to see whether the problem of their flatness could be solved in this framework. Moreover, it is interesting to check how much the halo mass is reduced and to compare the reduction with respect to the classical Newtonian estimates with the decreasing of $\\Omega_{CDM}$ obtained above. To this aim, low surface brightness (LSB) galaxies are ideal candidates since they are likely dark matter dominated so that systematic uncertainties on the luminous matter modelling have only a minor impact on the fitting procedure. Moreover, the stellar mass\\,-\\,to\\,-\\,light ratio of LSB galaxies is well constrained so that we may fix this quantity thus decreasing the degeneracy among the other parameters. Useful samples of LSB galaxies with detailed measurements of the rotation curve are yet available (see, for instance, \\cite{dBB02}) so that this kind of test may be easily implemented. In this same framework, it is also interesting to consider the velocity dispersion curves in elliptical galaxies where recent studies seem to indicate a dark matter deficit \\cite{Cap}.  Changing the gravitational potential does not only alter galaxies rotation curve, but also affects the clustering properties and thus leads to a different matter power spectrum. It is thus interesting to compare the predicted power spectrum with those measured from the SDSS galaxies in order to check the validity of the $\\Lambda$SDM model. A similar comparison has been recently performed by Shirata et al. \\cite{SSYS04} for two phenomenological modifications of the law of gravity. We stress, however, that their approach is purely empirical and, furthermore, assumes that the universe can still be described at large scales with the $\\Lambda$CDM model. Since in order to compute the power spectrum, one also needs the background Hubble parameter, it is important to use an expression for $H(z)$ that is consistent with the proposed modification of gravity. For the $\\Lambda$SDM model considered here, all the ingredients are at disposal so that a coherent calculation can be performed. It is worth noting that such a test is the only one capable of probing the model both at the galactic (through the gravitational potential) and cosmological (because of the use of $H$) scales at the same time.  As a concluding general remark, we would like to stress the need for tackling the dark matter and dark energy problem together taking care of what is the effect at the galaxy scale of any modification of the fundamental properties of one of these two components. In our opionion, this could be a valid approach in elucidating the problems connected to the dark side of the universe."
    },
    "0601/hep-ph0601014_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}} \\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{% PBH and DM from cosmic necklaces } \\author{% Tomohiro Matsuda\\footnote{E-mail:matsuda@sit.ac.jp} } \\address{% $^1$Theoretical Physics Group, Saitama Institute of Technology, Saitama 369-0293, Japan }  \\abstract{Cosmic strings in the brane Universe have recently gained a great interest. I think the most interesting story is that future cosmological observations distinguish them from the conventional cosmic strings. If the strings are the higher-dimensional objects that can (at least initially) move along the compactified space, and finally settle down to (quasi-)degenerated vacua in the compactified space, then kinks should appear on the strings, which interpolate between the degenerated vacua. These kinks look like ``beads'' on the strings, which means that the strings turn into necklaces. Moreover, in the case that the compact manifold is not simply connected, the string loop that winds around a non-trivial circle is stable due to the topological reason. Since the existence of degenerated vacua and a non-trivial circle is the common feature of the brane models, it is important to study cosmological constraints on the cosmic necklaces and their stable winding states in the brane Universe.} ",
        "introduction": "In this talk we will explain the cosmological consequences of the production of Dark Matter(DM) and Primordial Black Hole(PBH) from the loops of the cosmic necklaces. To begin with, I think it is fair to explain why necklaces\\cite{vilenkin_book} are produced in brane models, since in many papers it is discussed that ``only strings are produced in the brane Universe''\\cite{tutorial}. Of course I think this claim is not wrong, however somewhat misleading for non-specialists. To explain what is misleading in the ``standard scenario'', we have a figure in Fig.\\ref{fig:fig0}. In general, the distance between branes may appear in the four-dimensional effective action as a Higgs field of the effective gauge dynamics. At least in this case, it is natural to consider the cosmological defects coming from the spatial variation of the Higgs field, which corresponds to the ``deformation'' of the branes\\cite{matsuda_deformation}. Is the spatial variation of the Higgs field unnatural in the brane Universe? The answer is, of course, no. One should therefore consider at least two different kinds of defects in brane models: One is induced by the brane creation that is due to the spatial variation of the tachyon condensation, while the other is induced by the brane deformation that is due to the spatial variation of the brane distance. Along the line of the above arguments, it is possible to construct Q-ball's counterpart in brane models\\cite{matsuda_Q-balls}, which can be distinguished from the conventional Q-balls by their decay process. We therefore have an expectation that strings can be distinguished from the conventional ones, if one properly considers their characteristic features.  Now let us discuss about the validity of the conventional Kibble mechanism. Of course the Kibble mechanism is an excellent idea that explains the nature of the cosmological defect formation. However, if there is oscillation of the brane distance that may be induced by the brane inflation or by a later phase transition that changes the brane distance, the four-dimensional counterpart of the brane distance (i.e. the Higgs field) oscillates in the effective action. In the four-dimensional counterpart, defect production induced by such oscillation is already discussed by many authors, including the production of the sphaleron domain walls which otherwise cannot be produced in the Universe\\cite{matsuda_deformation}. The defect production induced by such oscillation may or may not be explained by the Kibble mechanism, however it should be fair to distinguish it from the ``conventional'' Kibble mechanism.  Let us summarize the above discussion about the defect production in the brane Universe. Actually, it is possible to produce all kinds of defects in the brane Universe, however it is impossible to produce defects other than the strings {\\bf simply as the result of the brane creation that is induced by the conventional Kibble mechanism}. One should therefore be careful about the assumption that is made in the manuscript, which may or may not be explicit. The necklaces are produced as the hybrid of the brane creation and the brane deformation. It should be noted that the stable loops of the necklaces that we will discuss in this talk may appear in the four-dimensional gauge dynamics, irrespective of the existence of the branes\\cite{050906x}. The stabilization of the necklace loops is first discussed in Ref.\\cite{matsuda_necklace} for brane models and in Ref.\\cite{050906x} for necklaces embedded in four-dimensional gauge dynamics.  In order to produce necklaces in the brane Universe, the motion in the compactified direction is important. I know that in the ``standard scenario'' it is sometimes discussed that the position of the strings are fixed by the potential that is induced by the supersymmetry breaking, and the position is a homogeneous parameter of the Universe because all the decay products (Typically, they are F, D, and $(p,q)$ strings) lie (at least initially) along the same plane of the original hypersurface on which the tachyon condensation took place. However, in this case one may hit upon the idea that the potential for a string cannot be identical to all the other kinds of the strings. One may therefore obtain many kinds of strings that may move independently along different hypersurfaces, with exponentially small intersection ratios. Moreover, I think it is not reasonable(but may be possible) to assume that the string motion is utterly restricted by the potential even in the most energetic epoch just after inflation. Please remember that in general the moving (inflating) brane carries kinetic energy, and the brane annihilation should be an energetic process, although one may admit that there could be exceptional scenarios. I therefore think that the decay products should have kinetic energy, which is enough to climb up the potential hill at least just after brane inflation. \\begin{figure}[h] \\begin{picture}(640,330)(0,0) \\resizebox{16cm}{!}{\\includegraphics{JGRGFig.eps}} \\end{picture} \\caption{We show how necklaces are produced despite the ``standard'' arguments. } \\label{fig:fig0} \\end{figure} ",
        "conclusions": ""
    },
    "0601/gr-qc0601085_arXiv.txt": {
        "abstract": "Quantum gravity is expected to be necessary in order to understand situations where classical general relativity breaks down. In particular in cosmology one has to deal with initial singularities, i.e.\\ the fact that the backward evolution of a classical space-time inevitably comes to an end after a finite amount of proper time. This presents a breakdown of the classical picture and requires an extended theory for a meaningful description. Since small length scales and high curvatures are involved, quantum effects must play a role. Not only the singularity itself but also the surrounding space-time is then modified. One particular realization is loop quantum cosmology, an application of loop quantum gravity to homogeneous systems, which removes classical singularities. Its implications can be studied at different levels. Main effects are introduced into effective classical equations which allow to avoid interpretational problems of quantum theory. They give rise to new kinds of early universe phenomenology with applications to inflation and cyclic models. To resolve classical singularities and to understand the structure of geometry around them, the quantum description is necessary. Classical evolution is then replaced by a difference equation for a wave function which allows to extend space-time beyond classical singularities. One main question is how these homogeneous scenarios are related to full loop quantum gravity, which can be dealt with at the level of distributional symmetric states. Finally, the new structure of space-time arising in loop quantum gravity and its application to cosmology sheds new light on more general issues such as time. ",
        "introduction": "\\label{section:introduction}  \\begin{flushright} \\begin{quote} {\\em Die Grenzen meiner Sprache bedeuten die Grenzen meiner Welt.}  ({\\em The limits of my language mean the limits of my world.}) \\end{quote}  {\\sc Ludwig Wittgenstein}  Tractatus logico-philosophicus \\end{flushright}  While general relativity is very successful in describing the gravitational interaction and the structure of space and time on large scales \\cite{Will}, quantum gravity is needed for the small-scale behavior. This is usually relevant when curvature, or in physical terms energy densities and tidal forces, becomes large. In cosmology this is the case close to the big bang, and also in the interior of black holes. We are thus able to learn about gravity on small scales by looking at the early history of the universe.  Starting with general relativity on large scales and evolving backward in time, the universe becomes smaller and smaller and quantum effects eventually become important. That the classical theory by itself cannot be sufficient to describe the history in a well-defined way is illustrated by singularity theorems \\cite{SingTheo} which also apply in this case: After a finite time of backward evolution the classical universe will collapse into a single point and energy densities diverge. At this point, the theory breaks down and cannot be used to determine what is happening there. Quantum gravity, with its different dynamics on small scales, is expected to solve this problem.  The quantum description does not only present a modified dynamical behavior on small scales but also a new conceptual setting. Rather than dealing with a classical space-time manifold, we now have evolution equations for the wave function of a universe.  This opens a vast number of problems on various levels from mathematical physics to cosmological observations, and even philosophy.  This review is intended to give an overview and summary of the current status of those problems, in particular in the new framework of loop quantum cosmology. ",
        "conclusions": ""
    },
    "0601/astro-ph0601521_arXiv.txt": {
        "abstract": "We report the parallax and proper motion of millisecond pulsar J0030+0451, one of thirteen known isolated millisecond pulsars in the disk of the Galaxy. We obtained more than 6 years of monthly data from the 305 m Arecibo telescope at 430 MHz and 1410 MHz. We measure the parallax of PSR J0030+0451 to be $3.3\\pm0.9$\\,mas, corresponding to a distance of $300\\pm90$\\,pc. The Cordes and Lazio (2002) model of galactic electron distribution yields a dispersion measure derived distance of 317 pc which agrees with our measurement. We place the pulsar's transverse space velocity in the range of 8 to 17 km~s$^{-1}$, making this pulsar one of the slowest known. We perform a brief census of velocities of isolated versus binary millisecond pulsars.  We find the velocities of the two populations are indistinguishable. However, the scale height of the binary population is twice that of the isolated population and the luminosity functions of the two populations are different.  We suggest that the scale height difference may be an artifact of the luminosity difference. ",
        "introduction": "A pulsar parallax can be combined with a measurement of the pulsar's dispersive delay (the Dispersion Measure or DM) to provide an accurate measure of the free electron density along the line of sight (LOS).  PSR J0030+0451 is one of fewer than a dozen pulsars to have its parallax measured via timing \\citep{Kaspi94, Camilo94_1713, Sandhu97, Toscano99, Wolszczan00, Jacoby03, Hotan04, Loehmer04, Splaver05}.  Another dozen have been measured via VLBI (e.g. see \\nocite{Brisken02, Chatterjee04} Brisken et al. 2002, Chatterjee et al. 2004, and also this URL\\footnote{ http://www.astro.cornell.edu/$\\sim$shami/psrvlb/parallax.html}). These measurements are important because they give us most of our knowledge about the galactic thermal electron distribution \\citep{Cordes02, Toscano99, Taylor93}.  In addition, PSR J0030+0451 presents a rare evolutionary case as an isolated millisecond pulsar (MSP). In the most popular MSP evolutionary model, MSPs are formed via accretion of matter from a companion star.  The incoming matter adds to the pulsar's angular momentum, i.e., the pulsar is ``spun up.''  However, isolated MSPs present a conundrum:  they were presumably spun up, yet they are without a companion which would have done so. One possible scenario is that the pulsar has ablated its companion \\citep{Ruderman89}.  We expect MSPs to have lower velocities than the regular population, because the kick from the supernova progenitor had to be small enough to leave the binary intact.  To produce an isolated MSP, the binary must remain intact during and after the supernova, but then after the spin-up phase the companion must leave the system or be evaporated. Several authors have debated whether isolated MSP velocities are lower, higher, or indistinguishable from those of the general population of MSPs. \\cite{McLaughlin04b} suggest we might expect isolated MSPs to have higher velocities. They argue that if isolated MSPs are formed by ablation, we would expect them to form from the tighter binaries which are more susceptible to ablation. The correlation between tight binaries and higher velocities is suggested by \\citet{Tauris96}. \\citet{McLaughlin04b} present the argument for faster velocities for isolated MSPs as a counterpoint to their timing proper motion and scintillation measurements which suggest the opposite, as do the measurements of \\citet{Johnston98} and \\cite{Toscano99}. \\cite{Hobbs05}, however, find the velocities of the populations to be indistinguishable. We present a measurement of the transverse velocity of PSR J0030+0451 which is unusually small, even compared to the isolated MSP population.   We reconsider the question of the velocity of isolated MSPs as compared to the binary MSP population.  Owing to its small timing residual, $\\sim$1~$\\mu$s, PSR J0030+0451 is a good candidate for membership in the Pulsar Timing Array (PTA), which is a collection of pulsars that will be used for detecting gravitational radiation \\citep{Jaffe03, Jenet04}. For this reason, continued refinement of the timing model of PSR J0030+0451 is important. In fact, the PTA, as it is conceived, is an interferometer, so a variety of baselines will be important to its operation. PSR J0030+0451 may be particularly useful in this regard because it has large angular separation from PSRs B1855+09, J1713+07 and J0437-4715, which are among the most stable and precise pulsars \\citep{Kaspi94, vanStraten01, Lommenthesis}.  In \\S2 we present a significant refinement to the timing model previously published \\citep{Lommen00}. We discuss our measurements of parallax and proper motion in \\S\\ref{sec:pm_px}. We include the effects of the solar wind in our analysis, which we discuss in \\S4. In \\S5,\\S6, and \\S7 we discuss the space velocity of J0030+0451, corrections to its measured period derivative, and the implications of its measured distance for the local interstellar medium (LISM). In \\S8 we summarize our conclusions. ",
        "conclusions": "We have measured the parallax of PSR J0030+0451 to be $3.3\\pm0.5$\\,mas. We have measured its proper motion to be $-5.74\\pm 0.09$\\,~mas~yr$^{-1}$ in the plane of the ecliptic and have established an upper limit on its motion out of the plane at 10 mas~yr$^{-1}$. The is one of the lowest velocities measured for any pulsar and is noteworthy even within the relatively low-velocity millisecond pulsar population.   Combining proper motion data from this pulsar with the collection of existing MSP proper motion measurements, we find the statistical properties of the transverse velocities of isolated and binary MSPs are indistinguishable from each other. We do, however, find that the average $z$-height of the isolated MSPs is half that of the binary MSPs. We suggest that a luminosity difference between the two classes of objects, such as that suggested by \\citet{Bailes97}, \\citet{Kramer98}, and \\citet{Hobbs04}, is the simplest way to account for both the observed difference in $z$-height and the similarity of the velocity distributions."
    },
    "0601/astro-ph0601651_arXiv.txt": {
        "abstract": "We present an analysis of 13 of the best quality Ultraluminous X-ray source (ULX) datasets available from \\xmmn European Photon Imaging Camera (EPIC) observations.  We utilise the high signal-to-noise in these ULX spectra to investigate the best descriptions of their spectral shape in the 0.3--10 keV range.  Simple models of an absorbed power-law or multicolour disc blackbody prove inadequate at describing the spectra.  Better fits are found using a combination of these two components, with both variants of this model - a cool ($\\sim 0.2$ keV) disc blackbody plus hard power-law continuum, and a soft power-law continuum, dominant at low energies, plus a warm ($\\sim 1.7$ keV) disc blackbody - providing good fits to 8/13 ULX spectra.  However, by examining the data above 2 keV, we find evidence for curvature in the majority of datasets (8/13 with at least marginal detections), inconsistent with the dominance of a power-law in this regime.  In fact, the most successful empirical description of the spectra proved to be a combination of a cool ($\\sim 0.2$ keV) classic blackbody spectrum, plus a warm disc blackbody, that fits acceptably to 10/13 ULXs.  The best overall fits are provided by a physically self-consistent accretion disc plus Comptonised corona model ({\\sc diskpn + eqpair}), which fits acceptably to 11/13 ULXs.  This model provides a physical explanation for the spectral curvature, namely that it originates in an optically-thick corona, though the accretion disc photons seeding this corona still originate in an apparently cool disc.  We note similarities between this fit and models of Galactic black hole binaries at high accretion rates, most notably the model of \\citet{donek05}.  In this scenario the inner-disc and corona become energetically-coupled at high accretion rates, resulting in a cooled accretion disc and optically-thick corona.  We conclude that this analysis of the best spectral data for ULXs shows it to be plausible that the majority of the population are high accretion rate stellar-mass (perhaps up to 80-$M_{\\odot}$) black holes, though we cannot categorically rule out the presence of larger, $\\sim 1000$-\\Msun intermediate-mass black holes (IMBHs) in individual sources with the current X-ray data. ",
        "introduction": "\\einstein X-ray observations were the first to reveal point-like, extranuclear sources in some nearby galaxies with luminosities in excess of $10^{39} \\ergsec$ \\citep{fabbiano89}. Subsequently, many of these so-called Ultraluminous X-ray sources (ULXs) have displayed short and long term variability, which suggests they are predominantly accreting objects (see \\citealt{milcol04} and references therein). However, the observed luminosities of most ULXs exceed the Eddington limit for spherical accretion onto a stellar-mass ($\\sim$10-$ M_{\\odot}$) black hole (BH). In fact, their luminosites are intermediate between those of normal stellar mass BH X-ray Binaries (BHBs) and Active Galactic Nuclei (AGN).  Therefore, the accretion of matter onto {\\it intermediate-mass\\/} black holes (IMBHs, of $\\sim$$10^2$--$10^4$ $ M_{\\odot}$) provide an attractive explanation for ULXs, and could represent the long sought-after `missing link' between stellar mass BHs and the supermassive BHs in the nuclei of galaxies. However, the large populations of ULXs associated with sites of active star formation (\\eg in the Cartwheel galaxy, \\citealt{gao03}) demand rather too high formation rates of IMBHs if they are to explain the ULX class as a whole \\citep{king04}. An alternative to accreting IMBHs is that ULXs may be a type of stellar-mass BHB with geometrically \\citep{king01} or relativistically \\citep*{koerding02} beamed emission, such that their intrinsic X-ray luminosity does not exceed the Eddington limit.  Another possibility is that they are stellar-mass BHBs that can achieve truly super Eddington luminosities via slim \\citep{ebisawa03} or radiation pressure dominated \\citep{begelman02} accretion discs.  As ULXs are probably the brightest class of X-ray binary fueled by the accretion of matter onto a BH\\footnote{Although the most likely reservoir of fuel for an ULX is a companion star, others have been suggested, for example the direct accretion of matter from molecular clouds \\citep{krolik04}.}, a knowledge of the properties of Galactic BHBs could be vital in interpreting their characteristics. Traditionally the X-ray spectra of BHBs have been fitted empirically with two components, namely a power-law continuum and a multicolour disc blackbody (MCD) component (\\citealt{mitsuda84}; \\citealt{makishima86}). In the standard picture, the power-law component is thought to represent inverse-Compton scattering of thermal photons from the accretion disc by hot electrons in a surrounding corona. As such, the power-law represents the hard tail of the X-ray emission while the MCD component models the soft X-ray emission from the accretion disc.  The MCD model itself has been formulated based on the best known model for accretion onto BHs (\\ie the thin accretion disc model, \\citealt{shakura73}).  It has long been recognised that Galactic BHBs demonstrate various X-ray spectral states which are defined by the balance of these two components (\\ie power-law and MCD) at any one time.  The three most familiar X-ray bright states are the low/hard (LH), high/soft (HS, also described as `thermal dominated') and the very high (VH, or `steep power-law') states (see \\citealt{mcclintock03} for further details).  At lower mass accretion rates, a BHB usually enters the LH state where their X-ray emission is dominated by a hard power-law component ($\\Gamma$$\\sim$1.7), thought to arise from Comptonisation of soft photons by a hot optically thin corona\\footnote{However this is still a topic of debate, with the main alternative for the X-ray power-law emission being synchrotron emission from the radio jet that is associated with this state (\\eg \\citealt{falcke95}; \\citealt*{markoff01}).}.  In this state, the disc is either undetected (\\eg \\citealt{belloni99}) or appears truncated at a much larger inner radius and hence cooler than the parameters derived for the soft state (\\citealt{wilms99}, \\citealt{mcclintock01}).  The soft X-ray state is generally seen at a higher luminosity (\\ie the HS state) and is best explained as $\\sim$1 keV thermal emission from a multi-temperature accretion disc (\\ie modelled with a MCD component). In this state, the spectrum may also display a hard tail that contributes a small percentage of the total flux.  The VH state is in many cases the most luminous state and is characterised by an unbroken power-law spectrum extending out to a few hundred keV or more.  The photon index is typically steeper ($\\geq 2.5$) than found in the LH state and generally coincides with the onset of strong X-ray quasi-periodic oscillations (QPOs).  A MCD component may also be present in the VH state and the \\exosat era demonstrated that some of the QPOs occur when both disc and power-law components contribute substantial luminosity \\citep{vanderklis95}.  The idea of ULXs as analogues to Galactic BHBs in the HS state was supported by \\asca observations, which revealed that their 0.5--10 keV spectra were successfully fitted with the MCD model with relatively high disc temperatures (1.0--1.8 keV, \\citealt{makishima00}).  As such, the ULXs were considered to be mass-accreting BHs with the X-ray emission originating in an optically-thick accretion disc.  In fact, the use of the MCD model to describe these spectra permits one to obtain an `X-ray--estimated' BH mass, $M_{\\rm XR}$, from the following equation (cf. \\citealt{makishima00} equations (5)--(8)).  \\begin{equation} M_{\\rm XR} = {{\\xi\\kappa^2}\\over{8.86\\alpha}} \\frac{D}{\\sqrt{{\\rm cos} i}} \\sqrt{\\frac{f_{\\rm bol}}{2\\sigma T^4_{\\rm in}}} \\; M_{\\odot} \\end{equation}  Where $D$ is the distance to the X-ray source, which has an inclination $i$, a full bolometric luminosity (from the MCD model) of $f_{\\rm bol}$ and an observed maximum disc colour temperature $T_{\\rm in}$.  In addition $\\sigma$ is the Stefan-Boltzmann constant, $\\kappa$ is the ratio of the colour temperature to the effective temperature (`spectral hardening factor'), and $\\xi$ is a correction factor reflecting the fact that $T_{\\rm in}$ occurs at a radius somewhat larger than $R_{\\rm in}$ (here, we assume that $R_{\\rm in}$ is at the last stable Keplerian orbit).  \\cite{makishima00} use values of $\\xi = 0.412$ and $\\kappa = 1.7$, though other work has found different values for the spectral hardening factor (e.g. $\\kappa = 2.6$ for GRO J1655-40, \\citealt{st03}).  Finally, $\\alpha$ is a positive parameter with $\\alpha = 1$ corresponding to a Schwarzschild BH.  However, the masses inferred from the \\asca data and Equation (1) are far too low to be compatible with the large BH masses suggested by their luminosities (assuming Eddington-limited accretion), for standard accretion discs around Schwarzchild BHs.  \\citet{makishima00} suggested that this incompatibility could be explained if the BHs were in the Kerr metric (\\ie rapidly rotating objects), allowing smaller inner disc radii and hence higher disc temperatures.  \\chandra observations have provided some support for the \\citet{makishima00} results, with some ULX spectra being consistent with the MCD model (\\eg \\citealt{roberts02}).  However, \\chandra also revealed that some ULX spectra showed a preference for a power-law continuum rather than the MCD model (\\eg \\citealt{strickland01}; \\citealt{roberts04}; \\citealt{terashima04}). It has been suggested that this preference for a power-law spectrum could be interpreted in terms of the LH state seen in Galactic BHB candidates, relativistically beamed jets or emission from a Comptonised accretion disc in the VH state.  As well as these single component models, \\asca and \\chandra spectroscopy have also suggested the presence of two component spectra for some ULXs, comprising a MCD with a power-law component.  For example, previous \\asca analyses hinted at evidence for IMBHs, \\ie cool accretion disc components (see below), but these observations were not sensitive enough to statistically require two component modelling (\\eg \\citealt{colbert99}).  Similar results have been obtained by {\\it Chandra}, \\eg ULXs in NGC 5408 \\citep{kaaret03} and NGC 6946 \\citep{roberts03}.  Conversely, \\chandra spectra of the Antennae ULXs (\\citealt{zezas02a}; \\citealt{zezas02b}) revealed an accretion disc (MCD) component consistent with the high temperature \\asca results (\\ie $kT_{\\rm in}$$\\sim$1 keV), together with a hard power-law component ($\\Gamma$$\\sim$1.2).  It is only recently, using high quality {\\it XMM-Newton}/EPIC spectroscopy of ULXs, that it has been demonstrated that the addition of a soft thermal disc component to a power-law continuum spectrum provides a strong statistical improvement to the best fitting models to ULX data (\\eg \\citealt{miller03}; \\citealt*{miller04a}). These particular observations have provided strong support for the IMBH hypothesis by revealing disc temperatures in these sources up to 10 times lower than commonly measured in stellar mass BHBs, consistent with the expectation for the accretion disc around a $\\sim$1000-\\Msun IMBH\\footnote{It is common for the generic range of masses for IMBHs to be quoted as 20--$10^6$ \\Msun.  The lower limit comes from a consideration of the measured masses of BHs in our own Galaxy \\citep{mcclintock03}, and a theoretical limit for the mass of a BH formed from a single massive star \\citep{fryer01}.  However, more recent population synthesis analyses show that BHs of up to $\\sim 80$-\\Msun may be formed in young stellar populations \\citep{belczynski04}.  Hence, when we refer to IMBHs in this paper we refer specifically to the larger $\\sim$1000-\\Msun IMBHs implied by the cool accretion disc measurements.} (cf. Equation (1)).  However, in a few cases, {\\it XMM-Newton}/EPIC observations have revealed a more unusual two-component X-ray spectrum.  A detailed analysis of such a source is presented in \\citet*{stobbart04}. In this case, the ULX is the brightest X-ray source in the nearby (1.78 Mpc) Magellanic-type galaxy NGC 55 and, although its X-ray luminosity only marginally exceeds $10^{39} \\ergsec$, it represents one of the highest quality ULX datasets obtained to date. The initial low flux state data were best fitted with an absorbed power-law continuum ($\\Gamma$$\\sim$4), while a subsequent flux increase was almost entirely due to an additional contribution at energies $> 1$ keV, adequately modelled by a MCD component ($kT_{\\rm in}$$\\sim$0.9 keV).  Whilst this accretion disc component is reasonable for stellar-mass BHs, the dominance of the power-law continuum at soft X-ray energies is problematic.  Such a soft power-law cannot represent Comptonised emission from a hot corona, as one would not expect to see the coronal component extend down below the peak emissivity of the accretion disc, where there would be insufficient photons to seed the corona. Alternative sources of seed photons for the corona are unlikely; for example, the incident photon flux of the secondary star at the inner regions of the accretion disc is too low to provide the seeding (cf. \\citealt{roberts05}).  It also seems unlikely that the power-law emission could arise from processes at the base of a jet, as these are typically represented by much harder photon indices than measured here ($\\Gamma$$\\sim$1.5--2; \\citealt{mnw05} and references therein). Indeed, with the possible exception of NGC 5408 X-1 \\citep{kaaret03}, there is no evidence that ULXs do display bright radio jets, though this cannot be excluded by current observations \\citep{kcf05}.  The possibility of the soft component resulting from an outflow of material from the accretion disc may also be discounted as this would produce a thermal spectrum rather than a power-law continuum.  Although this spectral description has not been seen in Galactic systems, a second case has been reported independently for the nearest persistent extragalactic ULX (M33 X-8) by \\citet{foschini04}.  The non-standard model provided the best fit to this ULX with $\\Gamma$$\\sim$2.5 and $kT_{\\rm in}$$\\sim$1.2 keV.  However, this source is also at the low luminosity end of the ULX regime with \\lx$\\sim$2$\\times 10^{39}\\ergsec$. A further possible case, in a more luminous ULX, has arisen from the \\xmmn data analysis of NGC 5204 X-1 \\citep{roberts05}.  In this case the authors show that there is spectral ambiguity between the non-standard fit ($\\Gamma$$\\sim$3.3, $kT_{\\rm in}$$\\sim$ 2.2 keV) and the IMBH model ($\\Gamma$$\\sim$2.0, $kT_{\\rm in}$$\\sim$ 0.2 keV), with both providing statistically acceptable fits to the data.  Even more recently, two additional examples of this spectral form have been uncovered in an \\xmmn survey of ULXs by \\citet{feng05}.  Although it is difficult to derive a literal physical interpretation from the non-standard model, it does provide an accurate empirical description of ULX spectra in some cases, and as such it has the potential to provide new insights into the nature of these sources. Therefore in this paper we re-evaluate current data in an attempt to determine the best spectral description for the shape of high quality ULX spectra, and ask what consequences this has for the idea of ULXs as accreting IMBHs. The paper is structured as follows: Sec.~\\ref{sec_sample} -- introduction to the ULX sample; Sec.~\\ref{sec_obs} -- details of the observations and data reduction; Sec.~\\ref{sec_spectra} -- description of the spectral analysis; Sec.~\\ref{sec_lx_kt} -- a comment on the luminosity and inner disc temperature relationship of these ULXs; Sec.~\\ref{sec_discussion} -- a discussion of our results; and finally Sec.~\\ref{sec_conclusions} -- our conclusions. ",
        "conclusions": "\\label{sec_conclusions}  We have conducted a detailed examination of the X-ray spectral shapes in a sample of the highest quality \\xmmn EPIC ULX datasets available to us.  Most notably, more than half of the ULXs show at least marginal evidence for curvature in their 2--10 keV spectra, which is somewhat unexpected if they are to be interpreted as the accreting $\\sim 1000$-\\Msun IMBHs suggested by modelling the soft spectral components as accretion discs.  Physical modelling shows that this curvature is likely to originate in optically-thick coronae, which in turn leads to interpretations of the ULXs in terms of high accretion-rate stellar-mass (or slightly larger) BHs operating at around the Eddington limit.  However, while we conclude that it is likely that the general ULX population does not have a large contribution from IMBHs, we obviously cannot rule out the possibility that some ULXs do possess IMBHs.  Perhaps the best candidate on the basis of our spectral fitting is M81 X-9, which is well fitted by cool disc plus power-law/optically-thin corona models, and does not show explicit curvature in its 2--10 keV spectrum, though even this ULX may be fitted using a hot ($\\sim 2.2$ keV) accretion disc plus soft excess model.  Clearly, it is difficult to find unique solutions for these sources even with high quality \\xmmn EPIC data.  Ultimately, we may perhaps have to wait for radial velocity measurements from the optical counterpart of an ULX, leading to dynamical mass measurements of the compact accretor, before we have conclusive evidence whether any individual ULX does harbour an IMBH."
    },
    "0601/astro-ph0601698_arXiv.txt": {
        "abstract": "Binary supermassive black holes form naturally in galaxy mergers, but their long-term evolution is uncertain. In spherical galaxies, $N$-body simulations show that binary evolution stalls at separations much too large for significant emission of gravitational waves (the ``final parsec problem''). Here, we follow the long-term evolution of a massive binary in more realistic, triaxial and rotating galaxy models. We find that the binary does not stall. The binary hardening rates that we observe are sufficient to allow complete coalescence of binary SBHs in $10$ Gyr or less, even in the absence of collisional loss-cone refilling or gas-dynamical torques, thus providing a potential solution to the final parsec problem. ",
        "introduction": "When two galaxies containing supermassive black holes (SBHs) merge, a binary SBH forms at the center of the new galaxy. The two SBHs can eventually coalesce, but only after stellar- or gas-dynamical processes bring them close enough together ($\\lap 10^{-2}$ pc) that gravitational radiation is emitted. There is strong circumstantial evidence that rapid coalescence is the norm. For instance, no binary SBH has ever been unambiguously observed \\citep{LR05}. Furthermore, in a galaxy containing an uncoalesced binary, mergers would eventually bring a third SBH into the nucleus, precipitating a gravitational slingshot interaction that would eject one or more of the SBHs from the nucleus \\citep{MV90,VHM03}. This could produce off-center SBHs, and could also weaken the tight correlations that are observed between SBH mass and galaxy properties \\citep{FM00,G01,MH03}.  Unless the binary mass ratio is extreme, dynamical friction rapidly brings the smaller SBH into a distance $\\sim G\\mu/\\sigma^2$ from the larger SBH, where $\\mu\\equiv M_1M_2/(M_1+M_2)$ is the binary reduced mass and $\\sigma$ is the 1D velocity dispersion of the stars. At this separation -- of order $1$ pc -- the two SBHs begin to act like a ``hard'' binary, ejecting passing stars with velocities large enough to remove them from the nucleus. $N$-body simulations \\citep{MF04,SMM05,BMS05} show that continued hardening of the binary takes place at a rate that depends strongly on the number $N$ of ``star'' particles used in the simulation. As $N$ increases, the hardening rate falls, as expected if the binary's loss cone is repopulated by star-star gravitational encounters \\citep{Y02,MM03}. When extrapolated to the much larger $N$ of real galaxies, these results suggest that binary evolution would generally stall (the ``final parsec problem'').  To date, $N$-body simulations of the long-term evolution of binary SBHs have only been carried out using spherical or nearly spherical galaxy models. But it has been suggested \\citep{MP04} that binary hardening might be much more efficient in non-axisymmetric galaxies due to the qualitatively different character of the stellar orbits. Here, we test that suggestion by carrying out the first $N$-body simulations of massive binaries in strongly non-axisymmetric galaxy models. We find that the hardening rate is {\\it independent} of $N$ for particle numbers up to at least $0.4\\times 10^6$. To the extent that our galaxy models are similar to real merger remnants, these results imply that binary SBHs can efficiently harden through purely stellar-dynamical interactions in many galaxies, thus providing a plausible solution to the final parsec problem. ",
        "conclusions": ""
    },
    "0601/astro-ph0601301_arXiv.txt": {
        "abstract": "We present a \\chandra\\ study of mass profiles in 7 elliptical galaxies, of which 3 have galaxy-scale and 4 group-scale halos, demarcated at $10^{13}$\\msun. These represent the best available data for nearby objects with comparable X-ray luminosities. We measure $\\sim$flat mass-to-light (M/L) profiles within an optical half-light radius (\\reff), rising by an order of magnitude at $\\sim$10\\reff, which confirms the presence of dark matter (DM). The data indicate hydrostatic equilibrium, which is also supported by agreement with studies of stellar kinematics in elliptical galaxies. The data are well-fitted by a model comprising an NFW DM profile and a baryonic component following the optical light. The distribution of DM halo concentration parameters (c) {\\em versus} \\mvir\\ agrees with \\lcdm\\ predictions and our observations of bright groups. Concentrations are slightly higher than expected, which is most likely a selection effect. Omitting the stellar mass drastically increases c, possibly explaining large concentrations found by some past observers. The  stellar M/\\lk\\ agree  with population synthesis models, assuming a Kroupa IMF. Allowing adiabatic compression (AC) of the DM halo by baryons made M/L more discrepant, casting some doubt on AC. Our best-fitting models imply total baryon fractions $\\sim$0.04--0.09, consistent with models of galaxy formation incorporating strong feedback. The groups exhibit positive temperature gradients, consistent with the ``Universal'' profiles found in other groups and clusters, whereas the galaxies have negative gradients, suggesting a change in the evolutionary history of the systems around \\mvir$\\simeq 10^{13}$\\msun. ",
        "introduction": "The nature and distribution of dark matter (DM) in the Universe is one of the fundamental problems facing modern physics. Cold DM lies at the heart of  our  current (\\lcdm) cosmological   paradigm, which predicts substantial  DM halos for objects  at all mass-scales from galaxies to clusters.   Although  \\lcdm\\  has  been    remarkably  successful   at explaining large-scale features \\citep[\\eg][]{spergel03a,perlmutter99a}, observations of galaxies have been more problematical for the theory. Dissipationless dark matter simulations find that dark matter halos are well characterized by a ``Universal'' mass density profile \\citep[][hereafter NFW]{navarro97} over a wide range of Virial masses (\\mvir) \\citep[e.g.][]{bullock01a}. Low mass halos tend to form first in hierarchical cosmologies and are consequently more tightly concentrated than their later forming,  high mass counterparts. This tendency  produces a predicted correlation between the DM halo concentration parameter (c, which is ratio between Virial radius, \\rvir, and the characteristic scale of the density profile) and \\mvir \\citep{navarro97}.  However, since mass and formation epoch are not perfectly correlated, we expect a significant scatter at fixed Virial mass \\citep{jing00a,bullock01a,wechsler02a}. The tight link between halo formation epoch and concentration implies that the precise relation between c and \\mvir\\ is sensitive to the underlying Cosmological parameters, including \\sigmaeight\\ and the dark energy equation of state \\citep{kuhlen05a}, making an observational test of this relation potentially a very powerful tool for cosmology.   The mass profiles of galaxies also may provide valuable clues as to the way in which galaxies form  in DM halos. In particular, as baryons cool and collapse into stars, the associated increase in the central mass density should in turn modify the shape of the DM halo. This process is typically modelled assuming adiabatic contraction (AC) of the DM particle orbits \\citep[\\eg][]{blumenthal86a,gnedin04a}. If the galaxy halo subsequently evolves by major mergers, simulations are unclear as to whether these features would persist \\citep[\\eg][]{gnedin04a} or whether the merging process may destroy this imprint of star formation, or even mix the DM and baryons sufficiently to produce a {\\em total} gravitating mass profile more akin to NFW \\citep{loeb03a,elzant04a}.   Observational tests of the predictions of \\lcdm\\ have proven controversial. In clusters of galaxies there is overwhelming evidence for DM, and an increasing body of work verifying the predictions of \\lcdm. In particular recent, high-quality \\chandra\\ and \\xmm\\ observations have revealed mass profiles in remarkable agreement with the Universal profile from deep in the core to a large fraction of \\rvir\\ \\citep[\\eg][]{lewis03a,zappacosta06a,vikhlinin05b}, and a distribution of c {\\em versus} \\mvir\\ in good agreement with \\lcdm\\ \\citep{pointecouteau05a}. In galaxies, however, the picture is much less clear. Rotation curve analysis of low surface brightness (LSB) disk galaxies has suggested significantly less cuspy density profiles than expected \\citep[\\eg][]{swaters00a}. Although this discrepancy led to a serious discussion of modifications to the standard paradigm \\citep[\\eg][]{hogan00a,spergel00a,zentner02a,kaplinghat05a,cembranos05a}, recent  results, taking account of observational bias and the 3-dimensional geometry of the DM halos, have done much to resolve the discrepancy \\citep[\\eg][]{swaters03a,simon05a}. However, some significant discrepancies remain, not least of which is that the DM halos of these galaxies appear less concentrated than expected \\citep[\\eg][]{gonzalez00a,kassin06a}. A possible explanation is that LSB galaxies are preferentially found in low-concentration halos \\citep{bullock01a,bailin05a,wechsler05a}, making additional constraints at the galaxy scale extremely important. %  In many respects, kinematical mass measurements are far more challenging for early-type than spiral galaxies. As essentially pressure-supported systems little is known {\\em a priori} about the velocity anisotropy tensor of the stars in elliptical galaxies, which is problematical for the determination of the mass from stellar motions. Nonetheless, stellar kinematical measurements have widely been used as a means to measure the gravitating matter within $\\sim$the optical half-light radius (\\reff) of elliptical galaxies \\citep[\\eg][]{binney90a,vandermarel91a,gerhard01a}. These studies tend to find relatively flat mass-to-light (M/L) ratios within \\reff, implying that most of the matter within this radius is baryonic. Consideration of the tilt in the fundamental plane can also lead to the same conclusion \\citep{borriello03a}. In contrast, \\citet{padmanabhan04a} pointed out that dynamical M/L ratios within \\reff\\ are much larger than predicted by realistic stellar population synthesis models for stars alone, allowing \\gtsim 50\\% of the mass within \\reff\\ to be dark. %  Attempts to extend kinematical studies of elliptical galaxies to larger radii, where DM should be dominant, have proven controversial. In particular \\citet{romanowsky03a} argued against the existence of DM in a small sample of elliptical galaxies, based on planetary nebulae dynamics within $\\sim$5\\reff. We note that this sample was heavily biased towards very X-ray faint objects, which might hint at low-mass halos since they have not held onto their hot gas. In any case \\citet{dekel05a} pointed out that their conclusions were very sensitive to the uncertainty in the velocity anisotropy tensor, for plausible values of which the data were consistent with substantial DM halos. In fact globular cluster dynamics in one of these systems, NGC\\thin 3379, does imply a significant amount of DM \\citep{pierce06a,bergond06a}. As more kinematical studies of early-type galaxies at large radii are appearing, it is becoming clear that at least some elliptical galaxies host considerable DM halos \\citep[\\eg][]{statler99a,romanowsky05a}. There persist some questions, however, as to the extent to which all galaxies have DM halos consistent with \\lcdm. In particular \\citet{napolitano05a} argued that a substantial number of early-type galaxy halos appear  less concentrated than expected.  Gravitational lensing provides further evidence that, at least some, early-type galaxies possess substantial DM halos \\citep[\\eg][]{kochanek95a,fischer00a,rusin02a}. Since weak lensing of galaxies only provides useful mass constraints in a statistical sense, the relatively rare instances of strong lensing are required to study DM in individual systems. Nonetheless it has been possible in a few cases to decompose the mass into stellar and DM components, albeit with strong assumptions or additional observational constraints \\citep[\\eg][]{rusin03a,treu04a}.    X-ray observations of the hot gas in early-type galaxies provide a complementary means to infer the mass-profiles {\\em via} techniques similar to those used in studying clusters. Since the X-ray emission from early-type galaxies is typically not very bright, prior to the advent of  \\chandra\\ and \\xmm\\ this was limited by the relatively sparse information on the radial temperature and density profiles of the hot gas which could be determined by prior generations of satellites. Notwithstanding this limitation, large M/L ratios (consistent with substantial DM) were inferred for a number of X-ray bright galaxies, albeit with strong assumptions concerning the temperature and density profiles \\citep[\\eg][]{forman85a,loewenstein99b}. Using a novel technique which relied, instead, on the ellipticity of the X-ray halo, \\citet{buote94} were able robustly to detect DM  in the isolated elliptical NGC\\thin 720 \\citep[see also][]{buote96a,buote98d,buote02b}. Detailed measurements of the radial mass distribution were, however, largely restricted to a few massive systems, which may be entwined with a group halo \\citep[\\eg][]{irwin96,brighenti97a}. Nevertheless \\citet{brighenti97a} were able to decompose the mass profiles of two systems, NGC\\thin 4472 and NGC\\thin 4649, into stellar and DM components. \\citet{sato00a} investigated the \\mvir-c relation using \\asca\\ for a sample of objects ranging from massive clusters to $\\sim$3 elliptical galaxies. The limited spatial resolution of \\asca\\ necessitated some assumptions about the density  profiles and, crucially, the authors neglected any stellar mass component in their fits. This omission may explain the very steep \\mvir-c relation (with c$_{200}$\\gtsim 30 for the galaxies) found by these authors, in conflict with \\lcdm\\ \\citep{mamon05a}.  Although mass profiles of early-type galaxies are beginning to appear which exploit the improved sensitivity and resolution of \\chandra\\ and \\xmm, many of the most interesting constraints on DM are still restricted to massive systems, which may be at the centres of groups. For example, \\citet{fukazawa06a} reported \\chandra\\ and \\xmm\\ M/L profiles for $\\sim$50 galaxies and groups, confirming $\\sim$flat profiles within \\reff\\ which rise at larger radii. However, the constraints at large radii were dominated by the massive (group-scale) objects so the implications for the DM content of normal galaxies are unclear. Furthermore, the authors  included a substantial number of highly disturbed systems, in which hydrostatic equilibrium may be questioned, and failed to account for the unresolved sources which dominate the emission in the lowest-\\lx\\ objects in their sample\\footnote{Although the authors account for unresolved sources when measuring the gas temperature, they do not account for it when computing the gas density, where its effect is more pronounced}. Recently, however, detailed \\chandra\\ and \\xmm\\ mass profiles have begun to appear for isolated early-type galaxies, also confirming the presence of massive DM halos \\citep[\\eg][]{osullivan04b,khosroshahi04a}.   This paper is part of a series \\citep[see also][]{gastaldello06a,zappacosta06a,buote06b,buote06a} using high-quality \\chandra\\ and \\xmm\\ data to investigate the mass profiles of  galaxies, groups and clusters. This provides an unprecedented opportunity to place definitive constraints upon the \\mvir-c relation over $\\sim$2 orders of magnitude in \\mvir. In this paper, we focus on the temperature, density and mass profiles of seven galaxies and poor groups chosen from the \\chandra\\ archive. In order to compare to theory we perform spherically-averaged analysis, leaving a discussion of the ellipticities of the X-ray halos to a future paper. In \\S~\\ref{sect_targets} we discuss the target selection. The data-reduction is described in \\S~\\ref{sect_reduction} and the X-ray morphology is addressed in \\S~\\ref{sect_imaging}. We discuss the spectral analysis in \\S~\\ref{sect_spectra}, the mass analysis in \\S~\\ref{sect_mass}, the systematic uncertainties in our analysis in \\S~\\ref{sect_systematics} and reach our conclusions in \\S~\\ref{sect_discussion}. The three systems for which we find \\mvir$<10^{13}$\\msun\\ are optically isolated and so we refer to them as ``galaxies'', and the other systems in our sample as groups. We discuss this in more detail in  \\S~\\ref{discussion_groups}. In this paper, all error-bars quoted represent 90\\% confidence limits, unless otherwise stated, and we computed Virial quantities assuming a ``critical overdensity'' factor for the DM halos of $\\rho_{\\rm halo}/\\rho_{\\rm crit} = 103$ (where $\\rho_{\\rm halo}$ is the mean density of a sphere of mass \\mvir\\ and radius \\rvir). ",
        "conclusions": "\\label{sect_discussion} \\subsection{Hydrostatic equilibrium} Our fit results provide strong evidence that the gas is in hydrostatic equilibrium in these systems. Despite highly nontrivial temperature and density profiles, we were able to recover smooth mass profiles in remarkably good agreement with expectation for these systems, using two complementary techniques. If the gas is significantly out of hydrostatic equilibrium, this would represent a remarkable ``conspiracy'' between the density and temperature profiles. It is unsurprising that the gas is close to hydrostatic equilibrium in these systems, since we took care to choose objects with relaxed X-ray morphology. Based on N-body/ hydrodynamical analysis, X-ray measurements are expected to give reliable constraints on the DM in systems without obvious substructure \\citep{buote95a}.  Further support for hydrostatic equilibrium is provided by the general agreement between our measured \\mstars/\\lk\\ ratios and those predicted by SSP models, coupled with the agreement between the measured \\mvir-c relation and that expected. Similarly a comparison between our results and masses determined from stellar dynamics provides even more evidence that our measured mass profiles are reliable. Dynamically-determined \\mgrav/\\lb\\ within the B-band  \\reff\\ are typically found to be $\\sim$4--10 \\citep{gerhard01a,trujillo04a}. We found \\mgrav/\\lb\\ within the B-band \\reff\\ (taken from RC3 or \\citealt{faber89}) for our systems ranged from $\\sim$3 to $\\sim$8, in good agreement with this result. For NGC\\thin 4649, outside \\reff\\ there is excellent agreement between our measured \\mgrav/L profile and that obtained from globular cluster kinematics, although at small radii the X-ray data lie $\\sim$30\\% lower (K.\\ Gebhardt et al, in preparation). \\citet{vandermarel91a} constructed stellar kinematical models for 5 galaxies in our sample (NGC\\thin 720, NGC\\thin 1407 and NGC\\thin 4261, NGC\\thin 4472 and NGC\\thin 4649), under the assumption of a constant M/L profile. Strictly speaking a direct comparison cannot be made between their \\mgrav/\\lb\\ measurements and our results since our data indicate this assumption is incorrect. However, if we simply assume that these M/L ratios represent those integrated out to \\reff, the X-ray inferred masses vary from $\\sim$40\\% lower to $\\sim$10\\% higher than those from kinematics. \\citet{kronawitter00a} report \\mgrav/\\lb$\\sim$8$\\pm1.5$ for NGC\\thin 4472 within $\\sim$50\\arcsec, at which radius our X-ray determined value is $\\sim$50\\% lower. The discrepancies between the X-ray and dynamical masses are only modest (the X-ray mass being on average $\\sim$20\\% lower), indicating that the data must be close to hydrostatic equilibrium. Turbulence is expected to contribute only $\\sim$10\\% pressure support in clusters, which are believed to be more turbulent than galaxies, \\citep{rasia06a}. Therefore, on a case-by-case basis, the observed differences are most likely a manifestation of the mass-anisotropy degeneracy \\citep[\\eg][]{dekel05a}.  \\subsection{Mass profiles} We obtained detailed mass profiles for 3 galaxies and 4 group-scale systems, out to $\\sim$10\\reff. The data clearly show M/L profiles which are $\\sim$flat within \\reff\\ and rise considerably outside this range. This confirms the presence of substantial DM in at least some early-type galaxies and indicates that a stellar mass component dominates within $\\sim$\\reff. This is consistent with studies of stellar kinematics and similar to the mass decomposition analysis of \\citet{brighenti97a}.  The data are well-fitted by a model comprising a stellar mass (H90) component and an NFW DM profile. Omitting the stellar mass component led to systematically poorer fits, smaller \\mvir\\ and vastly larger c ($\\gg$20), which are inconsistent with the predictions of \\lcdm. This effect is easy to understand--- if we add a compact stellar mass component to an (extended) NFW profile, we increase the mass in the core which, by definition, makes the halo more concentrated. However, it is not entirely clear whether this effect, pointed out by \\citet{mamon05a}, can completely account for the significantly steeper \\mvir-c relation found by \\citet{sato00a}. Based on our analysis of group-scale halos \\citep{gastaldello06a} we found that the inclusion of the stellar mass component does not have a strong effect on c in most systems with \\mvir \\gtsim 2$\\times 10^{13}$\\msun, provided the data are fitted to a sufficiently large fraction of \\rvir. The data did not allow us to distinguish statistically between the simple NFW+stars model and scenarios in which the DM halo experiences adiabatic compression due to star formation (however, see \\S~\\ref{discussion_mass_to_light}), or the NFW profile was replaced with the alternative N04 profile.  Comparing our inferred \\mvir\\ and c  to the predictions of \\lcdm\\ we find general agreement. There is some evidence, however, that the concentrations are systematically higher than one would expect, although the error-bars are typically large. Such a trend is also seen in our analysis of groups \\citep{gastaldello06a}. Whilst the slope  of the \\mvir-c relation therefore implied by our data is difficult to explain by varying the cosmological parameters within reasonable limits \\citep{buote06b}, we suspect that the discrepancy can be resolved by taking into account the selection function of our galaxies.  Our data were not selected in a statistically complete manner and, by choosing objects with relatively relaxed X-ray morphologies we are probably selecting objects which have not had a recent major merger. This systematically biases us towards early-forming, hence higher concentration halos. In fact, it is striking that all three {\\em de facto} galaxies in our sample are relatively isolated systems (\\S~\\ref{discussion_groups}). Such systems preferentially might be expected to occupy high-c halos \\citep{zentner05a}, which does appear to be the case for 2 out of 3 of the galaxies. We will return to these issues in detail in \\citet{buote06b}.  \\subsection{Galaxies, Groups and Fossil Groups} \\label{discussion_groups} All three of the lowest-mass systems in our sample are very isolated optically. NGC\\thin 6482 matches the isolation criteria adopted to identify so-called ``fossil groups'' \\citep{khosroshahi04a}. NGC\\thin 4125 and NGC\\thin 720 are both listed as ``groups'' in G93, but closer inspection actually reveals they are also very isolated. Excepting the central galaxy, only one of the putative members of the NGC\\thin 720 ``group'' listed in the G93 catalogue  \\citep[which omits the dwarf galaxy population studied by][]{dressler86a}, actually lies within the projected \\rvir\\ (but outside 0.75$\\times$\\rvir) and it is 2.4 magnitudes fainter in B than the central galaxy. \\citeauthor{dressler86a} remarked upon the optical isolation of this galaxy. Of the two putative companion galaxies to NGC\\thin 4125 given in G93 which lie within the projected \\rvir\\ (but outside 0.67$\\times$\\rvir), both are much fainter (by 2.3 and 3.9 magnitudes, respectively) in B than the central galaxy. In contrast, the four remaining systems in our sample are much less optically isolated. \\citet{schindler99a} show the clear over-density of early-type galaxies around NGC\\thin 4649 and NGC\\thin 4472, and almost 60 group members are associated with these systems by G93. \\citet{gould93a} identified at least 10 members of the NGC\\thin 1407 group, from the dynamics of which he inferred a mass broadly consistent with our measured \\mvir\\ (\\S~\\ref{sect_ngc1407}). \\citet{helsdon03a} report 57 galaxies associated with the NGC\\thin 4261 group within $\\sim$1~Mpc projected radius, which is consistent with our measured \\rvir.   Rather than an isolated galaxy \\citet{khosroshahi04a} identify NGC\\thin 6482 as a ``fossil group''. Fossil groups are group-sized X-ray halos centred on essentially a single elliptical galaxy \\citep{ponman94,vikhlinin99a,jones03}. The typical interpretation of these objects is groups in which all of the \\lstar\\ members have merged. Confusingly, using almost the same selection criteria, \\citet{osullivan04b} classify the galaxy NGC\\thin 4555 as an ``isolated elliptical galaxy'' and posit a very different formation scenario. This object appears to be more massive than NGC\\thin 6482; the authors found \\mgrav $\\sim 3\\times 10^{12}$\\msun\\ within 60~kpc which, assuming an NFW profile with c$=$15 would imply \\mvir $\\sim 2\\times 10^{13}$\\msun. Nonetheless, both of these systems have more in common (both optically and in the X-ray) with each other, and the other isolated ellipticals in our sample, than, for example, the  massive (\\mvir\\gtsim$10^{14}$\\msun), hotter (kT$\\sim$2~keV) fossil groups considered by \\citet{vikhlinin99a}. We suspect that the distinction made between ``isolated elliptical'' and ``fossil group'' for these two systems is largely semantic, and consider NGC\\thin 6482 more properly an isolated  galaxy, too.   The clear division in the galaxy content of our sample clearly lends itself to the nomenclature  ``galaxies'' for the three lowest-mass systems, and ``groups'' for the others. Strikingly, this separation between galaxies and groups also appears consistent with a difference in temperature profiles  (\\S~\\ref{discussion_temp}). That this distinction appears commensurate with \\mvir$\\sim 10^{13}$\\msun\\ is suggestive that this mass-scale may be a useful yard-stick with which to compare to other systems. The error-bars on our mass estimates are sufficiently large that the 90\\% confidence regions of several of the objects (notably NGC\\thin 720, NGC\\thin 6482 and NGC\\thin 1407) actually straddle $10^{13}$\\msun. However, it is clear that {\\em on average}, the systems with \\mvir \\ltsim $10^{13}$\\msun\\ are galaxies. We note that the \\mvir\\ adopted here is that {\\em before} any tidal truncation which is almost certainly occurring as NGC\\thin 4472 and NGC\\thin 4649 merge with Virgo (their untruncated \\rvir\\ would stretch much of the distance to M\\thin 87). \\mvir\\ does not exactly correlate with formation epoch, so that lower-mass halos may still be in the process of forming (hence contain multiple galaxies of similar magnitude), and  more massive halos may contain single, dominant ellipticals (fossil groups). Nonetheless, classifying halos primarily on the basis of \\mvir\\ provides a straightforward way to locate them in the formation hierarchy. Traditionally, galaxy-like and group-like systems are distinguished on the basis of local over-densities of galaxies.  However, placing optically-identified groups into a cosmological context requires a firm understanding not only of the formation of DM halos but also how galaxies populate them, which is much less well-understood \\citep[\\eg][]{kravtsov04a}. This problem is compounded by the difficulties faced by optical group-finding algorithms in identifying very poor groups \\citep[\\eg][]{gerke05a}. Not only can a significant fraction of putative groups be chance superpositions of galaxies, particularly along filaments, but adjacent groups can be merged, such as happened for NGC\\thin 4649 and NGC\\thin 4472 in G93. If there are only a few identified members, small-number statistics and the treatment of interlopers can affect their interpretation \\citep[\\eg][]{gould93a}. To complicate matters further, some authors refer to {\\em any} over-density of galaxies as a group, even a Milky Way-sized galaxy and its dwarf satellites \\citep[\\eg][]{tully05a}.    \\subsection{Stellar Mass-to-Light Ratios} \\label{discussion_mass_to_light} Comparing our measured stellar M/L ratios to the predictions of simple stellar population (SSP) models, we found reasonable agreement provided one assumes a \\citet{kroupa01a} IMF. There is modest disagreement, even when the less-cuspy N04 DM model was adopted. Considering the uncertainties in the SSP modelling (discussed below), however, we believe this discrepancy is not significant. If we allowed the DM profile to be modified by adiabatic compression, we obtained substantially smaller \\mstars/\\lk\\ values from our data, (since it increases the cuspiness of the halo) which are more discrepant with the SSP models. This result casts doubt on AC being as significant an effect as currently modelled. However, the data alone did not allow us statistically to distinguish between the NFW+stars and AC NFW+stars models. Nonetheless, this result is joining a growing body of literature which similarly calls into question whether AC operates as predicted \\citep{zappacosta06a,kassin06a,sand04a}.   There are a number of major uncertainties in the computation of the stellar mass-to-light ratios from the SSP models. Specifically, the results are very dependent upon the assumed IMF, which is not confidently known in early-type galaxies. Furthermore there is some evidence that early-type galaxies frequently contain multiple stellar populations of different ages, including a significant young population \\citep[\\eg][]{rembold05a,nolan06a}. Depending on the mass fraction of the young component, this may substantially reduce \\mstars/\\lk\\ in the galaxy, hence possibly reconciling the data and the AC NFW+stars model. A small amount of star formation may also give rise to a population of stars which can dominate the light in the galaxy core, giving rise to significantly lower synthetic \\mstars/\\lk\\ than measured. This may be the case in NGC\\thin 720 (see \\S~\\ref{sect_mass_to_light}). More problematically, there are known to be significant abundance, or possible age, gradients in the stellar populations of early-type galaxies \\citep[\\eg][]{trager00a,kobayashi99a,rembold05a}, which would translate into stellar \\mstars/\\lk\\ gradients. Our simple modelling did not allow us to account for such an effect {\\em per se}. Although we suspect that such gradients will primarily lead to a \\mstars/\\lk\\ value which reflects an average for the galaxy, \\mstars/\\lk\\ does depend to some extent upon the shape of the assumed stellar potential. Properly taking account of this effect is beyond the scope of this present work, but may bring the synthetic M/L ratios and our results into better agreement. Clearly this is only one of a number of other systematic effects which may also reconcile the slight discrepancy (Table~\\ref{table_syserr}).  \\subsection{Baryon fractions} An interesting result from our analysis is that these systems, despite having masses \\gtsim 5$\\times 10^{12}$\\msun, do not appear in general to be baryonically closed. To some extent this trend was enforced by applying Eq~\\ref{eqn_baryons} to constrain the data. However, the excellent fits we obtained by this method, in conjunction with the good agreement between the measured \\mvir-c relation and the predictions of \\lcdm\\ and, crucially, our measurements at the group scale (which do not employ this restriction: \\citealt{gastaldello06a}), indicate that the inferred \\fbaryons ($\\sim$0.04--0.09; Table~\\ref{table_results}) are accurate. Furthermore, if we relaxed this constraint and instead restricted \\fbaryons\\ to a finite range, we also found that the data tended to favour modest values of \\fbaryons. In particular, for any given system, the measured \\mvir\\ and \\fbaryons\\ were strongly anti-correlated, so that our upper \\mvir\\ constraint is in part imposed by the {\\em lower} limit we place on \\fbaryons. Given the shapes of the \\mvir-c contours (Fig~\\ref{fig_confidence}), it is clear that good agreement with the \\mvir-c relation predicted from simulations tends, therefore, to require rather modest values of \\fbaryons. This would suggest that strong feedback plays an important role in the evolution of these objects.  \\subsection{Temperature profiles} \\label{discussion_temp} By inspection of the temperature profiles (Fig~\\ref{fig_temp}) it is immediately clear that, for all of the galaxy-scale systems in our sample the temperature profiles have negative gradients. In contrast the group-scale objects have positive temperature gradients, similar to observations of other X-ray bright groups and clusters \\citep{gastaldello06a,vikhlinin05b,piffaretti05a}. This radical difference in the temperature profiles seems consistent with our division of galaxies and groups at \\mvir $\\sim 10^{13}$\\msun. The origin of this distinct demarcation between objects around $10^{13}$\\msun\\ is unclear, however. Negative temperature gradients are expected for isolated galaxies containing relatively cool gas, such as that arising from stellar mass-loss. In the deep stellar potential well, compressive heating of the gradually inflowing gas can dominate over radiative cooling to produce a negative temperature slope. In contrast, if hotter ($\\sim$1--2~keV) baryons are allowed to flow in, radiative cooling dominates to produce a positive temperature gradient \\citep{mathews03a}. It is by no means clear, however, why the hot baryons appear to be present only in the systems with \\mvir\\gtsim $10^{13}$\\msun. One possibility is the local environment; all of the galaxy-scale objects are rather isolated, whereas the groups NGC\\thin 4472 and NGC\\thin 4649, in particular, are found in a relatively dense cluster environment, which could provide a reservoir of hot baryons. However, such an explanation cannot easily account for the positive temperature gradient in NGC\\thin 1407, which is comparatively isolated, or the isolated system NGC\\thin 4555, which appears only slightly more massive than our galaxies.   It is possible that selection effects may have played some role in the bimodal temperature profile behaviour, since both NGC\\thin 4125 and NGC\\thin 6482 are classified in \\ned\\ as LINERS, and NGC\\thin 720 has a dominant young stellar population (Appendix~\\ref{sect_stars}). However, none of these systems show strong X-ray morphology disturbances in the core, which might indicate a substantial energy input from star-formation or AGN activity. In any case, the cooling time in the core of NGC\\thin 720 is only $\\sim$200~Myr, substantially less than the implied time since the last major burst of star formation, and so the negative temperature gradient cannot simply be related to energy injection during a starburst. Furthermore, at least two of the group-scale systems also harbour AGN and do not show obvious negative temperature gradients in the core.  Another example of an object we believe to be a galaxy (rather than a group) which exhibits a negative temperature gradient  is the S0 NGC\\thin 1332 \\citep{humphrey04b}. A possible counter-example to this trend might by the ``isolated elliptical galaxy'' NGC\\thin 4555, which exhibits a temperature profile akin to the groups in our sample \\citep{osullivan04b}. However, as we discuss in \\S~\\ref{discussion_groups}, this probably has comparable \\mvir\\ to the groups.   Another intriguing feature of two of the group scale objects is a central temperature peak, similar to a feature we found in the cluster A\\thin 644 \\citep{buote05a}. In that system, we found a significant offset between the X-ray centroid and the emission peak in an otherwise fairly relaxed object. We suggested that both of these features may be related to the cD ``sloshing'' in the potential well of the cluster, which is relaxing following disturbance by, for example, a merger. We do not find obvious evidence of a similar offset in either NGC\\thin 1407 or NGC\\thin 4649. However, these groups may be in a comparably more relaxed (evolved) state than A\\thin 644. Alternatively, the central peaks may be related to past AGN activity heating the gas in the core of the galaxies, from which the system has had time to relax dynamically but not cool completely.    \\subsection{Is NGC\\thin 1407 a ``dark group''?} \\label{sect_ngc1407} Based on the group member dispersion velocity \\citet{gould93a} suggested that NGC\\thin 1407 may lie in a massive (\\gtsim a few $\\times 10^{13}$\\msun) DM halo. Although such a conclusion was strongly dependent on the association of the galaxy NGC\\thin 1400, which exhibits a large peculiar velocity, with the group, we can now confirm the presence of a substantial DM halo around this system. Both the temperature profile and our best-fit mass are similar to the bright X-ray group NGC\\thin 5044 \\citep{buote06a}, and yet it is almost 2 orders of magnitude fainter in \\lx. NGC\\thin 5044 appears to be close to baryonic closure \\citep{mathews05a}, and so has likely retained most of its large gaseous halo. On the other hand NGC\\thin 1407 is not baryonically closed (we estimate \\fbaryons$\\simeq$0.06) and so the loss of much of its hot gas envelope easily explains its lower \\lx/\\lb. Since the masses of the two systems are not considerably different, this points to substantial variation in the evolutionary history of these two groups. In particular, feedback may have operated more efficiently in evacuating the gas from NGC\\thin 1407.  \\citeauthor{gould93a}'s preferred mass estimate ($\\sim10^{14}$\\msun) would imply a remarkably high M/L ratio for the system (\\mvir/\\lb$\\sim$900\\msun/\\lsun), making NGC\\thin 1407 a bona fide ``dark group''. The existence of such an object would provide a valuable insight into the process of star formation in DM halos, as it would imply star formation was somehow inhibited in that system. This mass estimate is, however, considerably larger than our preferred value $\\sim 1.5 \\times 10^{13}$\\msun, which implies a more modest M/L ratio (\\mvir/\\lb$\\sim$140\\msun/\\lsun). To some extent, though, our constraint on \\fbaryons, which was necessary to obtain interesting \\mvir\\ constraints,  has probably enforced this behaviour. Such a restriction may not be valid in a system with an unusual star-formation history and so we experimented with freeing \\fbaryons. To enable \\mvir\\ to be constrained,  we restricted c to lie on the best-fit \\mvir-c relation found by \\citet{bullock01a}. The best-fitting mass, \\mvir$=(9.7^{+17.8}_{-6.2})\\times 10^{13}$\\msun, was in good agreement with \\citet{gould93a}'s values, but implies a baryon fraction only of $\\sim$0.003. Since this fit was statistically indistinguishable from the preferred model, we cannot determine which mass estimate is more likely."
    },
    "0601/astro-ph0601137_arXiv.txt": {
        "abstract": " ",
        "introduction": "The photometric data sets obtained in the course of large microlensing surveys provide an abundance of opportunities to carry out scientific investigations not related to lensing. In one such project, Sumi~(2004; S04 hereafter) presented extinction maps of the Galactic bulge fields observed in the Optical Gravitational Lensing Experiment (OGLE) during its second phase. These fields range in Galactic longitude between approximately $-11$~deg$<l<11$~deg and cover an area close to 11 deg$^2$. The purpose of this effort was to enable studies of Galactic structure with the same data set.  The method employed by S04 was to determine the location of the red clump (RC) giants in the $V-I, I$ color-magnitude diagram for a given sub-region of sky (presumably with homogeneous or nearly-homogeneous extinction), and use the observed color of the RC as a measure of the reddening. By locating the RC in the color-magnitude diagrams for sub-regions with different amounts of extinction, assuming the RC to be constant in luminosity as well as color, the reddening slope can be measured and extrapolated to zero reddening to obtain the total extinction (with some calibration). As suggested by \\citet{popo00} and confirmed by \\citet{udal03}, S04 also found that the measured reddening slopes $1.9 \\la R_{VI} \\la 2.1$ were significantly flatter than the ``standard'' value of about 2.5. Here $R_{VI}\\equiv A_V/E(V-I)$ is the usual ratio of total to selective absorption. This ``anomalously'' shallow extinction slope provides an explanation for previous claims (e.g., \\citealt{stut99}) that extinction corrected colors (derived according to a standard reddening law) of stars in the bulge were redder than their local counterparts.  The minimum light colors of fundamental mode RR Lyrae (RRab) stars have long been recognized as useful reddening measures \\citep{stur66} because of their low intrinsic dispersions and weak dependences on metallicity. The $V-I$ color was suggested to be a particularly good reddening standard by \\citet{mate95} based on the data of \\citet{liu89}. Recently, \\citet{day02} and \\citet{guld05} have studied the colors of field RR Lyrae stars and found the intrinsic minimum light color to be $(V-I)_{0, ml} = 0.58\\pm 0.02$ (consistent with the value used by \\citealt{mate95}), with a scatter of 0.024~mag.  S04 considered the possibility that the intrinsic colors of RC stars might vary somewhat from field to field, due to a possible weak dependence on metallicity; in order to test this he compiled a list of RRab stars. In S04's original analysis, the precision of the $V-I$ colors of the RR Lyrae stars was limited, because the available $V$-band measurements were averages of small numbers of observations taken at random (and unknown) phases. In this work we present and analyze the catalog of RRab stars in the Galactic bulge. Using individual $V$-band time series, we extract measurements of $(V-I)_{ml}$ for bulge RR Lyrae stars and use them to evaluate the field-to-field zero point accuracy of S04's extinction map, among other applications.  In \\S\\ref{sec:data} we discuss the OGLE data, in \\S\\ref{sec:selection} we present the sample selection procedure, and in \\S\\ref{sec:analysis} we analyze the light curves. In \\S\\ref{sec:vmi} we extract and analyze the $(V-I)_{ml}$ colors. In \\S\\ref{sec:disc} we discuss the extinction zero points, distance scale, and geometry of the inner Galaxy, and in \\S\\ref{sec:conc} we summarize our main conclusions. ",
        "conclusions": "\\label{sec:conc}  We have presented a catalog of 1888 fundamental mode RR Lyrae (RRab) stars extracted from the OGLE-II Galactic bulge data set, plus 25 double entries of stars detected in two fields for a total of 1913 entries. The catalog includes basic light curve parameters such as periods, amplitudes, mean magnitudes, $V-I$ measurements at minimum light, labels to indicate Blazhko phenomenology, and significant additional frequencies detected close to the main pulsation frequency, as well as extinctions derived from the map of S04 and various additional information useful for assessing the quality of the photometric data for individual stars. This data set has a variety of applications for studies of the inner Galaxy.  Our frequency analysis (\\S\\ref{sec:freq}) of the light curves has revealed a high incidence of the Blazhko phenomenon: 27.6\\% of the stars show some type of clear Blazhko behavior and an additional 4.8\\% show evidence of unstable main pulsation frequencies, which may be indicative of Blazhko periods longer than the time baseline of the observations (about four years). The Blazhko incidence rate we measure, which is only a lower limit, is somewhat higher than what has been found previously among RRab stars in the bulge, a fact that we attribute to the higher quality of our data set and the sensitivity of our search method. More strikingly, we have obtained a much higher ratio of BL2 stars (with symmetrical frequency triplets, i.e., one additional significant frequency on either side of the main pulsation frequency) to BL1 stars (with only a single additional significant frequency) than has been found previously. Within the limits of our sensitivity, we find BL2s to be about 1.7 times more common than BL1s, whereas previous studies of the bulge \\citep{mize03} and the LMC \\citep{alco03} found BL2/BL1 incidence ratios of 0.59 and 0.83 respectively. Since we believe our detections of additional frequencies to be robust, we attribute this large increase of the BL2/BL1 ratio to the specific and very sensitive method we employ to detect additional, symmetrical frequency components. Furthermore, we have identified several instances of RRab stars that have two pairs of symmetrical frequency triplets, that is, stars with two separate sets of BL2 sidebands. These frequency quintuplets are unevenly spaced. This phenomenon may be important for understanding the long-unsolved origin of the Blazhko effect.  From the comparison of $(V-I)_{ml}$ (minimum light) colors of bulge RR Lyrae stars, dereddened according to S04, with field RR Lyrae colors, we conclude that there is a discrepancy of 0.05--0.08~mag between the RR Lyrae-to-red clump (RC) color differential of the bulge population (measured from OGLE data) as compared to the local population. The sense of the effect is that the color separation is greater for the bulge population. We evaluate likely sources of systematic color errors and conclude that the color discrepancy is probably real. If this is correct then the most likely explanation seems to be the influence of metallicity on the color of the RC, or possibly the RR Lyrae, although there is some evidence against the latter possibility. We have observed a weak dependence of $(V-I)_{0,ml}$ on period (redder at longer periods), but find no evidence of a further dependence on metallicity as estimated from the light curves. The trend with period cannot account for the observed color discrepancy. While the color discrepancy does not affect the conclusion that the reddening slope toward the bulge is anomalously flat, it does cast some doubt on the zero-point accuracy of S04's reddening map. If the RR Lyrae are more reliable reddening tracers, then the $E(V-I)$ values reported by S04 should be reduced by approximately 0.05~mag, with some variation from field to field as indicated in Table~4. We additionally measure an unexpectedly high star-to-star scatter of $(V-I)_{0,ml}$ about 0.07~mag, which is larger than the photometric uncertainties and probably results from unresolved structure in the extinction or other random-type errors in the reddening map.  We exploit the approximately uniform mean $V$-band luminosities of the RRab stars to study the distance to and geometry of the inner Galaxy, and we compare these results to what S04 obtained from studying the RC. Since a coherent picture has not emerged, we summarize some relevant facts below.  \\begin{itemize}  \\item{The Galactocentric distance modulus has been measured to be $14.5\\pm0.1$~mag by \\citet{eise03} from the orbit of a star around the central black hole.}  \\item{Correcting for extinction according to S04, RC stars indicate a distance modulus to Baade's Window (which should be virtually identical to the Galactocentric distance modulus) of $14.86\\pm 0.04$~mag, measured in the $I$-band, assuming that they have the same luminosities as local RC stars. Likely corrections to the reddening zero point (and concomitant corrections to the extinction, assuming that the measured reddening slopes hold all the way to $E(V-I)=0$, for which there is no direct observational evidence) have the wrong sign to bring the distance modulus down.}  \\item{The distance modulus to Baade's Window measured from RRab stars is somewhere between approximately 14.4--14.8~mag, depending mainly on the absolute magnitude calibration adopted. This is measured in the $V$ band.}  \\item{Both the RC and the RR Lyrae stars clearly reveal the signature of the Galactic bar in fields with Galactic longitude $|l|<3$. While the sense of the effect is the same (brighter at positive longitude as compared to negative longitude), the slopes are different. The magnitude difference from $l\\approx -3^{\\circ}$ to $l\\approx 3^{\\circ}$ is approximately 0.35--0.4~mag for the RC and 0.2--0.25~mag for the RR Lyrae.}  \\item{Both the RC and RR Lyrae magnitudes differ by 0.1--0.2~mag between fields approximately symmetrically located above and below the Galactic plane, with the fields at positive latitude being fainter. This assumes there is no significant extinction zero-point mismatch between the fields on opposite sides of the Galactic plane.}  \\item{There is a significant mismatch between the RC and RR Lyrae magnitude trends between about $3^{\\circ}<l<10^{\\circ}$, increasing toward larger $l$, with the RR Lyrae becoming fainter and the RC becoming brighter; the discrepancy between the two trends reaches approximately 0.25~mag at $l\\approx 10^{\\circ}$.}  \\end{itemize}"
    },
    "0601/astro-ph0601071_arXiv.txt": {
        "abstract": "Fast rotation seems to be the mayor factor to trigger the Be phenomenon. Surface fast rotation can be favored by initial formation conditions, such as abundance of metals. We have observed 118 Be stars up to the apparent magnitudes $V=9$ mag. Models of fast rotating atmospheres and evolutionary tracks were used to interpret the stellar spectra and to determine the stellar fundamental parameters. Since the studied stars are formed in regions that are separated enough to imply some non negligible gradient of galactic metallicity, we study the effects of possible incidence of this gradient on the nature as rotators of the studied stars. ",
        "introduction": "In the present paper we would like to know whether the content of metals in the formation regions of stars can introduce some signature in the ($\\tau/\\tau_{\\rm MS},M/M_{\\odot}$) diagram [$\\tau/\\tau_{\\rm MS}\\!=$ fractional age the stars can spent in the main sequence (MS) evolutionary phase]. In fact, it was shown by Ste\\c pie\\'n (2002) that magnetic fields not exceeding 400 G can spin up early type stars in the PMS life, while stronger magnetic fields spin them down. However, since the mechanism acts through magnetic mass-accretion and magnetic-disc locking, the efficacy of the magnetic interaction can differ according to the metallic content in the star and its gaseous environment.\\par It has been shown that Be stars rotate at $\\Omega/\\Omega_c\\sim0.9$ (Fr\\'emat et al. 2005). This rotational rate can be attained in the MS evolutionary phase only if stars have high rotational rates early in the ZAMS. It is expected then that the presence of magnetic fields and its effectiveness at establishing high initial stellar surface rotations can be different according to the metallic content of the medium where they are formed: {\\it age vs. mass} distributions may then be somewhat different. Our aim is to study Be stars situated towards the {\\it galactic-center} and in the {\\it anti-center} directions and see whether some information can be drawn form the {\\it age vs. mass} distributions.\\par ",
        "conclusions": "The evolutionary tracks used to calculate masses and ages are for rotating stars (Meynet \\& Maeder 2000, Zorec et al. 2005).  The distribution of fractional ages ($\\tau/\\tau_{\\rm MS}\\!=$ age/time spent in the MS) against the stellar mass of all studied Be stars is shown in Fig.~\\ref{f1}a. In this figure we see that {\\it all studied stars, but one, lay below the TAMS}. Although most of the studied stars lay in the second half of the MS strip, a non negligible number of them is still in the first half of the MS evolutionary phase. Quite similarly to what was found in Zorec et al. (2005) {\\it a lack of stars with masses $M \\la 7M_{\\odot}$ below $\\tau/\\tau_{\\rm MS} =$ 0.5 is noticeable}. In our sample, stars with masses $M \\ga 12M_{\\odot}$ approach the TAMS. This may be due to their fast evolution and perhaps to some lack in our stellar sample of massive Be stars with ages $\\tau/\\tau_{\\rm MS} \\la$ 0.5. This lack can also be produced by a loss of angular due to their high mass-loss rates. They can become low rotators rapidly and thus be impeded to display the Be phenomenon any more.\\par In order to separate the stars in two sets carrying possible information on differences in initial matallicity content, we separated them into ``galactic-center\" and ``galactic anti-center\" groups. Fig.~\\ref{f1}b shows that there is no noticeable difference in the {\\it ``age vs. mass\"} distributions thus obtained. The most striking differences are: a) the number of Be stars in the \"galactic-center\" group outnumbers the ``anti-center\" one; b) there are younger Be stars in the anti-center direction. Finding a) may agree with the fact that low metallicity favors fast rotation (Maeder et al. 1999). However, the distinction between both groups could be more reliable if we could see differences in the number fractions N(Be)/N(B+Be).\\par"
    },
    "0601/astro-ph0601065_arXiv.txt": {
        "abstract": "{ We report on an {\\it XMM-Newton} observation of the supernova remnant (SNR) \\object{DEM~L241} in the Large Magellanic Cloud. In the soft band image, the emission shows an elongated structure, like a killifish, with a central compact source. The compact source is point-like, and named as XMMU~J053559.3$-$673509. The source spectrum is well reproduced with a power-law model with a photon index of $\\Gamma = 1.57$ (1.51--1.62) and the intrinsic luminosity is $2.2\\times 10^{35}~\\mathrm{ergs~s^{-1}}$ in the 0.5--10.0~keV band, with the assumed distance of 50~kpc. The source has neither significant coherent pulsations in $2.0\\times 10^{-3}$~Hz--8.0~Hz, nor time variabilities. Its luminosity and spectrum suggest that the source might be a pulsar wind nebula (PWN) in DEM~L241. The spectral feature classifies this source into rather bright and hard PWN, which is similar to those in Kes~75 and B0540$-$693. The elongated diffuse structure can be divided into a ``Head'' and ``Tail'', and both have soft and line-rich spectra. Their spectra are well reproduced by a plane-parallel shock plasma ({\\it vpshock}) model with a temperature of 0.3--0.4~keV and over-abundance in O and Ne and a relative under-abundance in Fe. Such an abundance pattern and the morphology imply that the emission is from the ejecta of the SNR, and that the progenitor of DEM~L241 is a very massive star, more than 20~M\\sun. This result is also supported by the existence of the central point source and an OB star association, LH~88. The total thermal energy and plasma mass are $\\sim 4\\times 10^{50}$~ergs and $\\sim 200$~M\\sun, respectively. } ",
        "introduction": "Supernovae (SNe) and supernova remnants (SNRs) shape and enrich the chemical and dynamical structure of the interstellar medium and clouds. X-ray studies give us plenty of information about hot plasma in SNRs with emission lines from highly ionized ions. Moreover, SNRs are believed to be cosmic ray accelerators around their pulsar and pulsar wind nebula (PWN), and/or shock fronts. The Magellanic Clouds (MCs) are the best galaxies for the systematic study of SNRs, thanks to the known distance \\citep[50~kpc;][]{feast1999} and the small absorption column. Another subject of interest is supernova explosions in starburst galaxies. The SNe of massive stars scatter light elements into such galaxies, and make an abundance pattern different from those in normal galaxies \\citep[e.g.,][]{umeda2002}. The Large Magellanic Cloud (LMC) is the nearest starburst galaxy \\citep{vallenari1996}, and thus we can examine the influence of discrete massive SNe on the galaxy. Now, more than 30 LMC SNRs are cataloged \\citep{williams1999} and the number may still increase \\citep[e.g.,][]{chu2004}. However, we have only few samples for which type and age are known. The pulsars and their nebulae are a clear evidence of core-collapsed origin and are good indicators of their age. Now, only four LMC SNRs are reported to have a pulsar and/or PWN; \\object{B0540$-$693} \\citep{manchester1993}, \\object{N157B} \\citep{wang2001}, \\object{B0532$-$710} \\citep{klinger2002,williams2005}, and \\object{B0453$-$685} \\citep{gaensler2003}. This is mainly because of a lack of spatial resolution in previous radio and X-ray observations. Thus, searching for new SNR containing pulsars and PWNe in the LMC is critical to carrying out a systematic study of the contributions from SNRs to interstellar medium in the LMC.   DEM~L241 (0536$-$67.6) was identified as an SNR by \\citet{mathewson1985}. They found that DEM~L241 shows the typical looped filamentary structure of an SNR with a size of $\\sim 2\\arcmin~$ in the optical image. The [\\ion{S}{ii}] to H$\\alpha$ ratio is 0.6 \\citep{mathewson1985}, which is typical of that for SNRs. Radio emission in 843~MHz was also reported by \\citet{mathewson1985}, which shows shell-like structure with similar size to the optical shell. {\\it Einstein} detected relatively bright X-rays \\citep[$2.7\\times 10^{35}~\\mathrm{ergs~s^{-1}}$ in the 0.15--4.5~keV band,][]{mathewson1985}. It appears to be a blow-out of the dense \\ion{H}{ii} region, N~59B, around the OB star association \\object{LH~88} \\citep{chu1988}. A study of the echelle spectrum by \\citet{chu1997} shows a relatively small expansion velocity and a large intrinsic velocity width, implying that the SNR may be expanding within the stellar-wind blown cavity, which suggests that the progenitor of DEM~L241 is likely to be a massive star. \\citet{nishiuchi2001} observed this SNR with {\\it ASCA}, and found that the spectrum of the remnant requires not only thermal emission but also a power-law component with $\\Gamma\\sim 1.6$. This result indicates that DEM~L241 has a pulsar and/or a PWN, or synchrotron X-ray emitting shells like SN~1006 \\citep{koyama1995,bamba2003}. However, the lack of spacial resolution of {\\it ASCA} prevented us from concluding the origin of the hard X-ray emission.  In this paper, we report the detailed X-ray analysis of DEM~L241 for the first time using {\\it XMM-Newton}, with the help of {\\it ASCA} data. \\S\\ref{sec:obs} summarizes the details of {\\it XMM-Newton} and {\\it ASCA} observations of DEM~L241. The analysis results are written in \\S\\ref{sec:results}. \\S\\ref{sec:discuss} is devoted to the discussion about the origin of the remnant. We assume the distance to the LMC to be 50~kpc \\citep{feast1999} in this paper. ",
        "conclusions": "\\label{sec:discuss}  \\subsection{The absorption}  The best-fit $\\mathrm{N_H{}^{LMC}}$ in our analysis is almost consistent with or slightly larger than the LMC absorption column estimated from the 21~cm line survey \\citep[$8.4\\times 10^{20}~\\mathrm{cm^{-2}}$;][]{rohlfs1984}, implying that this source is in the LMC. The small excess might be due to the \\ion{H}{i} excess reported in the \\ion{H}{i} map by \\citet{stavely-smith1998}. However, we can conclude nothing due to the lack of high sensitivity information around the SNR. More \\ion{H}{i} observations with better spatial resolution are required to discuss this issue quantitatively.  \\subsection{Source~1}  The central position of Source~1, the hard spectrum, and persistent luminosity indicate that Source~1 might be a pulsar and/or a PWN of DEM~L241. The lack of detection of coherent pulsation is not surprising, since the periods of ordinary rotation-powered pulsars detected with X-rays are $\\lse$0.1~sec. The hard spectrum also resembles the emission from background active galactic nuclei (AGNs), although the source shows no time variability. We estimate the probability of chance coincidence that there is a background AGN accidentally in the SNR region. With the $\\log N$---$\\log S$ relation of AGNs derived by \\citet{hasinger1998}, we found that the expected number of AGNs is only $1.1\\times 10^{-3}$ in the $1.5\\arcmin\\times 3\\arcmin$ region. Therefore, we concluded that Source~1 is an LMC member. The similarity of the absorption column for Source~1 to that of the Head region also support that. In the case that Source~1 is a pulsar and/or a PWN, this adds a new member to the sample of LMC PWNe \\citep{gaensler2003}. The hard X-ray flux detected by {\\it ASCA} \\citep{nishiuchi2001} is consistent with that of Source~1. Thus we consider that is not nonthermal emission from energetic electrons accelerated on the shell, but the central point source. The distance from Source~1 to the center of LH~88 is about 14\\arcsec, and the diaphragm diameter of LH~88 is 100\\arcsec~ \\citep{bica1996}, so the progenitor might be a member of the cluster.  The intrinsic luminosity of Source~1 is $2.2\\times 10^{35}~\\mathrm{ergs~s^{-1}}$ in the 0.5--10.0~keV band under the assumption that the distance to Source~1 is 50~kpc, and the photon index is 1.57 (1.51--1.62). The luminosity and hardness ranges of pulsars are wide, $\\log L_x \\sim $~31--36, $\\Gamma \\sim$~0.6--2.0, respectively \\citep[e.g.,][]{gotthelf2002,gaensler2003}. Recently, ``compact central sources'' have been discovered in several SNRs (Cas~A, \\citealt{chakrabarty2001}; Vela~Jr., \\citealt{kargaltsev2002}; Kes~79, \\citealt{seward2003}), which have rather small luminosities ($\\log L_x \\sim $~32--34) and soft spectra ($\\Gamma$ = 3--4). In the case that Source~1 is a pulsar or a compact central source, Source~1 is one of the brightest samples, and as bright as the Crab pulsar \\citep{pravdo1997}. Such bright pulsars certainly have PWNe, which are about 10~times brighter than the pulsars themselves and also have hard spectra \\citep[$\\log L_x = $33--37 and $\\Gamma = $1.3--2.3;][]{gotthelf2002}. Thus, we conclude that the emission is from a PWN of DEM~L241. Its photon index is rather small for ordinary PWNe. and similar to those of the PWN in \\object{Kes~75} \\citep{helfand2003} and \\object{B0540$-$693} \\citep{hirayama2002}. \\citet{helfand2003} notes that these pulsar systems change their energy into X-ray radiation very efficiently. Source~1 may be a new sample of such energetic PWNe. Source~1 is moderately bright for PWNe, which might imply that the system is middle-aged. Pulsars in such PWNe have periods of less than a few hundred msec and luminosities of about 10\\% of their nebulae. Therefore it is natural that we could not detect the coherent pulsation even in the case that Source~1 has a pulsar.  In order to confirm the origin of Source~1, it is essential to resolve the pulsar spatially and to detect coherent pulsations. Therefore, further observations are encouraged with better time and spatial resolution, and excellent statistics, in the X-ray and radio bands.  \\subsection{Diffuse emission}  The temperature and abundance pattern of the hot plasma in the Head and Tail regions are roughly statistically the same as each other, (Table~\\ref{tab:spec_diffuse}), implying that these plasmas have the same origin. The overabundant O and Ne relative to the average LMC indicate that the plasma emission is from the ejecta of the explosion, and that relative to Fe implies that the progenitor of DEM~L241 may be a core-collapsed SN \\citep{tsujimoto1995}. The Fe abundance is slightly smaller than the avarage value in the LMC \\citep[0.22][]{hughes1998}, it might be the local fluctuation of metal. The less abundant O relative to Ne might imply that the progenitor is very massive, $\\gse 20 M$\\sun~ \\citep{umeda2002}. \\citet{umeda2002} also insists that fast convective mixing in the progenitor star \\citep{sprut1992} makes more Ne and less O. Therefore, we concluded that the progenitor of DEM~L241 was very massive star. This conclusion is also supported by the presence of the energetic pulsar candidate and the OB star association, LH~88 \\citep{chu1988}. The abundance pattern of DEM~L241 resembles in that of a starburst galaxy \\object{M82} \\citep{tsuru1997}. This fact might imply that DEM~241 have been produced in a period of starburst activity of the OB star association. \\citet{umeda2000} and \\citet{nakamura2001} have shown that a large Si/O ratio may be signature of energetic SN explosion ($\\gse 10^{51}$~ergs). Thus, we need to determine Si abundance with low background observations.  The thermal emission in the Head region has a center-filled morphology even considering the contribution from Source~1. Note that the radius of the extent of the thermal emission is $\\sim$70\\arcsec~ while that of Source~1 corresponds to the {\\it XMM-Newton} PSF, $\\sim$5\\arcsec. The structure, the radio shell \\citep{mathewson1985}, and the large size of the SNR (see \\S\\ref{subsec:image}), classify the SNR as a mixed-morphology (MM) SNRs \\citep{rho1998}. However, the over-abundant light elements suggest that the emission is not from interstellar medium but ejecta as already mentioned, which is not common for ordinary MM SNRs. Then, we concluded that this SNR is an ejecta dominant SNR without X-ray emitting shells. The absence of a limb-brightened X-ray shell is expected, since this SNR is in a hot and tenuous bubble \\citep{chu1997} and it is hard to form the strong shock. The elongated shape of the Tail region, together with the lack of enhancement of [\\ion{S}{ii}] and radio continuum emission around Tail, might indicate that the plasma was ejected an-isotropically into a low-density region. It is natural to consider that the plasma in the Tail region is a blown-out on the basis of the [\\ion{S}{ii}] morphology. \\citet{chu1997} suggested that the SNR could be interacting with the inner walls of a bubble to produce the SNR signature. In such a case, the interacting region is heated and emits thermal X-rays, and the emission certainly traces the shock region, like the [\\ion{S}{ii}] emission. If the thermal emission is, on the other hand, center-filled, then we conclude that the thermal X-rays are not from the interacting region but from the ejecta.  In order to estimate physical parameters of the plasma, we assumed that the plasma in the Head region is distributed in a sphere with the radius of 70\\arcsec, excluding a spherical region with the radius of 30\\arcsec~ around Source~1 (total volume $V_{Head} = 5.6\\times 10^{59}~\\mathrm{cm}^3$), and that the Tail region has an ellipsoid distribution with the radii of 50\\arcsec$\\times$50\\arcsec$\\times$90\\arcsec~ ($V_{Tail} = 4.0\\times 10^{59}~\\mathrm{cm}^3$), and calculated the mean electron density ($n_e$), the age of the plasma ($t_p$), the thermal energy ($E \\simeq 3n_eVkT$), and the total mass of the plasma ($M = n_em_pV$, where $m_p$ is the proton mass). The results are summarized in Table~\\ref{tab:diffuse_sum}. The $t_p$ indicates that the SNR is rather aged, which is consistent with the large size of the SNR."
    },
    "0601/astro-ph0601586_arXiv.txt": {
        "abstract": "We present far-ultraviolet (FUV) spectral-imaging observations of the Vela supernova remnant (SNR), obtained with the Spectroscopy of Plasma Evolution from Astrophysical Radiation (\\emph{SPEAR}) instrument, also known as \\emph{FIMS}.  The Vela SNR extends $\\sim8^{\\circ}$ in the FUV and its global spectra are dominated by shock-induced emission lines.  We find that the global FUV line luminosities can exceed the 0.1--2.5 keV soft X-ray luminosity by an order of magnitude.   The global O~{\\small VI}:C~{\\small III} ratio shows that the Vela SNR has a relatively large fraction of slower shocks compared with the Cygnus Loop. ",
        "introduction": "The Vela supernova remnant (SNR) has been studied in great detail due to its proximity \\citep[250pc;][]{vela_distance} and its large angular diameter of $8^{\\circ}$ \\citep[]{VelaROSAT}.  The remnant is $\\sim$10,000 years old \\citep{vela_age} and its overall emission is dominated by the interaction of the SN blast wave with the interstellar medium (ISM), but also has a pulsar and plerionic nebula near its center.  Vela is one of two Galactic SNRs (the other being the Cygnus Loop) that has been extensively studied in the UV.  The UV emission from SNRs arises primarily in shocks driven by the supernova blast wave into interstellar clouds.  The shock velocities responsible for the UV emission lie in the range $\\sim$50-300 km s$^{-1}$, and heat the gas to temperatures of ~$10^5 - 10^6$ K.  The hot shocked gas can also be studied via absorption so long as there are suitable background continuum sources.  Absorption line studies of Vela using \\textit{Copernicus} showed the presence of O~{\\small VI} and N~{\\small V}, high ionization species expected in shocked gas, and also high velocity components of lower ionization species \\citep{Jenkins1976a,Jenkins1976b}.  Absorption line studies have also been carried out with \\textit{IUE} \\citep{Jenkins,Nichols}, \\emph{HST} \\citep{Jenkins1995,Jenkins1998} and \\textit{FUSE} \\citep{Slavin}.  These studies have shown the existence of fast shocks ($\\gtrsim$160 km s$^{-1}$) distributed widely but inhomogeneously across the face of the remnant, and also the variations in the dynamic pressure driving these shocks.  Emission line studies of Vela have been carried out using \\textit{IUE} \\citep{Vela_IUE}, \\textit{HUT} \\citep{RaymondVela}, \\textit{FUSE} \\citep{ravi2,ravi} and \\emph{Voyager 2} UVS \\citep[henceforth BVL]{vela_voyager}.  These observations, except for the ones obtained by \\emph{Voyager 2}, were of regions with angular extents of order an arc-minute and probed the properties of individual shock fronts.  BVL analyzed \\emph{Voyager 2} spectra of a $1.5^{\\circ} \\times 2.0^{\\circ}$ region in the northern part of Vela, which showed strong C~{\\small III} and O~{\\small VI} emission features.  They found variations in the flux ratios on scales of 0.2$^{\\circ}$.  They also concluded from the overall flux ratio between the two lines that slower shocks (those unable to produce O~{\\small VI}) were more prevalent in Vela than in the Cygnus Loop.  \\textit{FUSE} spectra of a few regions in Vela well separated from each other have strong C~{\\small III} lines and relatively weak or no O~{\\small VI}, and show the presence of slower shocks ($100 - 140$ km s$^{-1}$)  spread over the face of the remnant \\citep{raviproc}.  We present spectral images of the Vela SNR in several FUV lines and FUV spectra of the entire remnant.  The data were obtained with \\emph{SPEAR} (The Spectroscopy of Plasma Emission from Astrophysical Radiation), also known as \\emph{FIMS} (Far-ultraviolet Imaging Spectrograph).  \\emph{SPEAR}, launched Sepember 27, 2003 on the Korean satellite STSAT-1, is a dual-channel FUV imaging spectrograph (S channel 900 - 1150 \\AA, L channel 1350 - 1750 \\AA, $\\lambda/\\Delta\\lambda\\sim$550) with a large imaged field of view (S: 4.0$^{\\circ} \\times 4.6'$, L: 7.5$^{\\circ} \\times 4.3'$, spatial resolution $\\sim10'$) optimized for the observation of diffuse emission \\citep[see][for an overview of the instrument and mission.]{Instrument,Mission}.  A large effective field of view can be obtained by sweeping across the sky.  This combination of instrument properties yields a dataset that supplements data obtained in previous UV studies of the remnant.  The data allow us to estimate the total flux from Vela in several FUV lines.  In the following sections we present the observations (\\S2), discuss the spectral images (\\S3) and the total spectrum (\\S4).  The last section (\\S5) summarizes the importance of these data for the study of SNRs. ",
        "conclusions": "We have presented global FUV spectral images and spectra of the Vela SNR.  Our data indicate inhomogeneous shock-induced emission from the SNR surface.  The images show limb brightening and knots of emission.  The spectra are consistent with past FUV observations of Vela. The global FUV luminosities of emission lines can exceed the soft X-ray luminosity by an order of magnitude, providing an efficient cooling channel to the SNR.  The data will be used to compare the global FUV morphology with images in other wavebands and examine how the SNR evolves and interacts with the ambient ISM.  \\emph{SPEAR} spectra toward previously observed regions \\citep[ex.][]{RaymondVela,ravi2} will be modeled in detail to determine the physical properties of the regions.  Mapping the O~{\\small VI}:C~{\\small III} ratio will allow us to map the distribution of shock velocities across Vela.  Furthermore, the global \\emph{SPEAR} spectra can be compared with those of neighboring regions to investigate Vela's association with the Gum nebula."
    },
    "0601/astro-ph0601253_arXiv.txt": {
        "abstract": "Large surveys of very metal-poor stars have revealed in recent years that a large fraction of these objects were carbon-rich, analogous to the more metal-rich CH-stars. The abundance peculiarities of CH-stars are commonly explained by mass-transfer from a more evolved companion.  In an effort to better understand the origin and importance for Galactic evolution of Fe-poor, C-rich stars, we present abundances determined from high-resolution and high signal-to-noise spectra obtained with the UVES instrument attached to the ESO/VLT. Our analysis of carbon-enhanced objects includes both CH stars and more metal-poor objects, and we explore the link between the two classes. We also present preliminary results of our ongoing radial velocity monitoring. ",
        "introduction": "A subgroup of carbon stars, the CH stars, were first distinguished from other carbon stars 60 years ago (\\cite{keenan42}). Their chief characteristics are : strong CH absorption lines, strong Swan bands of C$_2$, strong resonance lines of Ba~II and Sr~II (later shown to reflect genuine overabundances of s-process elements), weak lines of the iron-group elements and high proper motion.  Since it was shown that all field CH stars were spectroscopic binaries (\\cite{mcclure90}), the binary scenario was accepted: CH stars have been polluted by a nearby companion, formerly on the AGB, now a defunct white dwarf.  But some recent results have cast doubts on this general picture. Conflicting evidences include the radial velocity monitoring of targets from the HK survey (\\cite{beers92}) which show no radial velocity variations (\\cite{preston01}), and the discovery of carbon but not s-process enhanced stars (e.g. \\cite{sneden96}) . Furthermore, the unexpectedly high fraction of carbon-enhanced stars among metal-poor stars discovered in the HK survey (about 25\\% in the metallicity range [Fe/H] $\\le$ -2.5 compared to 1-2 \\% among stars of higher metal abundances), make them crucial for the study of the early stages of Galaxy evolution.  \\begin{table} \\caption{The sample and the atmospheric parameters.} \\label{table} \\begin{center} \\leavevmode \\footnotesize \\begin{tabular}{l|c|c|c|c} \\hline star     &T$_{\\rm eff}$  & log~g &  [Fe/H]  & $^{12}$C/$^{13}$C  \\\\ \\hline CS 22891-171      &     5100  &  1.8  &  -2.31   &  5     \\\\ CS 22942-019      &     5100  &  2.5  &  -2.48   &  7     \\\\ CS 22945-017*    &     6400  &  4.3  &  -2.60   &  4     \\\\ CS 22953-003      &     4600  &  1.0  &  -3.27   & $>$5   \\\\ CS 22956-028*    &     6700  &  3.5  &  -2.38   &  3     \\\\ CS 30322-023      &     4000  &  0.0  &  -3.78   &  5     \\\\ HD 26*                &     5200  &  2.6  &  -0.71   &  ?     \\\\ HD 5424*            &     5000  &  3.0  &  -0.41   &  4    \\\\ HD 24035*          &     5000  &  3.7  &   0.16   &  4     \\\\ HD 168214          &     5200  &  3.5  &  -0.03   &  ?     \\\\ HD 187861*        &     4800  &  1.8  &  -2.21   & 10    \\\\ HD 196944*        &     5250  &  1.7  &  -2.21   &  ?     \\\\ HD 206983          &     4550  &  1.6  &  -0.93   &  4     \\\\ HD 207585*        &     5800  &  4.0  &  -0.35   & 10    \\\\ HD 211173          &     4800  &  2.5  &  -0.17   & 13    \\\\ HD 218875          &     4600  &  1.5  &  -0.63   & 100   \\\\ HD 219116          &     4800  &  1.8  &  -0.45   &  7      \\\\ HD 224959*        &     5000  &  2.0  &  -2.08   &  7     \\\\ HE 1419-1324     &     4900  &  1.8  &  -3.28   &  4      \\\\ HE 1410+0213* &     4550  &  1.0  &  -2.38   &  4      \\\\ HE 1001-0243*   &     5000  &  2.3  &  -3.12   & 30     \\\\ \\hline \\multicolumn{5}{l}{* probably binary, from radial velocity monitoring}\\\\ \\end{tabular} \\end{center} \\end{table} ",
        "conclusions": ""
    },
    "0601/astro-ph0601409_arXiv.txt": {
        "abstract": "We have obtained optical and infrared photometry of the quiescent soft X-ray transient XTE J1118+480. In addition to optical and $J$-band variations, we present the first observed $H$- and $K_s$-band ellipsoidal variations for this system. We model the variations in all bands simultaneously with the WD98 light curve modeling code. The infrared colors of the secondary star in this system are consistent with a K7V, while there is evidence for light from the accretion disk in the optical. Combining the models with the observed spectral energy distribution of the system, the most likely value for the orbital inclination angle is $68^{\\circ}\\pm2^{\\circ}$.  This inclination angle corresponds to a primary black hole mass of $8.53\\pm0.60 \\,M_{\\odot}$. Based on the derived physical parameters and infrared colors of the system, we determine a distance of 1.72$\\pm$0.10 kpc to XTE J1118+480. ",
        "introduction": "Transient low mass X-ray binaries (LMXBs) exhibit large and abrupt X-ray and optical outbursts that can be separated by decades of quiescence \\citep{csl97}.  For these systems, the compact object is a black hole or a neutron star, and the companion is normally a low-mass K- or M-type dwarf-like star \\citep{cc03}.  During their periods of quiescence, these systems are faint at X-ray, optical, and infrared (IR) wavelengths.  While the quiescent X-ray emission can be caused by accretion onto a compact object or thermal emission from the surface of a neutron star \\citep{garcia01,wijnands04}, the companion can dominate the luminosity at optical and IR energies.  During the binary orbit, the changing aspect of the tidally-distorted companion causes a periodic modulation of the optical and IR emission \\citep{gelino01}. Measurements of these ``ellipsoidal variations'' provide information about the physical parameters of the binary.  XTE J1118+480 ($\\alpha_{2000}$ = 11$^{\\rm h}$18$^{\\rm m}$10.85$^{\\rm s}$, $\\delta_{2000}$ = 48$^{\\circ}$02$\\arcmin$12.9$\\arcsec$) was discovered with the all-sky monitor (ASM) on the {\\it Rossi X-Ray Timing Explorer} by \\citet{rem00} on 2000 March 29, while \\citet{gar00} spectroscopically identified its 12.9 magnitude optical counterpart. Owing to its position in the Galactic halo, this high latitude ({\\it b} = +62$^{\\circ}$) system has been observed by numerous groups over many wavelength regimes.  Recent orbital parameters determined from optical spectra suggest an orbital period of 4.078 hr and a secondary star radial velocity semi-amplitude of 709$\\pm$7 km s$^{-1}$ \\citep{tor04}. These values imply a mass function of {\\it f(M)}=6.3$\\pm$0.2 M$_{\\odot}$, identifying the compact object as a black hole.  Determining a precise black hole mass requires an accurate measurement of the orbital inclination angle of the system.  As discussed in \\citet{gho01}, the best way to find the inclination angle in a non-eclipsing system is to model its infrared ellipsoidal light curves. In the IR regime, there is a smaller chance of contamination from other sources of light in the system.  While modeling several light curves from one wavelength regime helps to constrain model parameters, simultaneously modeling light curves that span more than one wavelength regime provides tighter constraints than modeling IR light curves alone.  Previous inclination estimates for XTE J1118+480 have come from modeling optical ellipsoidal variations as the system approached quiescence \\citep{mcc01, wag01, zur02}.  These inclination angles range from 55$^{\\circ}$\\ \\citep{mcc01} to 83$^{\\circ}$\\ \\citep{wag01}, and correspond to primary masses of 10 M$_{\\odot}$ and 6.0 M$_{\\odot}$, respectively. Since this system has been known to exhibit optical superhumps from the precession of an eccentric accretion disk on its way to quiescence \\citep{zur02}, it is important to determine the orbital inclination angle while XTE J1118+480 is in a truly quiescent state.  In order to determine an accurate orbital inclination angle for XTE J1118+480, we have obtained $B$-, $V$-, $R$-, $J$-, $H$-, and $K_s$-band light curves of the system while in quiescence, and simultaneously model them here with the WD98 light curve modeling code \\citep{wil98}.  To date, this is the most comprehensive ellipsoidal variation data set published for this system.  The modeled inclination angle is combined with recently published orbital parameters to determine a highly constrained mass of the black hole in this X-ray binary. ",
        "conclusions": "Figure \\ref{fig1} presents the resulting $B$-, $V$-, and $R$-band light curves of XTE J1118+480, while the final $J$, $H$, and $K_s$ differential light curves are presented in Figure \\ref{fig2}. Despite the expectation of detecting superhumps from the 52 day accretion disk precession period, there was no evidence of a superhump period when the data were run through a periodogram.  These results are consistent with the findings of \\citet{sha05} whose optical data taken on 2003 June suggested that if a superhump modulation existed, it was at the $<$0.50$\\%$ level.  \\citet{sha05} also observed stochastic variability and fast flares in their light curves.  However, the observed stochastic variability had the same magnitude as their photometric error bars, and despite the significant power in the power density spectrum, the flares seen in XTE~J1118+480 had roughly the same magnitude as the spread in the comparison star measurements. The 2003 and 2004 $V$- and $R$-band light curves presented here were consistent in shape and amplitude.  As Table \\ref{obstab} shows, two thirds of the nights that XTE~J1118+480 was observed, data was gathered for at least 75\\% of an orbit, with more than one full orbit covered on four of the nine nights.  No evidence for unequal maxima, flares, or light curve distortions were found in the unphased data.  The periodogram results are consistent with the results from \\citet{tor04} and therefore all data presented here have been phased to their ephemeris.  These are the first $H$- and $K_s$-band detections of ellipsoidal variations from this X-ray binary system.   \\subsection{Ellipsoidal Models}  The spectral type of the secondary star can be estimated by comparing its red optical spectrum with the spectra of stars with various spectral types from the same luminosity class. Alternatively, one can use a spectral energy distribution (SED), and published limits on the spectral type to not only derive an effective temperature of the secondary star, but to also estimate both the visual extinction and contamination level. Given that photometric data usually has a higher S/N than spectroscopic data, an effective temperature derived using photometry can be just as useful as a spectral type derived from a spectroscopic data set. To this end, we compared the observed optical/IR ($BVRJHK_s$) SED for XTE J1118+480 with the observed SEDs for K0V -- M4V stars \\citep{bb88,bes91,mik82}.  We present the best fitting SED in Figure \\ref{fig3}, and find that it predicts a visual extinction of A$_V$=0.065$\\pm$0.020 mags, and a secondary star spectral type of K7V. It also includes 60 -- 70\\% light from the accretion disk at $B$ and 30 -- 35\\% at $V$. A K5V gives a slightly worse fit, and predicts more $R$-band light than is observed, as well as a smaller amount of disk light in the $B$- and $V$-bands, inconsistent with previously published values. The visual extinction found here is consistent with the column density adopted by \\citet[$N_H=1.2 \\times 10^{20}$ cm$^{-2}$]{mcc04} based on three independent measurements, and the spectral type found here is consistent with those found through spectral fitting \\citep[K5/7 V]{mcc03,tor04}.  Light from the primary object or accretion disk in an X-ray binary will act to dilute the amplitude of the ellipsoidal variations of the secondary star. \\citet{tor04} estimate that the secondary star in the XTE J1118+480 system contributes roughly 55\\% of the total flux between 5800\\AA ~and 6400\\AA; however, at $R$ = 18.6, the system was not in true quiescence when their observations were taken. \\citet{zuriauc02} determined that XTE J1118+480 has a true quiescent $R$-band magnitude of 18.9. % The optical data presented here ($R$=19.00) are consistent with this value, and thus we contend that the system was in a quiescent state during our observations. In addition, Doppler imaging of the system did not detect any H$\\alpha$ emission from a hot spot or accretion stream \\citep{tor04}. Since our observations took place while XTE J1118+480 was in true quiescence, the optical data presented here are consistent with both of the \\citet{tor04} results.  The most difficult contamination source to extract in the case of XTE J1118+480 is the one that has the shallowest spectral slope. If the disk contamination in the infrared is based on the assumption that the optically thin disk radiates through free-free emission processes, and we therefore ascribe the entire $V$-band excess to free-free emission, then in the $K_s$-band, the contamination would be $\\sim$ 8\\%. An 8\\% contamination in the infrared bands would cause the observed orbital inclination angle of the system to be underestimated by 2$^{\\circ}$. What if we instead assume that any IR contamination is ascribed to a jet or some form of an advection dominated accretion flow (ADAF)? Models by \\citet{yua05} for XTE~J1118+480 in quiescence show that both the jet and ADAF flux in the IR are predicted to constitute significantly less than 8\\% of the companion flux.   Irradiation of the secondary star by the accretion disk will affect the symmetry of the ellipsoidal light curves.  \\citet{tor04} investigated the possibility of X-ray irradiation powering the H$\\alpha$ emission seen in their Doppler tomograms of the XTE J1118+480 system.  They concluded that it was improbable that this could be the case, and found that the strength of the H$\\alpha$ emission from the secondary star was comparable with other K dwarfs.  Similarly, we see no significant evidence for irradiation effects in the unphased data or the light curves presented here.  As in \\citet{gho01}, we simultaneously modeled the optical and infrared light curves of XTE J1118+480 with WD98 \\citep{wil98}.  See \\citet{gel01} for references and a basic description of the code, and \\citet{gelino01} for a more comprehensive description.  The data were run through WD98's DC program which uses a Simplex algorithm for initial parameter searches and a damped least squares algorithm for error minimization between the data and model light curves.  We ran the code for a semi-detached binary with the primary component's gravitational potential set so that all of its mass is concentrated at a point. The most important wavelength-independent input values to WD98 are listed in Table~\\ref{tab1}.  The models were run for a range of inclination angles with parameters for a K3V through an M1V secondary star.  The secondary star atmosphere was determined from solar-metallicity Kurucz models.  We used normal, non-irradiated, square-root limb darkening coefficients \\citep{van93}.  We also adopted gravity darkening exponents found by \\citet{cla00}, and mass ratios ranging from $q$ = 0.030 - 0.056 \\citep{har99,oro01}. We assumed that the secondary star was not exhibiting any star spots during our observations. The models were run with varying amounts of additional light in all bands. Solving for six consistent light curve solutions simultaneously allowed the rejection of many disk light scenarios.  With 118 degrees of freedom, the best fit model had a reduced $\\chi^2$ of 1.65.  While the error on the best fit inclination angle was dependent on combining the uncertainties from varying all of the model parameters simultaneously, we found that changing the spectral type of the secondary from a K5V to an M1V resulted in a change in the orbital inclination angle of 1$^{\\circ}$. Similarly, varying $q$ from 0.030 to 0.056 affected $i$ by $\\le 1^{\\circ}$.  We find that the best fitting $B$-, $V$-, $R$-, $J$-, $H$-, and $K_s$-band model has $i$ = 68$^{\\circ}$, and the parameters found in Table \\ref{tab2}.  Figure 1 presents this model for the optical bands, while Figure \\ref{fig2} presents this model for the $J$-, $H$-, and $K_s$-band light curves.   \\subsection{The Model and its Uncertainties}  Based on our optical/IR SED, as well as published values, we adopt a secondary star spectral type of K7 \\citep{tor04, mcc01} with a temperature of $T_{\\rm eff}$ = 4250 K \\citep{gra92}.  The corresponding gravity darkening exponent is $\\beta_1$ = 0.34.  Modeling six light curves simultaneously is a robust method for constraining the amount of extra light in the system. If we model the light curves individually and assume that all of the light in the system comes from the secondary star, the best fitting orbital inclination angle for XTE J1118+480 varies with wavelength.  Since the contaminating light does not have the same spectrum as the secondary star the light curves at each wavelength are diluted by different amounts, causing the best fit inclination for the affected bands to be different from the others.  This is the case for the $B$- and $V$-band light curves.  If a jet or other contaminant were to have a flat spectrum that only affected the IR, the results from the SED would most likely not match those obtained from spectral measurements, and we would not be able to determine a reasonable parameter set and inclination angle that is consistent throughout the data. When fit simultaneously, the best fit inclination angle is $i$ = 68$^{\\circ}$ with 62\\% disk light in $B$, and 31\\% disk light in $V$. These disk light contributions are consistent with those found through the SED fitting.  Using an estimate of 8\\% for the infrared accretion disk contamination in the system gives an inclination of 68$^{+2.8}_{-2}$$^{\\circ}$, however, if we artificially add in 8\\% or more infrared light, we are unable to obtain a reasonable solution for the orbital inclination angle that is consistent throughout the entire data set.  Therefore, based on the IR colors of the system ($J$-$K$=1.1$\\pm$0.3), the SED fit, and the results of the simultaneous light curve modeling, it appears unlikely that the infrared light curves are significantly affected by any such contamination. In order to determine the final error on the orbital inclination angle, we plotted the $\\chi^2$ values as a function of $i$. Therefore, based on the error in each of the model parameters including $q$, the spectral type (i.e. temperature) of the secondary star, the amount of observed disk light, as well as the photometric error bars, the orbital inclination angle is 68$^{\\circ}$$\\pm$2$^{\\circ}$. We combined the determined inclination angle with the orbital period (P = 0.1699167 $\\pm$ 1.72$\\times$10$^{-5}$ d), radial velocity of the secondary star (K$_2$ = 709 $\\pm$ 7 km; \\citet{tor04}), and and the mass ratio ($q$ = 0.0435 $\\pm$ 0.0100) to find the mass of the primary object.  A Monte Carlo routine was used to propagate the errors on the above quantities and gives a primary mass of 8.53$\\pm$0.60 M$_{\\odot}$, confirming it as a black hole.  The constraints on the mass of the black hole in this system presented here represent a considerable improvement over those previously published. \\citet{wag01} determined a mass of 6.0 -- 7.7 M$_{\\odot}$ from data obtained before XTE 1118+480 had entered a quiescent state ($R\\sim$18.3). \\citet{mcc01} gave an upper limit of $M_1 \\le$ 10 M$_{\\odot}$, and more recently, \\citet{mcc04} adopted a mass of $\\sim$ 8 M$_{\\odot}$ for their calculations of the thermal emission from the black hole in the system.  The 8.53 M$_{\\odot}$ black hole mass determined here is consistent with the adopted mass of \\citet{mcc04}, and falls nicely into the current observed black hole mass distribution.  Theoretical models by \\citet{fry01} predict that there should exist a greater number of 3 -- 5 M$_{\\odot}$ black holes than 5 -- 12 M$_{\\odot}$ black holes; however, most of the determined black hole masses thus far have fallen into a 6 -- 14 M$_{\\odot}$ range.  In fact, thus far, GRO J0422+32 is the only system with a compact object that falls into the 3 -- 5 M$_{\\odot}$ range \\citep[3.97$\\pm$0.95 M$_{\\odot}$]{gel03}.  Using the mass of the compact object and the orbital period, we computed the orbital separation of the two components in the system.  We then combined the separation with the mass ratio to find the size of the Roche lobe for the secondary star.  The temperature of the secondary and its Roche lobe radius were then used to find the secondary's bolometric luminosity and bolometric absolute magnitude. After accounting for the bolometric correction \\citep{bes91}, the distance modulus for the $J$, $H$, and $K_s$ bands were used to find an average distance of 1.72$\\pm$0.10 kpc. Consistent with results from the modeling and SED fitting, this calculation assumes that all of the IR light in the system originates from the secondary star. Table~\\ref{tab2} lists all of the derived parameters for the XTE~J1118+480 system.  Both the mass and radius of the secondary star are smaller than that of a K7V ZAMS star.  In addition, the infrared colors of the system and its position in the Galactic halo, support the notion that the secondary star in the XTE J1118+480 system may be evolved."
    },
    "0601/astro-ph0601315_arXiv.txt": {
        "abstract": "We have carried out large scale CO observations with a mm/sub-mm telescope NANTEN toward a far infrared loop-like structure whose angular extent is about 20$\\times$20 degrees around ($l$, $b$) $\\sim$ (109$\\degr$, $-$45$\\degr$) in Pegasus. Its diameter corresponds to $\\sim$ 25 pc at a distance of 100 pc, adopted from that of a star HD886 (B2IV) near the center of the loop.  We covered the loop-like structure in the $^{12}$CO ($J$ = 1--0) emission at 4$\\arcmin$--8$\\arcmin$ grid spacing and in the $^{13}$CO ($J$ = 1--0) emission at 2$\\arcmin$ grid spacing for the $^{12}$CO emitting regions.  The $^{12}$CO distribution is found to consist of 78 small clumpy clouds whose masses range from 0.04 $M_{\\sun}$ to 11 $M_{\\sun}$, and $\\sim$ 83\\% of the $^{12}$CO clouds have very small masses less than 1.0 $M_{\\sun}$.   $^{13}$CO observations revealed that 18 of the 78 $^{12}$CO clouds show significant $^{13}$CO emission. $^{13}$CO emission was detected in the region where the column density of H$_{2}$ derived from $^{12}$CO is greater than 5$\\times$10$^{20}$ cm$^{-2}$, corresponding to $A$v of $\\sim$ 1 mag, which takes into account that of H{\\small \\,I}.  We find no indication of star formation in these clouds in IRAS Point Source Catalog and 2MASS Point Source Catalog. The very low mass clouds, $M$ $\\leq$ 1$M_{\\sun}$, identified are unusual in the sense that they have very weak $^{12}$CO peak temperature of 0.5 K--2.7 K and that they aggregate in a region of a few pc with no main massive clouds; contrarily to this, similar low mass clouds less than 1 $M_{\\sun}$ in other regions previously observed including those at high Galactic latitude are all associated with more massive main clouds of $\\sim$ 100 $M_{\\sun}$. A comparison with a theoretical work on molecular cloud formation (Koyama \\& Inutsuka 2002) suggests that the very low-mass clouds may have been formed in the shocked layer through the thermal instability. The star HD886 (B2IV) may be the source of the mechanical luminosity via stellar winds to create shocks, forming the loop-like structure where the very low-mass clouds are embedded. ",
        "introduction": "High Galactic latitude molecular clouds (hereafter HLCs) are typically located at $\\mid$b$\\mid$ $\\gtrsim$ 20$\\degr$--30$\\degr$.  Since the Gaussiun scale height of CO is estimated to be $\\sim$ 100 pc in the inner Galactic disk  (e.g., Magnani et al. 2000), HLCs are likely located very close to the Sun, within a few hundred pc or less.  Their proximity to the Sun and the low possibility of overlapping with other objects along the line of sight enable us to study them with a high spatial resolution and to compare CO data unambiguously with the data at other wavelengths.  HLCs have lower molecular densities compared with dark clouds where the optical obscuration is significant.  Therefore, HLCs are often called as translucent clouds (e.g., van Dishoeck \\& Black 1988) and most of the known HLCs are not the sites of active star formation, although a few of them are known to be associated with T Tauri stars (e.g., Magnani et al. 1995; Pound 1996; Hearty et al. 1999).   Given the very small distances of HLCs, it is a challenging task for observers to make a complete survey for HLCs over a significant portion of the whole sky. $^{12}$CO ($J$ = 1--0) emission has been used to search for HLCs because the line emission in the mm band is strongest among the thermally or sub-thermally excited spectral lines of interstellar molecular species.  It is however difficult to cover an area as large as tens of square degrees subtended by some of the HLCs because of the general weakness of the $^{12}$CO emission, typically $\\sim$ a few K (e.g., Magnani et al. 1996), with existing mm-wave telescopes in a reasonable time scale.  HLCs have been therefore searched for by employing various large-scale datasets at other wavelengths including the optical obscuration (Magnani et al. 1985; Keto \\& Myers 1986), the infrared radiation (Reach et al. 1994), and the far-infrared excess over H{\\small \\,I} (=FIR excess)(Blitz et al. 1990; Onishi et al. 2001).  On the other hand, unbiased surveys in CO at high Galactic latitudes have been performed at very coarse grid separations of 1$\\degr$ resulting in a small sampling factor of a few \\% (Hartmann et al. 1998; Magnani et al. 2000).  Most recently, Onishi et al. (2001) discovered 32 HLCs or HLC complexes. This search was made based on the FIR excess, demonstrating the correlation among FIR excess clouds with CO clouds is a useful indicator of CO HLCs.   Previous CO observations of individual HLCs at higher angular resolutions show that HLCs exhibit often loop-like or shell-like distributions having filamentary features with widths of several arc min or less (Hartmann et al. 1998; Magnani et al. 2000; Bhatt 2000), and in addition that HLCs often compose a group, whose angular extent is $\\sim$ 10 degrees or larger.  In order to better understand the structure of HLCs and to pursue the evolution of HLC complexes, CO observations covering tens of square degrees at a high angular resolution are therefore crucial.  The past observations of such complexes of HLCs are limited to a few regions including Polaris flare (Heithausen \\& Thaddeus 1990), Ursa Major (Pound \\& Goodman 1997) and the HLC complex toward MBM 53, 54, and 55 (Yamamoto et al. 2003).  Pound \\& Goodman (1997) showed an arc-like structure of the molecular cloud system and suggested that the origin of such structures could be some explosive events.  Most recently, Yamamoto et al. (2003) carried out extensive observations of the molecular cloud complex including MBM 53, 54, and 55 and suggest that the HLCs may be significantly affected by past explosive events based on the arc-like morphologies of molecular hydrogen (see also Gir et al. 1994).   The region of MBM 53, 54, and 55 is of particular interest among the three, because it is associated with a large H{\\small \\,I} cloud of $\\sim$ 590 $M_{\\sun}$ at a latitude of $-$35 degrees and because there is a newly discovered HLC of 330$M_{\\sun}$, HLCG92$-$35, which is significantly H{\\small \\,I} rich with a mass ratio $M$(H$_{2}$)/$M$(H{\\small \\,I}) of $\\sim$ 1,  among the known HLCs (Yamamoto et al. 2003).  This cloud was in fact missed in the previous surveys based on optical extinction (Magnani et al. 1985). Subsequent to these observations we became aware of that the region is also very rich in interstellar matter as shown by the 100$\\mu$m dust features (Kiss et al. 2004). There is a loop-like structure shown at 100 $\\mu$m around ($l$, $b$) $\\sim$ (109$\\degr$, $-$45$\\degr$).  Toward the center of the loop, an early type star HD886(B2IV) is located and may play a role in creating the loop. Its proper motion is large at a velocity of a few km s$^{-1}$, suggesting that the stellar winds of the star might have continued to interact with the surrounding neutral matter over a few tens of pc in $\\sim$ a few Myr. Magnani et al. (1985) and Onishi et al. (2001) yet observed only a small part of this region. In order to reveal the large scale CO distribution of the region, we have carried out observations toward ($l$, $b$) $\\sim$ (109$\\degr$, $-$45$\\degr$) by $^{12}$CO ($J$ = 1--0) and $^{13}$CO ($J$ = 1--0) with NANTEN 4-meter millimeter/sub-mm telescope of Nagoya University at Las Campanas, Chile. We shall adopt the distance of 100 pc from the sun to the loop-like structure which is equal to the distance of the B2 star in the center of the loop, and is also a typical value for the HLCs. ",
        "conclusions": "We have made a large-scale survey of high Galactic latitude molecular clouds in the $J$ = 1--0 lines of $^{12}$CO and $^{13}$CO toward a large scale structure located around ($l$, $b$) $\\sim$ (109$\\degr$, $-$45$\\degr$) with NANTEN.  This survey spatially resolved the distribution of molecular gas associated with the large scale structure.  The main conclusions of the present study are summarized as follows:   \\begin{enumerate} \\item  The $^{12}$CO observation covered the entire large loop-like structure.  The loop-like structure consits of very small clumpy clouds. The $^{12}$CO clouds are concentrated on the north to north-west of the loop-like structure and toward the south of that.  We identified 78 $^{12}$CO clouds in the observed region.  The total mass is estimated to be $\\sim$ 64 $M_{\\sun}$ if we assume the conversion factor from CO intensity to $N$(H$_2$) as 1.0$\\times$10$^{20}$ cm$^{-2}$/(K km s$^{-1}$).  \\item  We performed $^{13}$CO observations in and around the whole area where the peak temperature of $^{12}$CO is more than 2.0 K.  We identified 33 $^{13}$CO clouds and derived physical properties under the assumption of LTE.  \\item  The mass spectra are well fitted by a power law, $dN/dM$ $\\propto$ $M^{-1.53\\pm0.13}$ for the $^{12}$CO clouds and $dN/dM$ $\\propto$ $M^{-1.36\\pm0.10}$ for the $^{13}$CO clouds. These spectral indices are similar to those derived in the other regions.  \\item  The size and the line width relation of $^{13}$CO clouds is fitted by a least-squares method, log($\\Delta V$) = (0.22$\\pm$0.43) $\\times$ log($R$) + (0.37$\\pm$0.52) (c.c.=0.23), but the correlation is not good.  \\item  Present $^{13}$CO clouds are far from the virial equilibrium, indicating that $^{13}$CO clouds are not gravitationally bound. $M_{\\rm vir}$ and $M_{\\rm LTE}$ relation can be fitted by a least-squares method as log($M_{\\rm vir}$) = (0.91$\\pm$0.30) $\\times$ log($M_{\\rm LTE}$) + (2.23$\\pm$0.29) (c.c.=0.66).  This index is slightly different from the indices in the other regions although the tendency that molecular clouds are more vilialized as the mass increases is consistent with the other regions.  \\item  There is no sign of star formation from the comparison of IRAS point sources and Point Source Catalog of Two-Micron All-Sky survey in the present region. This suggests that molecular clouds in this region are not the site of present star formation or the remnants of past star formation.  \\item  There may be two expanding shells in the present region as inferred from H{\\small \\,I}  although we cannot identify them from CO.  The total mechanical luminosity of HD886 during the last few $\\times$ 10$^{6}$ yr is comparable to the expanding energy of the northern expanding H{\\small \\,I} shell.  This indicates that some additional source of energy other than HD886 is needed to explain the expanding energy.  \\item  $^{13}$CO emission is significantly detected in the $^{12}$CO clouds having molecular column density greater than 5$\\times$10$^{20}$ cm$^{-2}$. This may be explained as that the $^{13}$CO emitting regions become significant when $A$v becomes larger than $\\sim$ 1 mag, marginally enough to shield the ultraviolet radiation to protect $^{13}$CO molecules.  \\item  There is a possibility that very small clouds have been formed in the shoked layer through the thermal instability. The stellar wind of HD886 may be the source to creat shocks, forming the loop-like structure where the very small clouds are embedded.   \\end{enumerate}"
    },
    "0601/astro-ph0601645_arXiv.txt": {
        "abstract": "\\par We present new rates for the \\iso{22}Ne$(\\alpha, n$)\\iso{25}Mg and \\iso{22}Ne$(\\alpha,\\gamma)$\\iso{26}Mg reactions, with uncertainties that have been considerably reduced compared to previous estimates, and we study how these new rates affect the production of the heavy magnesium isotopes in models of intermediate mass Asymptotic Giant Branch (AGB) stars of different initial compositions. All the models have deep third dredge-up, hot bottom burning and mass loss. Calculations have been performed using the two most commonly used estimates of the \\iso{22}Ne $+ \\, \\alpha$ rates as well as the new recommended rates, and with combinations of their upper and lower limits. The main result of the present study is that with the new rates, uncertainties on the production of isotopes from Mg to P coming from the \\iso{22}Ne $+ \\alpha$-capture rates have been considerably reduced. We have therefore removed one of the important sources of uncertainty to effect models of AGB stars.  We have studied the effects of varying the mass-loss rate on nucleosynthesis and discuss other uncertainties related to the physics employed in the computation of stellar structure, such as the modeling of convection, the inclusion of a partial mixing zone and the definition of convective borders. These uncertainties are found to be much larger than those coming from \\iso{22}Ne $+ \\, \\alpha$-capture rates, when using our new estimates. Much effort is needed to improve the situation for AGB models. ",
        "introduction": "\\par The origin of the stable magnesium isotopes, \\iso{24}Mg, \\iso{25}Mg and \\iso{26}Mg, is of particular interest to astrophysics because Mg is one of the few elements for which we can obtain isotopic information from stellar spectroscopy.  The ratio \\iso{24}Mg:\\iso{25}Mg:\\iso{26}Mg has been derived from high-resolution spectra of cool dwarfs and giants in the thin and thick disk of the Galaxy \\citep*{gl00,yong03b}, and for giants stars in the globular cluster (GC) NGC 6752 \\citep{yong03a}. These observations show that many of the stars, including relatively metal-poor stars ([Fe/H] $\\lesssim -1.0$), have non-solar Mg isotopic ratios\\footnote{The solar Mg isotopic ratios are \\iso{24}Mg:\\iso{25}Mg:\\iso{26}Mg = 79:10:11 \\citep{lodder03}.} with enhancements in the neutron-rich isotopes, \\iso{25}Mg and \\iso{26}Mg, compared to what is expected from galactic chemical evolution (GCE) models. The main stellar nucleosynthesis site for all three stable isotopes is hydrostatic burning in the carbon and neon shells of massive stars that explode as Type II supernovae \\citep{ww95,cl04}. The abundances of the neutron-rich Mg isotopes are further enhanced by secondary $\\alpha$-capture processes operating in the helium shell. The amount of \\iso{24}Mg produced does not strongly depend on the initial metallicity of the model and is an example of primary nucleosynthesis\\footnote{Primary production means that the species is produced from the hydrogen and helium initially present in the star and the amount produced is relatively independent of the metallicity, $Z$, whereas secondary production requires some heavier seed nuclei to be present, and the amount produced scales with $Z$.}, whereas the amounts of \\iso{25}Mg and \\iso{26}Mg produced scale with the initial metallicity of the star. This means that the Mg content of the  ejecta of low metallicity supernovae will be mostly \\iso{24}Mg with very little \\iso{25}Mg and \\iso{26}Mg produced, with typical ratios \\iso{24}Mg:\\iso{25}Mg:\\iso{26}Mg $\\approx$ 99.0:0.50:0.50 from a 25$\\Msun$ supernova model with metallicity $Z=0.01 Z_{\\odot}$  \\citep{ww95}. \\par Previous studies have shown that GCE models using ejecta from massive stars match the observational data well for [Fe/H] $> -1.0$ but severely underestimate \\iso{25,26}Mg/\\iso{24}Mg at lower metallicities \\citep*{timmes95}, indicating that another production site for the neutron-rich Mg isotopes at low metallicities is required to account for the observations.  Recently, \\citet{fenner03} included the predicted stellar yields of \\iso{25}Mg and \\iso{26}Mg from Asymptotic Giant Branch (AGB) stars \\citep{kara03} along with yields from Type II supernovae \\citep{ww95,lc02} into a GCE model of the solar neighborhood.  The GCE model with the AGB contribution could successfully match the Mg isotopic ratios of the metal-poor Galactic disk stars while the model without an AGB contribution could not. This result indicates that low-metallicity intermediate mass AGB stars may play an important role in the production of these species in galaxies and stellar systems. The production of the neutron-rich isotopes in AGB stars is also of interest in relation to the non-solar Mg isotopic ratios observed in giant stars in globular clusters \\citep{yong03a,yong05}. The non-solar Mg isotopic ratios observed in NGC 6752 have been attributed to AGB stars, but \\citet{fenner04} used a GCE model with tailor made AGB yields from \\citet{campbell05} and failed to match the abundance patterns observed in stars in that cluster. As \\citet{ventura05a,ventura05b} have pointed out, there are still many uncertainties that effect the stellar yields, so an AGB solution to the globular cluster anomalies cannot be ruled out at present. \\par Further motivation for the study of the production of the Mg isotopes in AGB stars is given by their relevance in the important current debate on the apparent variation of the fine-structure constant \\citep*{murphy01,ashen04,fenner05}, and the origin of pre-solar spinel grains, some of which show enhancements in both \\iso{25}Mg and \\iso{26}Mg compared to solar \\citep{zinner05}. \\par Briefly, intermediate-mass stars (initial mass $\\sim$4 to 8$\\Msun$) will enter the thermally-pulsing phase with a hydrogen (H)-exhausted core mass (hereafter core mass) $M_{\\rm c} \\gtrsim$ 0.8$\\Msun$, after experiencing the second dredge-up (SDU) \\citep{latt96,bgw99,herwig05}. The SDU brings the products of H-burning to the stellar surface (mostly \\iso{4}He and \\iso{14}N), and will slightly alter the composition of the Mg isotopic ratios, with an enrichment in \\iso{26}Mg at the expense of \\iso{25}Mg. For the massive, $Z = 0.0001$ models, the SDU is the first time the surface abundances are altered, because there is essentially no first giant branch phase \\citep{herwig04a}. Following the SDU, the He-burning shell becomes thermally unstable and flashes every few thousand years or so. The energy from the thermal pulse (TP) drives a convective pocket in the He-rich intershell, which thoroughly mixes the products of He-nucleosynthesis within this region. Following the TP, the convective envelope moves inward in mass and may reach the region previously mixed by the flash-driven convective pocket. This mixing event is known as the third dredge-up (TDU) and if it occurs, is responsible for enriching the envelope in material from the H-exhausted core.  Following the TDU  the star contracts and the H-shell is re-ignited and provides nearly all of the surface luminosity for the next interpulse period.  The thermal pulse -- TDU -- interpulse cycle may occur many times during the TP-AGB phase; how many times depends on a number of factors including the convective model which determines the surface luminosity and mass-loss rate and hence the total AGB lifetime \\citep{ventura05a}. \\par Hot bottom burning can also occur when the base of the convective envelope becomes hot enough to sustain proton-capture nucleosynthesis \\citep{latt96}. If temperature at the base of the envelope is sufficiently hot (over $\\sim 60 \\times 10^{6}$\\,K), the NeNa and MgAl chains may operate alongside the CNO cycle; \\iso{7}Li production is also possible via the Cameron-Fowler mechanism \\citep{sb92,latt96}, which operates at lower temperatures (typically $\\sim 30-40 \\times 10^{6}\\,$K). \\citet{frost98} noted that intermediate-mass AGB stars may become luminous, optically obscured carbon stars near the end of the TP-AGB, when mass loss has removed much of the envelope, extinguishing HBB but allowing dredge-up to continue. \\par \\citet{kara03} described in detail the various nucleosynthesis processes that alter the Mg isotopic ratios in AGB stars. To summarize, \\iso{25}Mg and \\iso{26}Mg are synthesized in the He-shell during thermal pulses by the reactions \\iso{22}Ne($\\alpha,n$)\\iso{25}Mg and \\iso{22}Ne($\\alpha,\\gamma$)\\iso{26}Mg, when the temperature exceeds about 300$\\times 10^{6}\\,$K. The amount of Mg produced depends on the thermodynamic conditions inside the pulse as well as on the composition of the intershell, which will have been altered by previous H and He-burning. Neutron captures, in particular the  \\iso{25}Mg($n, \\gamma$)\\iso{26}Mg reaction, can also alter the Mg isotopic ratio in the intershell, where the neutrons come from the \\iso{22}Ne($\\alpha, n$)\\iso{25}Mg reaction \\citep{herwig04b}. HBB can also significantly alter the surface Mg isotopic ratio via the activation of the MgAl chain, which can result in the destruction of \\iso{24}Mg if the temperature exceeds $\\sim 90 \\times 10^{6}\\,$K. \\par The stellar yields of \\iso{25,26}Mg presented in \\citet{kara03} and shown in figure~\\ref{fig:mg-previous} were calculated from models covering a range in mass (1$\\Msun$ to 6$\\Msun$) and metallicity ([Fe/H] $= 0, -0.3, -0.7$). From this figure we see that the most massive AGB models produce the most \\iso{25,26}Mg, as a consequence of higher temperatures in the He-shells compared to lower mass stars; we also notice an increase in production at a given mass with a decrease in metallicity.  The computationally demanding nature of AGB models precluded a detailed study in that paper of the effect of the major uncertainties (mass loss, nuclear reaction rates, convection). A recent comprehensive study by \\citet{ventura05a,ventura05b} demonstrated that the predictive power of AGB models is still seriously undermined by these uncertainties. The theory of convection has a significant effect on the structure and nucleosynthesis \\citep{ventura05a}, whilst varying the mass-loss rate results in larger changes to the stellar yields than varying the nuclear reaction rates \\citep{ventura05b}. However, the magnitude of the errors associated with the relevant nuclear reaction rates is still one of the key questions concerning predictions of magnesium production in intermediate-mass AGB models.  Whilst the rates of the \\iso{14}N($\\alpha,\\gamma$)\\iso{18}F and \\iso{18}O($\\alpha,\\gamma$)\\iso{22}Ne reactions are well determined \\citep{gorres00,dab03}, the two key $\\alpha$-capture reaction rates: \\iso{22}Ne($\\alpha,n$)\\iso{25}Mg and \\iso{22}Ne($\\alpha,\\gamma$)\\iso{26}Mg, suffer from large uncertainties at  the stellar energies appropriate for AGB stars \\citep{koehler}.  For example, at typical He-shell burning temperatures, $T \\approx 300 \\times 10^{6}\\,$K, the NACRE compilation \\citep{nacre} give an upper limit  to the \\iso{22}Ne($\\alpha,n$)\\iso{25}Mg reaction rate that is about 47 times larger than the recommended rate. At lower temperatures, the uncertainties are even larger. \\par The aims of this paper are two fold. First, we present new reaction rates for the two key $\\alpha$-capture reactions, with considerably reduced uncertainties compared to those given by the NACRE compilation. Second, we use these new rates within models of different metallicities ([Fe/H] $\\approx 0, -0.3, -0.7, -2.3$), for a typical mass (5$\\Msun$) that produces the Mg isotopes during the thermally-pulsing AGB (TP-AGB) phase. For each model we calculate the stellar yields and compare to previous nucleosynthesis calculations using older estimates of these reaction rates, including those by \\citet{herwig04b} and \\citet{ventura05a,ventura05b}. We also examine the effect of other model uncertainties, in particular the inclusion of a partial mixing zone and mass loss, on the stellar yields. \\par The paper is organized as follows. In \\S\\ref{section:method} we discuss the numerical method used for the stellar model calculations, including a discussion of the input physics used. In \\S\\ref{section:rates} we present new rates for the \\iso{22}Ne $+ \\alpha$  capture reactions. The results of the calculations are presented in \\S\\ref{section:results} and discussed in \\S\\ref{section:discussion}. ",
        "conclusions": "\\label{section:discussion}  \\par The main important result of the present study is that the reduction of the uncertainties on the \\iso{22}Ne $+ \\, \\alpha$ reaction rates has allowed us to considerably reduce the uncertainties coming from these rates on the production of isotopes from Mg to P in AGB stars of intermediate mass. The uncertainties on the Mg yields are now at a level of $\\sim 30$\\%, much lower than those obtained when using the NACRE upper or lower limits. The yields of \\iso{25}Mg and \\iso{26}Mg are 20\\% to 45\\% and 9\\% to 16\\%, respectively, smaller with the new rates, as compared to NACRE. The uncertainties have also been reduced for species heavier than Mg, where for example, the uncertainty on the production of P in solar metallicity models is now at a level of 35\\%, much lower than the 400\\% obtained when using the NACRE upper limit. These results are clearly illustrated by figures~\\ref{yield-compare1} and~\\ref{yield-compare2} and in table~\\ref{tab:mgyields}. \\par From our analysis it would appear that among the uncertainties related to the stellar models, those coming from the treatment of convection and of the mass loss are the largest. We described earlier that these are enormous when compared to the current uncertainties coming from the \\iso{22}Ne $+ \\, \\alpha$ reaction rates. Much effort is needed to improve the situation for AGB models, in particular with respect to convection, by trying to evaluate and reduce the uncertainties, perhaps by exploiting all the available observational constraints. \\par Our new evaluation of the \\iso{22}Ne $+ \\alpha$ reaction rates will encourage much future work, as these rates are important for many nucleosynthetic processes and sites. The \\iso{22}Ne($\\alpha, n$)\\iso{25}Mg reaction is an important source of neutrons both during the final evolutionary stages of massive stars, and during the AGB phase of low to intermediate mass stars. It is responsible for the production of heavy s-process elements in these environments \\citep{gallino98,rauscher02}. The new rate and its uncertainties have to be tested in relation to this process.  Moreover, the smaller uncertainties of the rates presented here with respect to those in NACRE, appear to rule out the possibility that the production of the relatively abundant p-only isotopes of Mo and Ru could be related to a high value of the \\iso{22}Ne($\\alpha, n$)\\iso{25}Mg rates \\citep{costa00}. \\par It is important to study the relative production of \\iso{25}Mg and \\iso{26}Mg,  as we have done in table~\\ref{tab:mgyields}, because both spectroscopic observations and the analysis of pre-solar grains are able to separate these two isotopes. Also the contribution of radioactive \\iso{26}Al to the abundance of \\iso{26}Mg has to be carefully evaluated. In particular, one pre-solar spinel grain \\citep[][OC2]{zinner05} appears to bear the signature of nucleosynthesis in intermediate mass AGB stars, with excesses in both  \\iso{25}Mg and \\iso{26}Mg. A future application of our present work will be to compare our detailed results to the composition of this grain, extending the study to the oxygen isotopic ratios in AGB stars and their uncertainties, as these are also measured in the grain."
    },
    "0601/astro-ph0601535_arXiv.txt": {
        "abstract": "{ {We present a multiwavelength study of the massive star forming region associated with IRAS 06055+2039 which reveals an interesting scenario of this complex where regions are at different stages of evolution of star formation. Narrow band near-infrared (NIR) observations were carried out with UKIRT-UFTI in molecular hydrogen and Br$\\gamma$ lines to trace the shocked and ionized gases respectively. We have used 2MASS $J H K_{s}$ data to study the nature of the embedded cluster associated with IRAS 06055+2039. We obtain a power-law slope of 0.43$\\pm$0.09 for the $K_{s}$-band Luminosity Function (KLF) which is in good agreement with other young embedded clusters. We estimate an age of 2 -- 3 Myr for this cluster. The radio emission from the ionized gas has been mapped at 610 and 1280 MHz using the Giant Metrewave Radio Telescope (GMRT), India. Apart from the diffuse emission, the high resolution 1280 MHz map also shows the presence of several discrete sources which possibly represent high density clumps. The morphology of shocked molecular hydrogen forms an arc towards the N-E of the central IRAS point source and envelopes the radio emission. Submillimetre emission using JCMT-SCUBA show the presence of a dense cloud core which is probably at an earlier evolutionary stage compared to the ionized region with shocked molecular gas lying in between the two. Emission from warm dust and the Unidentified Infrared Bands (UIBs) have been estimated using the mid-infrared (8 -- 21\\,$\\mu$m) data from the MSX survey. From the submillimetre emission at 450 and 850\\,$\\mu$m the total mass of the cloud is estimated to be $\\sim$ 7000 -- 9000 $\\rm M_{\\odot}$. } ",
        "introduction": "Massive stars are preferentially formed in dense cores of molecular clouds. They remain deeply embedded in the prenatal molecular gas and obscuring dust and their pre-main sequence time scales are much shorter compared to the low mass stars. The luminous high mass stars also affect the parent cloud. In addition, massive stars do not form in isolation but often in clusters and associations. All these factors contribute in making the study of the formation mechanisms of these systems very difficult. Multiwavelength studies, therefore, hold the potential to probe these complexes at different depths and unravel the least understood aspects of massive star formation. IRAS 06055+2039 (G189.78+0.34, RAFGL 5179) is a massive star forming region chosen from the catalog of massive young stellar objects by Chan et al. (1996). G189.78+0.34 is listed as an ultracompact (UC) HII region (Shepherd \\& Churchwell 1996; Bronfman et al. 1996). It belongs to the Gem OB1 molecular cloud complex and is a part of the extended HII region Sh 252. It is associated with S252 A which is one of the six compact radio sources in Sh 252 revealed from the 5 GHz aperture synthesis observations by Felli et al. (1977). IRAS fluxes yield a far infrared luminosity of $\\sim 10^{4}\\rm L_{\\odot}$ (Carpenter et al. 1995) for this star forming region. There are several kinematic distance estimates used in literature for this source which range from 1.25 kpc (Mirabel et al. 1987) to 2.9 kpc (Wouterloot \\& Brand 1989). In this paper we use the value of 2.6 kpc (Wu et al. 2001) which is the most widely used distance estimate for this source.  This high mass star forming region has been observed as part of many surveys. H$_{2}$O maser (K\\\"{o}mpe et al. 1989; Lada \\& Wooden 1979) and the 6.7 GHz methanol maser (Szymczak et al. 2000a) have been detected towards IRAS 06055+2039. Positive detection has also been made in SiO (Harju et al. 1998), CO (2 - 1) (Wu et al. 2001) and CS (2 - 1) (Bronfman et al. 1996). Zinchenko et al. (1998) in their study of dense molecular cores have also mapped this source in CS (2 - 1). CO maps of Shepherd \\& Churchwell (1996) do not show evidence of any bipolar outflows. Search for the 6 cm (Szymczak et al. 2000b) and 5 cm (Baudry et al. 1997) OH masers show negative results. The IRAS low resolution spectra show a relatively red continuum from 13 to 23\\,$\\mu$m with the presence of an emission feature at 11.3\\,$\\mu$m which is attributed to the presence of Polycyclic Aromatic Hydrocarbon (PAH) molecules (Kwok et al. 1997).  In this paper, we present a multiwavelength study of this star forming region. In Sect.\\,\\,\\ref{obvs.sec}, we present the narrow-band near-infrared (NIR) and radio continuum observations and a brief description of the related data reduction procedures. In Sect.\\,\\,\\ref{arch.sec}, we discuss other available datasets used in the present study. Section\\,\\,\\ref{results.sec} gives a comprehensive discussion on the results obtained and in Sect.\\,\\,\\ref{concl.sec}, we summarize the results. ",
        "conclusions": "\\label{results.sec} \\subsection{Embedded Cluster in the Near-Infrared} \\subsubsection{Radial Profile and Stellar Surface Number Density} \\label{ssnd.sect} The 2MASS $K_{s}$ - band image of the region around IRAS 06055+2039 is shown in Figure\\,\\,\\ref{2mass_k.fig}. We see the presence of a diffuse emission region harbouring an infrared cluster. \\begin{figure} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{fig4.ps}} \\caption{ The 2MASS $K_{s}$ - band image of the region around IRAS 06055+2039. The presence of a diffuse emission region and a NIR cluster is seen. The plus sign marks the position of the IRAS point source.} \\label{2mass_k.fig} \\end{figure}   This cluster is also listed in the catalog of embedded infrared clusters compiled by Bica et al. (2003). We use the 2MASS $JHK_{s}$ data to study the nature of this cluster. To determine the $K_{s}$ - band radial profile and the stellar surface number density (SSND), we have selected sources which are detected in the $K_{s}$ band. To estimate the cluster radius we select a large region of radius $\\rm 300\\arcsec$ centered on IRAS 06055+2039 ($\\rm \\alpha_{2000.0} = 06^{h} 08^{m} 32^{s}.1\\,;\\,\\delta_{2000.0} = +20^{\\circ} 39\\arcmin 18\\arcsec$). To account for the contribution from the field stars we select a control field ($\\rm \\alpha_{2000.0} = 06^{h} 09^{m} 52^{s}.0\\,;\\,\\delta_{2000.0} = +20^{\\circ} 39\\arcmin 18\\arcsec$) which is $\\sim$ $\\rm 20\\arcmin$ to the east of IRAS 06055+2039. Figure\\,\\,\\ref{cl_rad.fig} shows the radial profile of the stellar density in log-log scale. This profile was created by counting the number of stars in $\\rm 10\\arcsec$ annuli and normalizing by the annulus area. We fit two models to the surface density radial profile -- the King's profile and the inverse radius ($r^{-1}$) model. Neglecting the tidal radius, the King's profile can be written as \\begin{equation} f(r) = a + \\frac{f_{0}}{1+(r/r_{c})^{2}} \\end{equation} where, $f_{0}$ is the core concentration at radius zero, $r_{c}$ is the core radius and $a$ is a constant for the background offset. As seen in Fig.\\,\\,\\ref{cl_rad.fig}, both the models describe the density distribution fairly well. However, the $r^{-1}$ model has a better overall fit (reduced $\\chi^{2}$ = 1.2) as compared to the King's model (reduced $\\chi^{2}$ = 3.3). Several studies have shown that young embedded clusters have been fitted by $r^{-1}$ profiles (e.g McCaughrean \\& Stauffer 1994; Lada \\& Lada 1995) or, by both the inverse radius model and the King's model (e.g Horner et al. 1997; Teixeira et al. 2004; Baba et al. 2004). As discussed in Baba et al. (2004), the $r^{-1}$ dependence is likely to be a reminiscent of the parental cloud core whereas the King's profile represents systems in dynamical equilibrium. Hence, the better fitting of the $r^{-1}$ model could possibly suggest that the cluster associated with IRAS 06055+2039 is not yet in complete dynamical equilibrium.  Within errors, the cluster profile merges with the field star level at $\\rm \\sim 85\\arcsec$ which translates to $\\sim$ 1.1 pc at a distance of 2.6 kpc. We take this as the cluster radius. The background level as estimated from the control field is $\\sim \\rm 9\\, stars\\, pc^{-2}$ which is in reasonable agreement with the value of $\\rm 11.8\\pm1.6\\, stars\\, pc^{-2}$ yielded by the King's profile fitting. The King's profile fitting also gives a core radius $\\rm r_{c} \\,\\sim 0.1\\, pc$. The core radius is a scale parameter and depends mainly on cluster parameters like density, luminosity, total mass etc. Several studies have shown that the core radius of the cluster is also correlated with its age. In their study of young clusters in the LMC, Elson et al. (1989) show that the core radii increases between $\\sim 10^{6}$ and $10^{9}$ yr and then begin to decrease again. Such trends in core radius evolution has also been discussed by Wilkinson et al. (2003) for the LMC clusters and Mackey \\& Gilmore (2003) for clusters in the SMC. These authors have shown that apart from the general increasing trend, the spread in the core radii also increases with the age of the cluster. Teixeira et al. (2004) and Baba et al. (2004) derive core radius values of 0.05 and 0.08 pc for the clusters NGC 2316 and RCW 36 respectively. The age estimate for both these clusters is 2 -- 3 Myr. For a comparatively older (5 -- 10 Myr) cluster NGC 2282, Horner et al. (1997) obtain a core radius of 0.19 pc. The above values of core radii seem to suggest that the cluster associated with IRAS 06055+2039 is likely to be of the same age ($\\sim$ 2 -- 3 Myr) as NGC 2316 and RCW 36.  The number of stars detected in the $K_{s}$ band within $\\rm 85\\arcsec$ radius of the cluster is 114. The total background population is 34 [$\\rm = \\,9 (background) \\times \\pi \\times (1.1)^{2} (area \\,of\\, cluster)$]. Hence, the total number of cluster members is estimated to be 80. This yields a space density of $\\sim$ 14 stars\\,pc$^{-3}$. However, it should be noted here that this density is a lower limit as we are not completely sampling the stellar population at fainter magnitudes due to the 2MASS sensitivity limits.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fig5.ps}} \\caption{The radial profile of the surface number density for the cluster associated with IRAS 06055+2039 in log-log scale. Also plotted are the two fitted models -- King's model (solid) and inverse radius (dashed). As compared to the King's model, the inverse radius model fits the radial profile much better. The horizontal dotted line corresponds to the background field star level which is $\\sim$ 9 stars pc$^{-2}$. Statistical errors are shown. } \\label{cl_rad.fig} \\end{figure}  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fig6.ps}} \\caption{The contour map of the stellar surface number density obtained by counting stars in a $\\rm 10\\arcsec \\times 10\\arcsec$ ( $\\sim$ 0.1 $\\times$ 0.1 pc) grid for the cluster associated with IRAS 06055+2039. The contours are from 30 to 140 stars pc$^{-2}$ in steps of 20 stars pc$^{-2}$. The lowest contour is at the 3$\\sigma$ level. The positions of the IRAS point source (plus), the radio peak (open triangle) and the sub-mm peak (star) are shown in the figure.} \\label{cl_sdm.fig} \\end{figure}  Figure\\,\\,\\ref{cl_sdm.fig} shows the surface density map of the region around IRAS 06055+2039. We see the presence of two peaks. The prominent peak which lies to the west coincides spatially with the central bright source. The secondary peak, which lies $\\rm \\sim 50\\arcsec$ to the east of the main peak, coincides well with the departure seen in the surface density profile from the model profiles around a radial distance of $\\sim$ 0.6 pc. This secondary peak is situated close to the edge of the sub-mm peaks presented in Fig.\\,\\,\\ref{jcmt.fig}. The stellar surface density distribution exhibits a centrally condensed-type structure rather than a hierarchical-type structure (Lada \\& Lada 2003). This is consistent with the fact that the derived radial profile fits reasonably well both to the King's model as well as the inverse radius model. Hierarchical-type complexes do not follow any well defined profile. Though at a less significant level, centrally condensed clusters have also been seen to show the presence of structures. For example, Lada \\& Lada (1995) show the presence of satellite subclusters in the outer regions of IC 348. This is consistent with the double peaked structure that we see for the cluster associated with IRAS 06055+2039.  \\subsubsection{Colour-Colour (CC) and Colour-Magnitude (CM) Diagrams} \\label{cccmd.sect} The ($H-K$) versus ($J-H$) CC diagram for the cluster associated with IRAS 06055+2039 is shown in Fig.\\,\\,\\ref{cc.fig}. In this figure we have plotted 53 sources which have good quality photometric magnitudes in all the three $JHK_{s}$ bands (2MASS `read-flag' value of 1 -- 3). Henceforth, for the analysis of 2MASS data, we will be using only good quality photometric data which have the above `read-flag' values. For clarity we have classified the CC diagram into three regions (e.g. Sugitani et al. 2002; Ojha et al. 2004a \\& b). The ``F\" sources are located within the reddening bands of the main sequence and the giant stars. These sources are generally considered to be either field stars, Class III objects or Class II objects with small NIR excess. ``T\" sources populate the region redward of the ``F\" region but blueward of the reddening line corresponding to the red end of the T Tauri locus. These sources are classical T Tauri stars (Class II objects) with large NIR excess or Herbig AeBe stars with small NIR excess. Redward of the ``T\" region is the ``P\" region which are mostly protostar-like Class I objects and Herbig AeBe stars. Majority of sources in our sample are almost equally distributed in the  ``F\" and the ``T\" regions, whereas, only four lie in the ``P\" region. A total of 18 out of 80 (22\\%) sources show infrared excess (i.e. sources populating the ``T\" and the ``P\" regions). However, it is important to note here that this NIR excess fraction is just the lower limit as several cluster members detected in the $K_{s}$ band were not detected in the other two shorter wavelength bands and hence are not part of this sample. The NIR excess in pre-main sequence (PMS) stars is due to the optically thick circumstellar disks/envelopes. These disks/envelopes become optically thin with age hence the fraction of NIR excess stars decreases with age. For very young ($\\le$ 1 Myr) embedded clusters the fraction is $\\sim$ 50 \\% (Lada et al. 2000; Haisch et al. 2000) and it decreases to $\\sim$ 20 \\% for more evolved (2 -- 3 Myr) clusters (Haisch et al. 2001; Texeira et al. 2004). The fraction of stars with NIR excess seen in this cluster suggests an upper limit 2 -- 3 Myr on the age which is reasonably consistent with that suggested from the core radius value. The age estimates for the cluster NGC 2175 associated with the extended HII region Sh 252 are 2 Myr (Grasdalen \\& Carrasco 1975) and 1 -- 2 Myr (K\\\"{o}mpe et al. 1989) which agrees rather well with our estimates for the cluster associated with IRAS 06055+2039 which also belongs to the Sh 252 complex.  We have calculated the extinction by de-reddening the stars in the CC diagram. The stars are shifted to a line drawn tangential to the turn-off point of the main sequence locus (see Fig.\\,\\ref{cc.fig}). The amount of shift gives an estimate of the extinction of individual stars. The extinction values range from $A_{V}$ $\\sim$ 0 to 13 mag with an average foreground extinction of $A_{V}$ $\\sim$ 7 mag. The range of extinction values obtained shows up as the spread of stars along the reddening band in the CC diagram. This indicates that the cluster is partially embedded (Teixeira et al. 2004). Also according to Lada \\& Lada (2003), low extinction values ($A_{V}$ $\\sim$ 1 -- 5 mag) are typical of partially embedded clusters. \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fig7.ps}} \\caption{Colour-colour diagram of the infrared cluster in the IRAS 06055+2039 region. The two solid curves represent the loci of the main sequence (thin line) and the giant stars (thicker line) derived from Bessell \\& Brett (1988). The long-dashed line is the classical T Tauri locus from Meyer et al. (1997). The parallel dotted lines are reddening vectors with the crosses placed along these lines at intervals corresponding to five magnitudes of visual extinction. We have assumed the interstellar reddening law of Rieke \\& Lebofsky (1985) ($A_{J}/A_{V}$ = 0.282; $A_{H}/A_{V}$ = 0.175 and $A_{K}/A_{V}$ = 0.112). The short-dashed line represents the locus of the Herbig AeBe stars (Lada \\& Adams 1992). The plot is classified into three regions namely ``F\", ``T\" and ``P\" (see text for details). The colours and the curves shown in the figure are all transformed to the Bessell \\& Brett (1988) system. The solid line shown is drawn tangential to the turn-off point of the main sequence locus. The arrow points to the position corresponding to the central brightest star in the cluster. The mean photometric errors are shown in the upper right corner.} \\label{cc.fig} \\end{figure}  Figure\\,\\,\\ref{cmd.fig} shows the ($H-K$) versus $K$ colour-magnitude (CM) diagram for 79 sources with good quality $HK$ magnitudes. Using the zero age main sequence (ZAMS) loci and the reddening vectors, we estimate the spectral type of the brightest star in the cluster to be $\\sim$ B0.5. This is the central IRAS point source and from the NIR estimates seems to be the most massive star in the cluster. A similar estimate for the spectral type is also obtained from the analysis of the ($J-H$) versus $J$ CM diagram.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fig8.ps}} \\caption{Colour-magnitude diagram of the infrared cluster in the IRAS 06055+2039 region. The nearly vertical solid lines represent the zero age main sequence (ZAMS) loci with 0, 20 and 40 magnitudes of visual extinction corrected for the distance. The slanting lines show the reddening vectors for each spectral type. The magnitudes and the ZAMS loci are all plotted in the Bessell \\& Brett (1988) system. The arrow points to the position corresponding to the central brightest star in the cluster. The mean photometric errors are shown in the upper right corner.} \\label{cmd.fig} \\end{figure}  \\subsubsection{$K_{s}$ - Band Luminosity Function} We use the 2MASS $K_{s}$ - band star counts to derive the $K_{s}$ - band luminosity function (KLF) for the embedded cluster. In order to obtain the KLF of the cluster it is essential to account for the background and foreground source contamination. For this purpose we use both the Besan\\c{c}on Galactic model of stellar population synthesis (Robin et al. 2003) and the observed control field star counts. We use the same control field as described in Sect\\,\\,\\ref{ssnd.sect}.  Star counts were predicted using the Besan\\c{c}on model in the direction of the control field. We have checked the validity of the simulated model by comparing the model KLF with that of the control field and found both the KLFs to match rather well. As mentioned in the previous section, the average foreground extinction is determined to be A$_{V}$$\\sim$ 7 mag. Hence, assuming spherically symmetric geometry, the background population is then seen through a cloud with extinction upto A$_{V}$$\\sim$ 14 mag (7 $\\times$ 2). Model simulations with A$_{V}$ = 0 mag and d $<$ 2.6 kpc gives the foreground contamination. The background population is generated with A$_{V}$ = 14 mag and d $>$ 2.6 kpc. We determine the fraction of the contaminating stars (foreground + background) over the total model counts. This fraction is used to scale the observed control field and subsequently the star counts of the modified control field are subtracted from the KLF of the cluster to obtain the final corrected KLF which is shown in the left panel of Fig.\\,\\,\\ref{klf_klf_slope.fig}.  \\begin{figure*} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{fig9l.ps} \\includegraphics{fig9r.ps}} \\caption{Left: The corrected $K_{s}$ - band luminosity function (KLF) for the cluster around IRAS 06055+2039 is shown (solid line). The dotted line is the luminosity function without foreground/background correction. Right: The KLF shown as logN versus the $\\rm K_{s}$ magnitude. The straight line is the least squares fit to the data points in the magnitude range 11.5 -- 14.5.} \\label{klf_klf_slope.fig} \\end{figure*}  The right panel of Fig.\\,\\,\\ref{klf_klf_slope.fig} shows the KLF plotted as $\\rm log N$ versus the $K$ magnitude. The KLF can be written as a power-law \\begin{equation} \\frac{dN(K_{s})}{dK_{s}} \\propto 10^{\\alpha K_{s}} \\end{equation} where, the left side of the equation denotes the number of stars per unit magnitude bin and $\\alpha$ is the slope of the power-law. A linear least squares fitting algorithm is used to fit the above power-law to the KLF in the magnitude range $K_{s} = 11.5 \\,\\rm to \\,14.5$. We obtain a value of $\\alpha$ = 0.43$\\pm$0.09 for the cluster. The least square fitting is done taking into account the statistical errors on individual data points.  Within the quoted errors the estimate of the power-law slope is consistent with the average value of slopes ($\\alpha \\sim 0.4$) obtained for other young clusters (Lada et al. 1991; Lada \\& Lada 1995; Lada \\& Lada 2003). The power-law slope values obtained for other embedded clusters like W3 main and NGC 7538 are significantly lower ($\\alpha$ $\\sim$ 0.17 to 0.33 -- Ojha et al. 2004a \\& b; Balog et al. 2004). However, it should be noted here that these clusters are much younger ($\\la$ 1 Myr) and the surveys are deeper ($K_{s}\\le 17.5$) which probe the low mass stellar population down to $\\sim \\rm 0.1M_{\\odot}$.  We estimate the masses of the sources in the cluster by comparing them with the evolutionary models of Palla \\& Stahler (1999). Figure\\,\\,\\ref{mass_spec.fig} shows the ($J-H$) versus $J$ CM diagram for the cluster field. We use the $J$-band magnitudes rather than $K_{s}$ because it is less affected by emission from circumstellar material. The solid curve represents the ZAMS isochrone for a 2 Myr cluster from Palla \\& Stahler (1999). Majority of sources have a typical mass of $\\sim \\rm 2 M_{\\odot}$ which we assume to be representative of the stellar population in the cluster. For a log-normal IMF, the power law slope ($\\gamma$), which is 1.35 for a Salpeter IMF, is variable and is given by $\\rm \\gamma = 0.94 + 0.94\\,log(m_{\\star})$, where $m_{\\star}$ is the stellar mass (Miller \\& Scalo 1979; Lada et al. 1993). Using this relation, we estimate $\\gamma \\sim 1.2$ for this mean mass. The power-law slope for the mass to luminosity relation is $\\beta$ $\\approx 1$ for clusters of age $\\sim 10^{6}$ Myr (Simon et al. 1992; Lada et al. 1993). The corresponding slope of the KLF, $\\alpha$ ($=\\gamma/2.5\\beta$), is 0.48 which is consistent with the value obtained from the least squares fit to the KLF. As is seen from Fig.\\,\\ref{mass_spec.fig}, the mass of the majority of sources in the cluster are below 2.5 $\\rm M_{\\odot}$ and the lowest mass limit is $\\sim$ 0.4 $\\rm M_{\\odot}$ from our sample.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fig10.ps}} \\caption{The mass spectrum for the cluster associated with IRAS 06055+2039. The solid curve represents the model isochrone from Palla \\& Stahler (1999) for a 2 Myr PMS stellar population for the mass range 0.1 -- 3 $\\rm M_{\\odot}$. The slanted dotted lines are the reddening vectors for 2.5 and 0.2 $\\rm M_{\\odot}$ PMS stars respectively. Sources which lie above the reddening vector for 2.5 $\\rm M_{\\odot}$ are luminous massive stars. The arrow points to the position corresponding to the central massive star in the cluster.} \\label{mass_spec.fig} \\end{figure}  \\subsection{Spatial Distribution of UIBs from MSX Data} \\label{uib.sect} We have used the scheme developed by Ghosh \\& Ojha (2002) to extract the contribution of UIBs (due to the Polycyclic Aromatic Hydrocarbons (PAHs)) from the mid-infrared images in the four MSX bands. The emission from each pixel is assumed to be a combination of two components. The first is the thermal continuum from dust grains (gray body) and the second is the emission from the UIB features falling within the MSX bands. The scheme assumes the dust emissivity to follow the power law of the form $\\epsilon_{\\lambda} \\propto \\lambda^{-1}$ and the total radiance due to the UIBs in band C to be proportional to that in band A. A self consistent non-linear chi-square minimization technique is used to estimate the total emission from the UIBs, the temperature and the optical depth. The spatial distribution of emission in the UIBs with an angular resolution of $\\rm \\sim 18\\arcsec$ (for the MSX survey) extracted for the region around IRAS 06055+2039 is shown in Fig.\\,\\,\\ref{pah.fig}. \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fig11.ps}} \\caption{Spatial distribution of the total radiance in the UIBs for the region around IRAS 06055+2039 as extracted from the MSX four band images. The contour levels are at 5, 10, 20, 30, 40, 50, 60, 70, 80, 90 and 95 \\% of the peak emission of $\\rm 4.1\\times10^{-5}$ W\\,m$^{-2}$\\,Sr$^{-1}$. The peak position of the UIB distribution (filled triangle), the 1280 MHz radio emission (open star) and the IRAS point source position (plus) are also indicated.} \\label{pah.fig} \\end{figure}  Comparing the morphology with the radio continuum maps, we see that the emission from the UIBs is much more extended though the gross morphologies are similar with a relatively steep intensity gradient towards the N-E and a smoothly decreasing intensity distribution to the S-W. The peak position of the UIB distribution matches rather well with the radio peak.  The integrated emission from the region around IRAS 06055+2039 in the UIB features within the band A of MSX (viz., 6.2, 7.7 \\& 8.6\\,$\\mu$m) is found to be 2.85 $\\times 10^{-12}$ Wm$^{-2}$ (see Fig.\\,\\,\\ref{pah.fig}). For comparison, we have estimated the emission in individual UIB features from the IRAS LRS spectrum (covering 8 -- 22\\,$\\mu$m), for this source. The total emission in the 7.7 and 8.6 $\\mu$m features is $\\sim$ 7.33 $\\times 10^{-12}$ Wm$^{-2}$, in the 11.3\\,$\\mu$m feature is $\\sim$ 1.32 $\\times 10^{-12}$ Wm$^{-2}$ and an upper limit for the 12.7\\,$\\mu$m feature is 5.9 $\\times 10^{-13}$ Wm$^{-2}$. The last value is an upper limit due to possible contamination from the Ne [II] line at 12.8\\,$\\mu$m. Hence, the UIB emission extracted from the MSX band A is $\\la$ 39\\% of the estimate from LRS. This is reasonable considering the larger effective field of view for the latter.  \\subsection{Emission from Ionized Gas} The radio continuum emission from the ionized gas associated with the region around IRAS 06055+2039 at 1280 and 610 MHz is shown in Fig.\\,\\,\\ref{1280_610.fig}. The integrated flux densities from these maps are 183 and 282 mJy at 1280 and 610 MHz respectively. It should be noted here that the flux densities are obtained by integrating upto 3\\,$\\sigma$ level, where $\\sigma$ is the rms noise of the maps. The integrated flux densities obtained from the GMRT maps are consistent with the results of Felli et al. (1977) and White \\& Gee (1986). Felli et al. (1977) get values of 230 and 205 mJy at 1415 and 4995 MHz respectively. White \\& Gee (1986) estimate the total integrated flux density at 5 GHz to be 198 mJy. The contour maps display a cometary morphology with a bright arc-shaped edge on the N-E side and a smoothly decreasing intensity distribution on the opposite side. The cometary morphology is more clearly seen in the 610 MHz map. Such a morphology implies that the HII region is ionization bounded towards the N-E and density bounded towards the S-W. The position of the brightest radio peak matches with the central bright star seen in the infrared images. Similar morphology is seen from the 4995 MHz map of Felli et al. (1977). They report the presence of several dense clumps. According to them the radio continuum emission from this HII region (S252 A) is spatially separate from the larger extended HII region associated with Sh 252. CO line observations of Lada \\& Wooden (1979) show that this compact HII region (S252 A), the associated H$_{2}$O maser and CO bright spot are nearly coincident and located near the interface of the molecular cloud with the extended S252 HII region. The detection of the H$_{2}$O maser also implies the relative youth of the region. Lada \\& Wooden (1979) also suggest S252 A to be in the early stages of stellar evolution in which a massive star (or stars) has reached the main sequence and created a compact HII region within its parental molecular cloud. VLA maps at 15 and 5 GHz of White \\& Gee (1986) also trace similar cometary morphologies.  Using the low frequency flux densities at 1280 and 610 MHz from our GMRT observations and 5 GHz data from White \\& Gee (1986), we derive the physical properties of the compact core of the HII region associated with IRAS 06055+2039. Mezger \\& Henderson (1967) have shown that for a homogeneous and spherically symmetric core, the flux density can be written as \\begin{equation} S = 3.07 \\times 10^{-2} T_{e} \\nu^{2} \\Omega (1 - e^{-\\tau(\\nu)}) \\end{equation} where, \\begin{equation} \\tau(\\nu) = 1.643a \\times 10^{5} \\nu^{-2.1} (EM) T_{e}^{-1.35} \\end{equation} where, $S$ is the integrated flux density in Jy, $T_{e}$ is the electron temperature in K, $\\nu$ is the frequency of observation in MHz, $\\tau$ is the optical depth, $\\Omega$ is the solid angle subtended by the source in steradians and $EM$ is the emission measure in $\\rm cm^{-6} pc$. $a$ is a correction factor and we use a value of 0.99 (using Table 6 of Mezger \\& Henderson 1967) for the frequency range 0.6 -- 5 GHz and $T_{e} = \\rm 8000 K$. The two GMRT maps are convolved to a common angular resolution of $\\rm 12\\arcsec \\times 12\\arcsec$ which is the resolution of the 5 GHz map of White \\& Gee (1986). In our case since the core is unresolved, $\\Omega$ is taken as this synthesized beam size (i.e $\\Omega = 1.133 \\times \\theta_{a} \\times \\theta_{b}$, where $\\theta_{a}$ and $\\theta_{b}$ are the half power beam sizes). The peak flux densities of the core in the 0.6 -- 5 GHz frequency range appear to lie in the optically thin region. Using these peak flux densities we derive the emission measure for an estimated electron temperature.  The electron temperature ($ T_{e}$) of HII regions is known to increase linearly with Galacto-centric distance ($ D_{G}$) (Deharveng et al 2000 and references therein). This is due to the decrease in heavy element abundance with $ D_{G}$ which results in higher $ T_{e}$. The values of $ T_{e}$ derived by Omar et al. (2002) for a sample of three Galactic HII regions are also consistent with the relationship given in Deharveng et al. (2000). Assuming $ D_{G}$ as 10 kpc (Shirley et al. 2003) for IRAS 06055+2039 (S252A), we obtain a value of $\\sim$ 8000 K for the electron temperature. Using this value of the electron temperature, the peak flux densities were used to fit the above equations (Eqns. 3 \\& 4). The best fit value for the emission measure is $\\rm 8.8 \\pm 0.4 \\times 10^{4} cm^{-6} pc$. We get an estimate of $\\rm 1.05 \\times 10^{3} cm^{-3}$ for the electron density ($n_{e}$ $= (EM/r)^{0.5}$, $r$ being the core size which in this case corresponds to the synthesized beam size). These values agree reasonably well with the estimates of Felli et al. (1977) for the brightest peak (A3) of the component S252 A (see Fig. 5 of Felli et al. 1977). They derive a value of $\\rm 9.9 \\times 10^{4} cm^{-6} pc$ and $\\rm 5.95 \\times 10^{2} cm^{-3}$ for $EM$ and $n_{e}$ respectively. They assumed an electron temperature of $\\rm 10^{4} K$ and a distance of 2 kpc. It should also be noted here that the 5 GHz map of Felli et al. (1977) has a larger beam size ($\\rm \\sim 8\\arcsec\\times21\\arcsec$).   Taking the total integrated flux density of 183 mJy at 1280 MHz and using the formulation of Schraml \\& Mezger (1969) and the table from Panagia (1973; Table II), we estimate the exciting star of this HII region to be of the spectral type B0 -- B0.5. This is consistent with the spectral class obtained by Felli et al. (1977) and White \\& Gee (1986). The FIR flux densities from the IRAS PSC yield a luminosity of $\\sim$ 10$^{4}\\rm L_{\\odot}$ which implies an exciting star of spectral type B0.5, in good agreement with radio measurements.  In addition to the diffuse emission seen in our 1280 MHz map, we also detect a few discrete sources probably representing high density clumps which are listed in Table\\,\\,\\ref{radio_sources.tab}. We designate them as R1, R2,....R7 and their positions are marked in Fig.\\,\\ref{1280_610.fig}. Three such dense clumps were also detected in the 5 GHz map of Felli et al. (1977). The position of R1 is spatially coincident with the position of the central bright and massive ($\\sim$ B0.5) star seen in the infrared cluster and is possibly the exciting source of the HII region. The other dense clumps could also be possible discrete radio sources but the resolution of our map makes it difficult to comment on their nature. \\begin{table} \\caption{Discrete sources extracted from the 1280 MHz map of the region associated with IRAS 06055+2039.} \\label{radio_sources.tab} \\begin{tabular}{c|c|c|c} \\hline\\hline Source     & RA (2000.0)   & DEC (2000.0)   & Peak Flux density \\\\ & (h m s)       & (d m s)        & (mJy/beam)  \\\\ \\hline R1             & 06 08 32.16    & +20 39 18.7    & 5.7       \\\\ R2             & 06 08 31.95    & +20 39 23.1    & 5.4 \t     \\\\ R3             & 06 08 32.79    & +20 39 16.1    & 4.3       \\\\ R4             & 06 08 31.04    & +20 39 24.9    & 3.3       \\\\ R5             & 06 08 30.92    & +20 39 30.3    & 2.8       \\\\ R6             & 06 08 31.39    & +20 39 02.1    & 3.2       \\\\ R7             & 06 08 32.20    & +20 39 01.7    & 2.8       \\\\ \\hline \\end{tabular} \\end{table} These seven clumps contribute $\\sim$ 5\\% of the total integrated emission from the ionized region around IRAS 06055+2039, the remaining being of diffuse nature.  \\subsection{Emission from Shocked Neutral Gas} The rotational-vibrational line of molecular hydrogen (H$_{2}$ (1-0)S1, 2.12\\,$\\mu$m) traces the shocked neutral gas at the interface between the ionized and the molecular gas. In the photo-dissociation regions (PDRs), the molecular emission traces the first neutral layer beyond the ionization front. In Fig.\\,\\,\\ref{h2_brg_rad.fig}, we compare the morphologies of the 610 MHz continuum map with the continuum subtracted narrow-band H$_{2}$ (left) and Br$\\gamma$ (right) images. It is interesting to note that the shocked molecular hydrogen envelopes the radio emitting region. The Br$\\gamma$ image shows the presence of a faint diffuse emission which correlates well with the cometary morphology of the radio continuum emission. The H$_{2}$ arc which traces the ionization front lies beyond the Br$\\gamma$ emission. Comparison with the 1280 MHz continuum map also shows similar morphology. From the position of the H$_{2}$ arc, we estimate the radius of the HII region to be $\\sim$ 0.4 pc. \\begin{figure*} \\resizebox{\\hsize}{!}{\\includegraphics{fig12l.ps} \\includegraphics{fig12r.ps}} \\caption{610 MHz radio contours overlaid over the continuum subtracted H$_{2}$ (left) and Br$\\gamma$ (right) images for the region associated with IRAS 06055+2039. The contour levels are same as in Fig.\\,\\,\\ref{1280_610.fig}. The plus sign in both the images marks the position of the IRAS point source.} \\label{h2_brg_rad.fig} \\end{figure*}  \\subsection{Emission from Dust: Temperature, Optical Depth and Dust Mass} As discussed in Sect.\\,\\,\\ref{uib.sect}, the MSX images were used to obtain the spatial distribution of temperature and optical depth ($\\tau_{10}$) of warm dust with the assumption that the dust is optically thin and the dust emissivity follows the power law of the form $\\epsilon_{\\lambda} \\propto \\lambda^{-1}$ (Mathis et al. 1983; Scoville \\& Kwan 1976). We obtain peak values of $1.4\\times10^{-4}$ and 155 K for $\\tau_{10}$ and the dust temperature respectively. Figure\\,\\,\\ref{msx_tautemp.fig} shows the optical depth and the mid-infrared dust temperature maps. \\begin{figure*} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{fig13l.ps}\\includegraphics{fig13r.ps}} \\caption{Left: The spatial distribution of dust optical depth $\\tau_{10}$ in the region around IRAS 06055+2029. The contour levels are at 5, 10, 20, 30, 40, 50, 60, 70, 80 and 90 \\% of the peak value of $\\rm 1.4 \\times 10^{-4}$. Right: The spatial distribution of the mid-infrared dust temperature in this region. The contour levels are for the 100,115,125,135,140,150 and 154 K. The peak temperature is 155 K. The plus sign in both the images marks the position of the IRAS point source.} \\label{msx_tautemp.fig} \\end{figure*} Morphologically, the spatial distribution of the UIBs and the optical depth contours are similar with the intensity peaks matching rather well. This indicates the presence of higher dust densities near the embedded cluster. The temperature distribution shows lower values near the centre with a plateau like feature running diagonally along the S-E and N-W direction. Higher temperatures are towards the periphery with the peak seen towards the north. The optical depth and temperature are inherently anti-correlated. Hence, we see that the region has higher optical depth at the centre with the values decreasing outwards. Several peaks seen in the $\\tau_{10}$ map could possibly indicate the clumpy nature of the region. As will be evident later, we see similar trends for the $T\\rm(12/25)$ and $\\tau_{25}$ maps derived from the IRAS-HIRES images.  The MSX Point Source Catalog (MSX PSC) lists two mid-infrared sources which fall within the radio nebulosity of IRAS 06055+2039. We designate them as M1 and M2. The MSX PSC flux densities for these two sources are listed in Table\\,\\,\\ref{msx_psc.tab}. \\begin{table} \\caption{Flux densities for the MSX point sources possibly associated with IRAS 06055+2039} \\label{msx_psc.tab} \\begin{tabular}{c|c c} \\hline\\hline & \\multicolumn{2}{|c}{Flux density (Jy)}\\\\ \\hline Wavelength ($\\mu$m) & M1$^{1}$        & M2$^{2}$\t\\\\ \\hline 8.3                 & 3.52            & 0.97               \\\\ 12.1                & 6.24            & 1.40                \\\\ 14.7                & 2.42            & 4.19                \\\\ 21.3                & 6.66            & 13.98              \\\\ \\hline \\end{tabular}  {\\tiny $^{1}$G189.7672+00.3407 ($\\alpha_{2000}= 06^{h}08^{m}33.02^{s}$; $\\delta_{2000} = 20^{\\circ}39\\arcmin32.04\\arcsec$)\\\\ $^{2}$G189.7677+00.3376 ($\\alpha_{2000}= 06^{h}08^{m}32.40^{s}$; $\\delta_{2000} = 20^{\\circ}39\\arcmin25.20\\arcsec$)\\\\} \\end{table} Comparing the MSX mid-infrared colours $\\rm F_{21/8}$, $\\rm F_{14/12}$, $\\rm F_{14/8}$ and $\\rm F_{21/14}$ of these two sources with study of the Galactic plane population by Lumsden et al. (2002), we infer that M1 is possibly a Herbig AeBe or a foreground star, whereas, M2 falls in the zone occupied mostly by compact HII regions. The NIR counterpart from 2MASS catalog for M2 is the central bright IRAS point source. The mid- and near-infrared colours $\\rm F_{21/8}$, $\\rm F_{8/K}$, $\\rm F_{21/12}$ and $\\rm F_{K/J}$ of this source are also consistent with compact HII regions.  The IRAS-HIRES maps (at 12, 25, 60 and 100\\,$\\mu$m) were also used to obtain the spatial distribution of warm and cold dust colour temperatures ($T\\rm (12/25)$, $T\\rm (60/100)$) and optical depths ($\\tau_{25}$, $\\tau_{100}$). We assume the dust emissivity to follow the power law of the form $\\epsilon_{\\lambda} \\propto \\lambda^{-1}$. Figure\\,\\,\\ref{hires_1234.fig} shows the dust temperature and optical depth maps. The maps for the optical depth $\\tau_{25}$ and colour temperature $T\\rm (12/25)$ represent the warmer dust component. The distribution is centrally dense with the optical depth peak and hence lower derived temperature at the centre. This is similar to the distribution seen in Fig.\\,\\,\\ref{msx_tautemp.fig}. On the other hand $\\tau_{100}$ and $T\\rm (60/100)$ distributions are from a relatively colder component which probably forms an envelope around the warmer dust. Unlike the warmer dust temperature distributions (obtained from mid-infrared emission), the $T\\rm (60/100)$ distribution has its peak at the centre with the temperature decreasing towards the periphery. \\begin{figure*} \\centering \\resizebox{\\hsize}{!}{\\includegraphics{fig14ul.ps} \\hskip -2 cm\\includegraphics{fig14ur.ps}} \\resizebox{\\hsize}{!}{\\includegraphics{fig14ll.ps} \\hskip -2 cm\\includegraphics{fig14lr.ps}} \\caption{Upper panel -- The dust optical depth ($\\tau_{25}$) (left) and colour temperature ($T\\rm(12/25)$) (right) maps of the region around IRAS 06055+2039. The optical depth contours are 5, 10, 20, 30, 40, 50, 60, 70, 80, 90 \\% of the peak value of $2.2\\times10^{-5}$. The dust temperature contour levels are at 140, 145, 150, 155, 160, 170 and 180 K from the centre to the periphery. The peak value is 189 K. Lower panel -- The dust optical depth ($\\tau_{100}$) (left) and colour temperature ($T\\rm(60/100)$) (right) maps of the region around IRAS 06055+2039. The optical depth contours are 10, 20, 30, 40, 50, 60, 70, 80, 90 \\% of the peak value of $1.2\\times10^{-2}$ from the centre to the periphery. The dust temperature contour levels are at 20, 30, 33, 39, 46, 53, 59 and 62 K. The peak value is 66 K. The plus sign in all the images marks the position of the IRAS point source.} \\label{hires_1234.fig} \\end{figure*} For $\\tau_{25}$ and $\\tau_{100}$, we obtain peak values of $2.2\\times10^{-5}$ and $1.2\\times10^{-2}$ respectively. The dust temperature distributions for $T\\rm(12/25)$ and $T\\rm(60/100)$ peak at 189 and 66 K respectively. Schreyer et al. (1996) derive a value of 31.2 K for $T\\rm(60/100)$ and $1.21\\times10^{-3}$ for $\\tau_{100}$ from the IRAS PSC flux densities. This difference is due to the different resolution of the raw IRAS and the HIRES processed maps. Scaling the $\\tau_{25}$ peak value, we obtain a value of $\\tau_{10}$ $\\sim$ $5.6\\times10^{-5}$. The difference between this value and that obtained from the MSX data could be a result of different beam sizes and/or an inhomogeneous medium. Using the peak value of $\\tau_{100}$ distribution, derived from the IRAS-HIRES maps, we estimate the warm dust mass to be $\\sim$ 6 $\\rm M_{\\odot}$.  We use the emission at submillimetre wavebands to study the cold dust environment in the region around IRAS 06055+2039. The spatial distribution of the sub-mm emission is shown in Fig.\\,\\,\\ref{jcmt.fig}. The angular resolutions are $\\rm 7\\arcsec.8$ and $\\rm 15\\arcsec.2$ for the 450 and 850\\,$\\mu$m wave bands respectively. The main source (central dense core which covers the region upto $\\sim$ 25\\% of the peak intensity) seems elongated in both the maps. Apart from this dense core, several dust clumps are seen which could have probably formed due to the fragmentation of the original cloud.  The dust mass can be estimated from the following relation: \\begin{equation} M_{dust} = 1.88 \\times 10^{-4} \\left(\\frac{1200}{\\nu}\\right)^{3+\\beta} S_{\\nu}(e^{0.048\\nu/T_{d}} - 1) d^{2} \\end{equation} This is taken from Sandell (2000) and is a simplified version of Eqn. 6 of Hildebrand (1983). The above equation assumes the standard Hildebrand opacities (i.e. $\\rm \\kappa_{1200GHz} = 0.1 cm^{2}g^{-1}$). Here, $S_{\\nu}$ is the flux density at frequency $\\nu$, $T_{d}$ is the dust temperature which we assume to be 20 K (Klein et al. 2005; Mueller et al. 2002), $\\beta$ is the dust emissivity index and is taken to be 2 (Hildebrand 1983) and $d$ is the distance to the source in kpc. The flux densities are obtained from the JCMT-SCUBA maps shown in Fig.\\,\\,\\ref{jcmt.fig}. To obtain the flux density of the entire cloud, we have integrated upto the last contour (which is at 5\\% of the peak value). Using the above relation, we estimate dust masses of $\\sim$ 70 and 90 $\\rm M_{\\odot}$ from the 450 and 850\\,$\\mu$m maps respectively. Assuming a gas-to-dust ratio of 100, the above values translate to total masses of 7000 and 9000 $\\rm M_{\\odot}$ for the cloud from the 450 and 850\\,$\\mu$m maps respectively. We also estimate the total mass of only the central dense core to be $\\sim$ 875 and 1250 $\\rm M_{\\odot}$ from 450 and 850\\,$\\mu$m maps respectively. This source (S252A) has also been studied by Mueller et al. (2002) at 350\\,$\\mu$m and more recently by Klein et al. (2005) at 850 and 1300\\,$\\mu$m. Scaling to the distance and gas-to-dust ratio assumed by us, the corresponding mass from Mueller et al. (2002) is $\\sim$ 600 $\\rm M_{\\odot}$ and from Klein et al. (2005) is $\\sim$ 100 $\\rm M_{\\odot}$. Comparing the mass derived from the 850\\,$\\mu$m map, our estimate is very close to the mass obtained by Mueller et al. (2002) considering the fact that they have assumed a higher dust temperature (29 K). The mass estimate from the 450\\,$\\mu$m maps is $\\sim$ 35 \\% lower which could have been affected by the large atmospheric extinction correction applied to the data. The mass derived by Klein et al. (2005) are relatively lower. This could be possibly due to the fact that the flux density values presented by Klein et al. (2005) are lower compared to our values. From the CS line maps, Zinchenko et al. (1998) derive a cloud mass of 3132 $\\rm M_{\\odot}$. It should be noted here that the CS maps are from a larger region. Furthermore, a point worth discussing here is the effect of varying $T_{d}$ and $\\beta$. Exploring the range of $T_{d}$ (20 -- 40 K) and $\\beta$ (1 -- 2), we infer that the mass estimates can vary by upto a factor of $\\sim$ 8.  \\subsection{Comparison of the Different Components Associated with IRAS 06055+2039} In Fig.\\,\\,\\ref{all.fig}, we present the various components of the region associated with IRAS 06055+2039. The plot displays the contour maps of the ionized gas and the emission from dust overlaid on the 2MASS $K_{s}$-band image. We show the contour plots of 610 MHz radio emission, 850\\,$\\mu$m cold dust emission and mark the peak positions for the 100 and 12\\,$\\mu$m emission from warm dust. The 850\\,$\\mu$m emission core lies to the S-E of the ionized region with the warmer dust in between. The 12 and 100\\,$\\mu$m peaks lie relatively closer to the radio peak. The ionized region is seen to be close to the edge of the molecular cloud. Comparing with Fig.\\,\\,\\ref{h2_brg_rad.fig}, we note that the shocked molecular gas lies in between the ionized region and the dense molecular core. The central region of the infrared cluster is located within the HII region which is at the edge of the molecular cloud. \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{fig15.ps}} \\caption{Contour maps of the emission from the ionized gas at 610 MHz (thick line) and dust emission at 850\\,$\\mu$m (thin line) are overlaid on the 2MASS $K_{s}$ band image for the region around IRAS 06055+2039. The dashed arc represents the position of shocked H$_{2}$. The peak positions of the IRAS-HIRES 12\\,$\\mu$m (open triangle) and 100\\,$\\mu$m (open star) maps are indicated. The cross marks the position of the methanol and water masers.} \\label{all.fig} \\end{figure}  The ionized region and the dense molecular core as seen in the radio and the sub-mm maps respectively are possibly at different stages of evolution. The dense molecular core seen in our sub-mm maps does not show any radio emission (down to level of the rms noise which is 0.4 mJy/beam). The sub-mm peak is spatially offset from the FIR peaks and there are no MSX or NIR counterparts seen. This indicates a very early evolutionary stage for this dense and massive molecular core. It is most probably a very early protocluster candidate and we are sampling the initial collapse phase of the star forming core before the formation of the UCHII region (Williams et al. 2004 and references therein). This fact is further supported by the positions of the water and methanol masers which are seen to be coincident with the peak of the molecular core. The central peak of the CS map (Zinchenko et al. 1998) and the peaks of the 450 and 850\\,$\\mu$m JCMT-SCUBA maps match with the position of the masers. More specifically, it is known that methanol maser sites are generally radio quiet (as is the case here) and trace high mass star forming protoclusters in very early evolutionary phases (Minier et al. 2005). On the other hand, the ionized region which is associated with IRAS 06055+2039, has FIR emission, free-free radio emission and has NIR and MSX counterparts. This could probably indicate that the source is at a later evolutionary stage in between an evolved protocluster and a young cluster. Here, massive stars have started forming and a detectable HII region has been created. The cluster is partially embedded in the parental cloud. The sub-mm emission is weak here. A small subcluster is seen close to the edge of the sub-mm core spatially coincident with the secondary peak mentioned in Sect.\\,\\,\\ref{ssnd.sect}. Thus, from this multiwavelength study of the region associated with IRAS 06055+2039, we see the signatures of different evolutionary stages at different locations."
    },
    "0601/astro-ph0601425_arXiv.txt": {
        "abstract": "{The blue Main Sequence (bMS) of $\\omega$ Cen implies a ratio of helium to metal enrichment $\\Delta Y/\\Delta Z \\approx 70$, which is a major enigma. We show that  rotating models of low metallicity stars, which  account for the anomalous abundance ratios of extremely metal poor stars, are also useful for understanding the very high $\\Delta Y/\\Delta Z$ ratio in $\\omega$ Cen. Models of massive stars with moderate initial rotation velocities produce stellar winds with large He-- and N--excesses, but without  the large  C-- (and O--) excesses made by very fast rotation, in agreement with the observed chemical abundance ratios in $\\omega$ Cen.  It is still uncertain whether the abundance peculiarities of $\\omega$ Cen result from the fact that  the high velocity contributions of supernovae escaped the globular cluster, usually considered as  a tidally  stripped  core of a dwarf galaxy. Another  possibility is a general dominance of wind ejecta at very low $Z$, due to the formation of black holes. Some abundance and isotopic ratios like $Mg/Al$, $Na/Mg$, $Ne/N$,  $^{12}C/^{13}C$, $^{16}O/^{18}O$ and $^{17}O/^{18}O$ may allow us to further discriminate between these scenarios and between the AGB and massive star contributions.     \\keywords stars:Omega Centauri -- Helium -- stars: evolution  } \\titlerunning{The high Helium Sequence in $\\omega$ Cen} ",
        "introduction": "The globular cluster $\\omega$ Centauri has remarkable properties: it is the most massive globular cluster in the Galaxy and is often interpreted as the remaining core of an ancient dwarf  galaxy (Bekki \\& Freeman \\cite{BekkiF03}), a possibility supported by the  density profile of the cluster (Ideta \\& Makino \\cite{IdetaM04}). Dynamical studies support a formation of $\\omega$  Cen from one of the  small progenitor galaxies of  the Milky Way (e.g. Gnedin et al. \\cite{Gnedin02}).  The stars in $\\omega$ Cen show a wide spread in metallicity (Norris \\& Da Costa \\cite{NorrisDC95}) from $[Fe/H] = -2$ to $-0.5$. Among its many  abundance peculiarities, $\\omega$ Cen shows a large N--excess, an overabundance of s--elements relatively to Fe (Norris \\& Da Costa \\cite{NorrisDC95}) and  an unusually low $[Cu/Fe]$ ratio relatively to other metal poor stars (Cunha et al. \\cite{Cunha02};  McWilliam \\& Smecker--Hane \\cite{McWilliam05}), which is interpreted as a relative lack of contributions from supernovae SNIa (Cunha et al. \\cite{Cunha02}).  The finding of a double sequence in the globular cluster $\\omega$ Cen by Anderson (\\cite{Anderson97}; see also Bedin et al. \\cite{Bedin04}, Gratton \\cite{Gratton05}) and the  further interpretation of the bluer sequence by a strong excess of helium  constitutes a major enigma  for stellar and galactic evolution (Norris \\cite{Norris04}). The interpretation in terms of an He excess is  convincing and supported by stellar models as well by the morphology of the horizontal branch stars (Piotto et al. \\cite{Piotto05}). The great problem is that the bluer sequence with a metallicity $[Fe/H]= -1.2$  or $Z =  2 \\cdot 10^{-3}$ implies an He--content $Y=0.38$ (0.35-0.45), i.e. an He--enrichment $\\Delta Y = 0.14$ (see Norris \\cite{Norris04}). In turn, this demands a relative helium to metal enrichment  $\\Delta Y/\\Delta Z$ of the order of   70 (Piotto et al. \\cite{Piotto05}; Gratton \\cite{Gratton05}). At the opposite, the system of globular clusters has a constant  $Y=0.250$ (Salaris et al. \\cite{Salaris04}), while $[Fe/H]$ varies a lot, which implies a ratio $\\Delta Y/\\Delta X(\\mathrm{Fe})=0$, where $X(\\mathrm{Fe})$ is the iron mass fraction. The value   $\\Delta Y/\\Delta Z=70$ is  enormous and more than one order of magnitude larger than the current value of $\\Delta Y/\\Delta Z= 4-5$ (Pagel et al. \\cite{Pagel92}) obtained from extragalactic  HII regions.  A value of 4 -- 5  is  consistent with the chemical yields from supernovae (Maeder \\cite{Maeder92}) forming  black holes above about 20 -- 25 M$_{\\odot}$.  The subject of the present work is to examine whether this extreme value of $\\Delta Y/\\Delta Z$ can be accounted for by models  of rotating stars at very low metallicity. These models, which have the same physical ingredients as the models successfully used at solar $Z$, well account (Meynet et al. \\cite{MEMZfirst}) for the abundance anomalies observed in extremely metal poor halo stars.  Sect. 2  collects the relevant observational determinations of the chemical abundances in $\\omega$ Cen and mentions the possible interpretations of  the observed abundance peculiarities. In Sect. 3, we show results of  models of rotating stars and in Sect. 4 we compare the results to observations of the bMS in $\\omega$ Cen. Sect. 5 gives the conclusions. ",
        "conclusions": "The wind contributions of low $Z$ massive rotating stars  are able to produce the high $\\Delta Y/\\Delta Z$ observed in the  bMS sequence in  $\\omega$ Cen.  The observations tend to favour an origin of the high helium observed by contributions from massive stars of intermediate rotation velocities.  At this stage, it is not clear whether the dominance of wind contribution is a general feature of low $Z$ stars or whether the tidal evaporation  experienced by $\\omega$ Cen has   enabled it to lose most of the supernovae ejecta, keeping the enrichment from stellar winds.  Some critical abundance and isotopic ratios may offer further signatures of the contributions of AGB and massive stars."
    },
    "0601/astro-ph0601613_arXiv.txt": {
        "abstract": "The subpulse drifting phenomenon in pulsar radio emission is considered within the partially screened inner gap model, in which the sub-Goldreich-Julian thermionic flow of iron ions or electrons coexists with the spark-associated electron-positron plasma flow. We derive a simple formula that relates the thermal X-ray luminosity $L_{\\rm x}$ from the spark-heated polar cap and the \\EB\\ subpulse periodicity $\\hat{P}_3$ (polar cap carousel time). For PSRs B0943+10 and B1133+16, the only two pulsars for which both $\\hat{P}_3$ and $L_{\\rm x}$ are known observationally, this formula holds well. For a few other pulsars, for which only one quantity is measured observationally, we predict the value of the other quantity and propose relevant observations that can confirm or discard the model. Then we further study the detailed physical conditions that allow such partially screened inner gap to form. By means of the condition $T_{\\rm c}/T_{\\rm s}>1$ (where $T_{\\rm c}$ is the critical temperature above which the surface delivers a thermal flow to adequately supply the corotation charge density, and $T_{\\rm s}$ is the actual surface temperature), it is found that a partially-screened gap (PSG) can be formed given that the near surface magnetic fields are very strong and curved. We consider both curvature radiation (CR) and resonant inverse Compton scattering (ICS) to produce seed photons for pair production, and find that the former is the main agency to produce gamma-rays to discharge PSG. ",
        "introduction": "Pulsar radio emission typically occurs in a form of periodic series of narrow bursts of radiation. These burst often have a complex structure, in which higher order periodicities can be found. The individual pulses consist of one, few to several subpulses. In some pulsars the subpulses demonstrate a very systematic drift across the pulse window. If the pulses are folded with the basic pulsar period then the drifting subpulses form amazingly spectacular patterns called drift-bands.  The phenomenon of drifting subpulses is a long standing puzzle in the pulsar research and its solution would likely result in deeper understanding to the nature of pulsar radiation. It is generally believed that this phenomenon is inherently associated with the so-called inner acceleration region above the polar cap, in which the magnetospheres plasma does not corotate with the neutron star surface. The first model based on this idea was proposed by \\citet[][ RS75 henceforth]{rs75}. The predictions of RS75 model were successfully compared with a handful of pulsars known to show this phenomenon at that time. A decade later, \\citet[][ R86 hereafter]{r86} compiled a list of about 40 drifting pulsars, but drifting subpulses were still regarded as some kind of exceptional phenomenon. However, recently \\citet[][ WES06 hereafter]{wes06} presented the results of a systematic, unbiased search for subpulse modulation in 187 pulsars and found that the fraction of pulsars showing drifting subpulse phenomenon is likely to be larger than 55\\%. They identified 102 pulsars with drifting subpulses in their sample, with a large fraction of newly discovered drifters. The authors concluded that the conditions required for the drifting mechanism to work cannot be very different from the emission mechanism of radio pulsars. WES06 then suggest that the subpulse drifting phenomenon is an intrinsic property of the pulsar emission mechanism, although drifting could in some cases be very difficult or even impossible to detect due to insufficient signal-to-noise ratio. It is therefore essential to attempt to unravel the physical conditions that can lead to formation of an inner acceleration region above the polar cap that could lead to development of the subpulse drift phenomenon.  The classical vacuum gap model of RS75, in which spark-associated sub-beams of subpulse emission circulate around the magnetic axis due to $\\mathbf{E}\\times\\mathbf{B}$ drift of spark plasma filaments, provides the most natural and plausible explanation of drifting subpulse phenomenon. However, despite its popularity, it suffers from the so-called binding energy problem. Namely, under canonical conditions the surface charges (ions or electrons) are likely to be directly pulled out of the surface so that a pure vacuum gap is difficult to form. The alternative steady flow polar cap models, the so-called space-charge-limited flow models \\citep{as79,hm98}, cannot give rise to the intermittent ``sparking'' behavior, which seems necessary to explain the subpulse drift phenomenon in radio pulsars. \\citet[][ GM01 hereafter]{gm01} revisited the binding energy problem of RS75 model and argued that the formation of the vacuum gap (VG) is, in principle possible, although it requires a very strong non-dipolar surface magnetic fields, much stronger than a canonical dipolar component inferred from the observed spindown rate. Once the binding is strong enough to prevent the thermionic emission at the full space charge limited flow, the inevitable \\EB\\ drift of plasma filaments will result in the observable subpulse drift phenomenon. It has been known for a long time that in order to allow all radio pulsars to produce electron-positron pairs (the necessary condition for coherent radio emission), the near-surface magnetic fields must include multipole components dominating over global dipole field \\citep{rs75,as79,zhm00}. Growing observational evidence of non-dipolar structure of surface magnetic field \\footnote{It is worth mentioning that although RS75 implicitly assumed non-dipolar surface magnetic fields to treat the $\\gamma$-$B$ pair production processes, in their calculations of many other physical quantities (such as the surface charge density) they still used dipolar form, presumably for the sake of simplicity.} accumulates, and the suggestion that such a sunspot-like fields form during the early proto-neutron star stage has been proposed \\footnote{Another possible mechanism of creating small scale anomalies of surface magnetic fields was proposed by \\citet{grg03}. They argued that due to a Hall drift instability, the poloidal magnetic field structures can be generated from strong subsurface toroidal fields.} \\citep[e.g.][]{ug04}. \\citet[][ GM02 hereafter]{gm02} calculated the non-dipolar VG model for 42 drifting subpulse pulsars tabulated by R86 and argued that VG can be formed in all considered pulsars, provided that the actual surface magnetic field was close to $10^{13}$ G independently of the value of the canonical dipolar magnetic field.  Although the binding energy problem could be, at least in principle, resolved by assuming an appropriately strong surface magnetic field, yet another difficulty of the RS75 model was that it predicted a much too fast \\EB\\ drift rate. Motivated by this issue, \\citet[][ GMG03 hereafter]{gmg03} developed further the idea of the inner acceleration region above the polar cap by including the partial screening by a sub-Goldreich-Julian thermal flow from the surface due to the spark-associated polar cap heating. This idea was first introduced by \\citep[ CR80 henceforth]{cr80}, who argued that even with thermionic ions included in the flow, the condition above the polar cap is close to that of a pure vacuum gap. A similar model was also invoked by \\citet{um95,um96}. GMG03 reanalyzed this model and argued that a thermostatic self-regulation should keep the surface temperature just few percent below the critical ion temperature at which the gap potential drop is completely screened. This results in more than 90 \\% of screening due to thermionic emission. Given a similar gap height, the actual potential drop, and hence the \\EB\\ drift rate, is about 10 \\% of that of the pure vacuum case. This is still above the threshold for the magnetic pair production. Similar to the avalanche pair production cascade introduced by RS75, the discharge of the growing potential drop above the polar cap would also occur in the form of a number of sparks. In this paper we call such an inner accelerator a partially screened gap (PSG henceforth). The latest XMM-Newton observation of the drifting pulsar PSR B0943+10 \\citep[][ ZSP05 hereafter]{zsp05} reveals a possible hot spot with the surface area much smaller than the conventional polar cap. This is consistent with the polar cap heating from such a PSG, which at the same time gives the right \\EB\\ drift rate. This lends strong support to the PSG model.  In this paper we study PSG model in greater detail and explore the physical conditions for the model to work. We also apply this model to a new set of 102 pulsars from WES05 and show that it can work in every case, provided that the surface non-dipolar magnetic field is strong enough, even stronger than $10^{13}$ G suggested by GM02 and GMG03. Our new treatment is a combination of those used in GM02 and GMG03. Our working hypothesis is that drifting subpulses manifest the existence of a thin inner acceleration region, with an acceleration length scale much shorter than the polar cap size. The ultra-high accelerating potential drop discharges via a number of localized $\\mathbf{E}\\times\\mathbf{B}$ drifting sparks, as has been envisioned by RS75 . These sparks produce isolated columns of electron-positron plasma that stream into the magnetosphere to generate radio-beams of the observed subpulse emission due to some kind of plasma instability (see the \\S8 for more discussion). Due to charge depletion with respect to the co-rotational Goldreich-Julian (1969) value, sparks experience an unavoidable \\EB\\ drift with respect to the polar cap surface. As a consequence, the spark-associated sub-beams of radio emission perform a slow circumferential motion that is responsible for the observed subpulse drift. This model, which is often called ``a pulsar carousel model'' \\citep{dr99}, is examined in this paper \\footnote{Other suggestions of subpulse drifting as phenomena occurring outside the inner gap have been made \\citep{kmm91,sa02,w03,gmm05,fkk06}, but the connection between the radio drifting rate and the X-ray properties in those models is not yet clear, and we do not discuss them in the current paper.}. We are particularly interested in the thermal effect associated with the surface bombardment by back-flowing particles produced by sparks. The intrinsic drift rate (manifested by the tertriary drift periodicity) and the polar cap heating rate (manifested by the thermal X-ray luminosity) should be correlated with each other, since they are determined by the same value of the accelerating potential drop. The properties of drifting subpulses are discussed in \\S2 and the properties of charge depleted acceleration region above the polar cap are discussed \\S3. We find a specific relationship between the appropriate observables and conclud that it holds for a number of pulsars for which good quality data is available (especially PSR B0943+10, ZSP05). It turns out that this relationship depends only on the observational quantities and thus it is a powerful tool for testing this theoretical model. Although the number of pulsars that have all necessary data for such testing is small at the moment, the clean prediction holds the promise to ultimately confirm (or discard) the PSG model in the future. In \\S\\S4-7 we analyze this model in a more detailed manner to investigate the microscopic conditions (e.g. near-surface configuration, radiation mechanism, etc) that are needed to form such PSGs. The paper is summarized in \\S8. ",
        "conclusions": "The phenomenon of drifting subpulses has been widely regarded as a powerful tool for understanding the mechanism of coherent pulsar radio emission. RS75 first proposed that drifting subpulses are related to \\EB\\ drifting sparks discharging the high potential drop within the inner acceleration region above the polar cap. The subpulse-associated streams of secondary electron-positron plasma created by sparks were penetrated by much faster primary beam. This system was supposed to undergo a two-stream instability, which should lead to generation of the coherent radiation at radio wave lengths. However, careful calculations of the binding energy (critical for spark ignition) and the growth rate of the two-stream instability have shown that neither the sparking discharge nor the two-stream instability were able to work in a way proposed by RS75. Nonetheless, qualitatively their idea was still considered attractive, at least to the authors of this paper who have been continuing efforts to search for mechanisms that would actually make the RS75 model to work. In this paper we applied the partially screened gap model proposed by GMG03 to PSR B0943+10, a famous drifter for which ZSP06 successfully attempted to measure the thermal X-ray from hot polar cap. We derived a simple relationship between the X-ray luminosity $L_{\\rm x}$ from the polar cap heated by sparks and the tertiary periodicity $\\hat{P}_3$  of the spark-associated subpulse drift observed in radio band. This relationship reflects the fact that both the drifting (radio) and the heating (X-rays) rates are determined by the same value of the electric field in the partially screened gap. As a consequence of this coupling equations~(\\ref{Lx}) and (\\ref{LxE}) are independent of details of the acceleration region. In PSRs B0943$+$10 and B1133+16, the only two pulsars for which both $L_{\\rm x}$ and $\\hat{P}_3$ are measured, the predicted relationship between observational quantities holds very well. We note that $\\hat{P}_3$ is very difficult to measure and its value is known only for four pulsars: B0943$+$10, B1133+16, B0826$-$34 and B0834$+$06. The successful confrontation of the predicted X-ray luminosity with the observations in PSRs B0943$+$10 and B1133+16 encourages further tests of the model with future X-ray observations of other drifting pulsars. This is particularly relevant to PSR B0834$+$06, whose predicted X-ray luminosity is much higher than in PSR B0943+10, while the distance to both pulsars is almost identical. It is worth mentioning that due to relatively poor photon statistics it is still not absolutely clear whether the X-ray radiation associated with polar cap of PSR B0943+10 is thermal or magnetospheric (or both) in origin. However, the case of Geminga pulsar B0633+17 strongly supports thermal origin. Observations of PSR B0834$+$06 with {\\it XMM-Newton} should help to resolve this question ultimately.  In both the steady SCLF model and the pure vacuum gap model, the potential increases with height quadratically at lower altitudes. However the growth rate is significantly different - the latter is faster by a factor of $R/r_{\\rm p}$. It is well known that in the pure vacuum case (RS75), the potential grows so fast that a primary particle could quickly generate pairs with a high multiplicity, and that some backward returning electrons generate more pairs and soon a ``pair avalanche'' occurs and the potential is short out by a spark. In the PSG model we are advocating, the potential drops by a factor of $\\eta$. For essentially all the cases we are discussing, this $\\eta$ value is much larger than the $r_{\\rm p}/R$ value required for the steady SCLF to operate, although it is much less than unity. The steady state condition is not satisfied, and the the gap is more analogous to a vacuum gap, i.e. the pair discharging happening intermittently. In fact, the partially screened potential drop is still above the threshold for the magnetic pair production, which in strong and curved surface magnetic field is a condition necessary and sufficient for the sparking breakdown.  The original RS75 pure VG model predicts much too high a subpulse drift rate and an X-ray luminosity to explain the case of PSR B0943+10 and other similar cases. Other available acceleration models predict too low a luminosity and the explanation of drifting subpulse phenomenon is generally not clear at all (see ZSP05 for more detailed discussion). In summary, the bolometric X-ray luminosity for the space charge limited flow \\citep{as79} pure vacuum gap (RS75) and partially screened gap (GMG03) is $(10^{-4} - 10^{-5})\\dot{E}$ \\citep{hm02}, $(10^{-1} - 10^{-2})\\dot{E}$ (ZSP05), and $\\sim 10^{-3}\\dot{E}$ (this paper), respectively. The latter model also predicts right \\EB\\ plasma drift rate. Thus, combined radio and X-ray data are consistent only with the partially screened VG model, which requires very strong (generally non-dipolar) surface magnetic fields. Observations of the hot-spot thermal radiation almost always indicate bolometric polar cap radius much smaller than the canonical Goldreich-Julian value. Most probably such a significant reduction of the polar cap size is caused by the flux conservation of the non-dipolar surface magnetic fields connecting with the open dipolar magnetic field lines at distances much larger than the neutron star radius.  Our analysis suggests the following pulsar picture: In the strong magnetic fields approaching $10^{14}$~G at the neutron star surface, the binding energy is high enough to prevent a full thermionic flow from the hot polar cap at the corotation limited level. A partially screened vacuum gap develops with the acceleration potential drop exceeding the threshold for the magnetic pair formation. The growth of this potential drop should be naturally limited by a number of isolated electron-positron spark discharges. As a consequence, the polar cap surface is heated by back-flowing particles to temperatures $T_{\\rm s}\\sim 10^6$~K, just below the critical temperature $T_{\\rm c}$ at which the thermionic flow screens the gap completely. The typical radii of curvature of the field lines ${\\cal R}$ is of the order of polar cap radii $r_{\\rm p}\\sim 10^3-10^4$~cm. The only parameter that is thermostatically adjusted in a given pulsar is the shielding parameter $\\eta=10^{-3}(B_{\\rm s}/B_{\\rm q}){\\cal R}^{-0.5}_6 P \\sim 0.001 T_6^{1.43}{\\cal R}^{-0.5}_6 P$, which determines the actual level of charge depletion with respect to the pure vacuum case ($\\eta$=1), and in consequence the polar cap heating rate as well as the spark drifting rate. It is worth to emphasize that $\\eta\\sim 0.1$ for longer period pulsars $P\\sim 1$ s and $\\eta\\sim 0.01$ for shorter period pulsars. Our calculations are consistent with PSR B0943$+$10 and few other drifting pulsars, for which the signatures of X-ray emission from the hot polar cap were detected.  The sparks operating at the polar cap generate streams of secondary electron-positron plasma flowing through the magnetosphere. These streams are likely to generate beams of coherent radio emission that can be observed in the form of subpulses. \\citet{u87} first pointed out that the non-stationarity associated with sparking discharges naturally leads to a two-stream instability as the result of mutual penetration between the slower and the faster plasma components. \\citet{am98} developed this idea further, calculated the growth rates, and demonstrated that the instability is very efficient in generating Langmuir plasma waves at distances of many stellar radii $r_{\\rm ins} \\sim 10^{4-5} h^{CR}=10^{7-8}$ cm, where $ h^{CR}$ is the height of the acceleration region described by equation (21) \\citep{mgp00}. This is exactly where pulsar's radio emission is supposed to originate \\citep[e.g.][]{kg98}. Conversion of these waves into coherent electromagnetic radiation escaping from the pulsar magnetosphere was considered and discussed by \\citet{mgp00} and \\citet{glm04}. These authors demonstrated that the nonlinear evolution of Langmuir oscillations developing in pulsar's magnetosphere leads to formation of charged, relativistic solitons, able to emit coherent curvature radiation at radio wavelengths. The component of this radiation that is polarized perpendicularly to the planes of dipolar magnetic field can escape from the magnetosphere \\citep[see][for observational evidence of such polarization of pulsar radiation]{lea01}. The observed pulsar radiation in this picture is an indirect consequence of sparking discharges within the inner acceleration region just above the polar cap. In light of this paper we can therefore argue that the coherent pulsar radio emission is conditional on the presence of strong non-dipolar surface magnetic fields at the polar cap, with a strength about $10^{13-14}$~G and radius of curvature of the order $10^{4}$~cm.  In the very strong surface magnetic field assumed within the accelerator, processes such as photon splitting \\citep{bh01} and bound pair creation \\citep{um95} may become important. It has been suggested that such processes could potentially delay pair creation and thus increase the height and voltage of the accelerator. For the photon splitting case, the delay is significant only if photons with both polarization modes split - a hypothesis in strong magnetic fields \\citep{bh01}. It could be possible that only one mode split \\citep{u02}. In such a case the gap height and potential of a PSG would not be significantly affected due to the high pair multiplicity in strong, curved magnetic fields. For bound pairs \\citep{um95,um96}, in the very hot environment near the neutron star surface (with temperatures exceeding 1 MK), it is possible that bound pairs could not survive long from photoionization. Following \\citet{bcp92} and \\citet{um96} one can roughly estimate the mean free path for bound pair photo-ionization. It turns out that for temperatures around (2-3) MK this mean free path is of the order of few meters, which is considerably less than the height of CR-driven PSG. In such a case, even if the bound pairs are initially produced, they would not significantly delay the pair formation. One could address this potential potential problem by referring to the detailed case study of PSR B0943+10 analyzed in this paper. This is the only pulsar in which we have full information concerning \\EB\\ drift rate and the polar cap heating rate. In the analysis we have used two methods. The first method is independent on details of accelerating region (such as height, potential drop, etc.) but is based only on the subpulse drift radio-data (eqs.~[\\ref{Lx},\\ref{LxE}]). The second one includes the detailed treatments of model parameters without considering the delaying effect by photon splitting and bound pairs (eqs.~[\\ref{Ts},\\ref{LxE2}]). Both methods give consistent results, as illustrated in Figures (1-4). We therefore conclude that the delaying processes, if any, are not significant in the PSGs given the physical conditions we invoke. Finally, we note that \\citet{um95} presented a model of the inner accelerator that shares some similar features with our model, although it is basically different. In their model the potential drop is self-regulated by partial screening close to the threshold for field ionization of bound pairs, so the surface temperature is maintained at the level necessary for this screening. In our model, the partial screening keeps the surface temperature slightly below the critical temperature at which the thermionic flow is space charge limited \\citet{as79}. In the light of results presented in this paper we claim that bound pairs do not affect the formation of inner accelerators in pulsars."
    },
    "0601/astro-ph0601339_arXiv.txt": {
        "abstract": "% The aim of this paper is to determine the initial rotational velocities required on the ZAMS for single stars to reach the critical velocity, sometimes called the break-up velocity, during the Main-Sequence (MS) phase. Some useful relations between $\\Omega/\\Omega_{\\rm crit}$, $\\upsilon/\\upsilon_{\\rm crit}$ ($\\upsilon$ is the velocity at the equator), the moments of inertia, the angular momenta, the kinetic energy in the rotation and various other basic physical quantities are obtained. ",
        "introduction": "The new results from VLTI on the shape of Achernar \\citep{Do05} show that detailed parameters of rotating stars become accessible to interferometric observations. In order to make valuable comparisons with models of rotating stars, we provide here a few basic data on rotating stellar models.  We also search the initial conditions required for a star to reach the critical limit during its Main-Sequence phase. When the surface velocity of the star reaches the critical velocity ({\\it i.e.} the velocity such that the centrifugal acceleration exactly balances gravity), one expects that equatorial ejections of matters ensue \\citep{Town04}. The physics involved in the ejection process, determining the quantity of mass ejected, the timescales between outburst episodes, the conditions required for a star to present such outbursts are important questions not only for understanding Be, B[e] or Luminous Blue Variable stars, but probably also for having a better knowledge on how the first stellar generations in the Universe behaved \\citep{EkTar}. Indeed, at lower metallicity, the stellar winds are weaker and much less angular momentum is removed at the surface. This favors an evolution to the critical limit and may have important consequences for the evolution of the first stellar generations. Also there are some indications that the distribution of the initial velocities contains more rapid rotators at lower metallicity \\citep{MG99}. Finally, realistic simulations of the formation of the first stars in the Universe show that the problem of the dissipation of the angular momentum is more severe at very low $Z$ than at the solar $Z$. Thus these stars might begin their evolution with a higher amount of angular momentum \\citep{Ab02}.  Rotation affects all the outputs of the stellar models. The reader will find reviews on the effects of rotation in \\citet{MMAA}, \\citet{Hel00}, \\citet{Ta04} and \\citet{MT}. In this work, we compute 112 different stellar models with initial masses equal to 3, 9, 20 and 60 M$_\\odot$, for metallicities $Z$ equal to 0, 0.00001, 0.002 and 0.020, and for values of the ratio of the angular velocity $\\Omega$ to the critical angular velocity $\\Omega_{\\rm crit}$ equal to 0.1, 0.3, 0.5, 0.7, 0.8, 0.9, 0.99. For the 3, 9 and 20 M$_\\odot$ models, the computation was performed until the end of the Main-Sequence phase or until the star reaches the critical velocity. For the 60 M$_\\odot$ stars, we have obtained all the models on the ZAMS, but only a subset of them were computed further on the Main-Sequence. We consider that a star arrives on the ZAMS, when a fraction of 0.003 in mass of hydrogen has been burned at the center. On the ZAMS, the star is supposed to have a solid body rotation. During the Main-Sequence phase the variation of $\\Omega$ inside the star is computed self-consistently taking into account the various transport mechanisms, {\\it i.e.}  convection, shear diffusion, meridional circulation and horizontal turbulence. The removal of angular momentum at the surface by the stellar winds and the changes of $\\Omega$ resulting from movements of contraction or expansion are also accounted for. A detailed description of the physical ingredients used in this grid of models will be given elsewhere (Ekstr\\\"om et al., in preparation), let us just mention here that the prescriptions for the opacity tables, the initial chemical compositions and the treatment of convection are as in \\citet{MMXI}.  \\begin{figure}% \\plottwo{meynet_fig1a.eps}{meynet_fig1b.eps} \\caption{{\\it Left:} Variations of the momentum of inertia on the ZAMS as a function of the initial masses for various metallicities $Z$ (squares: $Z=0.020$; triangles: $Z=0.002$; pentagons: $Z=0.00001$ and circles: $Z=0$). {\\it Right:} Variations of the gravitational energy on the ZAMS as a function of $M_{\\rm ini}$ for various $Z$. } \\label{moin} \\end{figure}  \\begin{figure}% \\plottwo{meynet_fig2a.eps}{meynet_fig2b.eps} \\caption{{\\it Left:} Variations of the polar radius on the ZAMS as a function of $M_{\\rm ini}$ for various metallicities. {\\it Right:} Fraction of the total mass (m$_r$ is the mass inside the radius $r$) as a function of the fraction of the total radius for solar metallicity models with $\\Omega/\\Omega_{\\rm crit}$=0.1. } \\label{radii} \\end{figure}   \\begin{figure}% \\plottwo{meynet_fig3a.eps}{meynet_fig3b.eps} \\caption{{\\it Left:} Variation of the critical equatorial velocity  on the ZAMS as a function of the initial mass for various metallicities. {\\it Right:} Variation of the critical rotation period on the ZAMS as a function of the initial mass for various metallicities. } \\label{vcrit} \\end{figure}  Based on the set of numerical results described above, we first present useful relations between basic quantities, and then explore the initial conditions required for a star to reach the critical velocity during it MS evolution. ",
        "conclusions": ""
    },
    "0601/astro-ph0601080_arXiv.txt": {
        "abstract": "{In  order  to   determine  the  incidence  of  black  hole accretion-driven nuclear activity in nearby galaxies, as manifested by their radio  emission, we  have carried out  a high-resolution Multi-Element Radio-Linked Interferometer Network (MERLIN) survey  of  LINERs  and  composite  LINER/H{\\sc ii}  galaxies  from  a complete magnitude-limited  sample of bright  nearby galaxies (Palomar sample)  with  unknown arcsecond-scale  radio  properties.  There  are fifteen radio  detections, of  which three are  new subarcsecond-scale radio core detections, all being candidate AGN.  The detected galaxies supplement  the already  known low-luminosity AGN  --  low-luminosity Seyferts,  LINERs and  composite LINER/H{\\sc  ii} galaxies  --  in the Palomar sample.  Combining all radio-detected Seyferts,  LINERs and composite LINER/H{\\sc  ii} galaxies (LTS sources), we obtain an  overall radio  detection rate of  54\\% (22\\% of  all bright nearby galaxies) and we estimate that at least $\\sim$50\\% ($\\sim$20\\% of all bright nearby galaxies)  are true AGN.   The radio  powers of the LTS galaxies allow the construction of a local radio luminosity  function. By comparing the luminosity function with those of selected moderate-redshift AGN,  selected from the 2dF/NVSS survey, we find that LTS sources naturally extend the RLF of powerful AGN down to powers of about 10 times that of Sgr A*. ",
        "introduction": "The search for low-luminosity active galactic nuclei (LLAGN) in nearby galaxies has been  the subject of many optical  surveys.  Results show that nuclear activity may be  a common phenomenon.  The Palomar survey (Ho, Filippenko  \\& Sargent  1995, 1997a, b)  has been very  useful in this  regard  by  providing  a  sensitive  magnitude-limited  (B$_{\\rm T}<$12.5~mag)  sample of almost 500 bright nearby  galaxies. About  half  of the sources are  emission-line nuclei,  classified as Seyferts,  LINERs or composite LINER/H{\\sc ii} galaxies,  the last category displaying both LINER and H{\\sc ii}  properties.  However, characterizing the powering mechanisms  of the  sources  is not  straightforward, particularly  in low-luminosity sources.  Many of these galaxies possess circum-nuclear star-forming  regions which  blend with  and  may even  drown out  the presence of a weak active galactic nucleus (AGN).  Optimally it is necessary to  pick spectral regions where the contrast between  any hypothetical LLAGN  component and  circum-nuclear stellar component  is maximized. X-rays are very useful in this regard as shown by the hard X-ray studies of LLAGN  (Terashima \\etal  2000; Terashima, Ho \\& Ptak 2000; Ho \\etal  2001; Terashima \\&  Wilson 2003). In  the  absence of  spectral  AGN signatures such  as  a Seyfert  or quasar-type continuum or broad  emission lines, radio observations can offer  an alternative method  for determining  the LLAGN  incidence in nearby  galaxies.   Measurements  of  radio flux,  compactness,  radio spectral   index\\footnote{F$_{\\nu}   \\propto  \\nu^{-\\alpha}$ throughout.}   and  brightness   temperatures  provide  the  necessary diagnostic tools for determining the nature of the radio emission.    \\begin{table*}  \\setcounter{table}{0}  \\scriptsize  \\begin{center}  \\begin{minipage}[!ht]{119mm}  \\caption{Target MERLIN sources. Col.  1  Source name. Col. 2  and 3: Optical position from  NASA/IPAC Extragalactic Database (NED). Col. 4: Adopted distance  from  Ho,  Filippenko \\& Sargent (1997a) with H$_{\\rm 0}$=75 km s$^{-1}$ Mpc$^{-1}$. Col. 5: Spectral Type from Ho, Filippenko \\& Sargent (1997a). L = LINERs, S = Seyferts, and T  = composite LINER/H{\\sc ii} galaxies.  Colons refer to uncertain (:)  or highly uncertain (::)  spectral classification.  The number 1.9 refers  to the presence of broad H$\\alpha$ detection and  2 refers to the absence of  broad H$\\alpha$  emission. Col. 6:  Hubble type  from Ho, Filippenko \\& Sargent (1997a).}  \\begin{tabular}[h!]{l c c c c c r}  \\hline \\hline  & RA(J2000)  & Dec(J2000) & $D$   &  &  \\\\  Galaxy &  (\\hrs \\ramin \\rasec) & (\\degree \\arcmin \\arcsec~) & (Mpc) & Spectral Type & Hubble Type  \\\\  (1)    & (2)     & (3)              & (4)         & (5)   & (6)  \\\\  \\hline  IC\\,356   & 04 07 46.9 & $+$69 48 45 & 18.1   & T2   & SA(s)ab pec  \\\\ IC\\,520   & 08 53 42.2 & $+$73 29 27 & 47.0   & T2:    & SAB(rs)ab?  \\\\ NGC\\,428  & 01 12 55.6 & $+$00 58 54 & 14.9 & L2/T2: & SAB(s)m    \\\\ NGC\\,488  & 01 21 46.8 & $+$05 15 25 & 29.3 &  T2:: & SA(r)b   \\\\ NGC\\,521  & 01 24 33.8 & $+$01 43 52 & 67.0 & T2/H: & SB(r)bc       \\\\ NGC\\,718  & 01 53 13.3 & $+$04 11 44 & 21.4 & L2 & SAB(s)a  \\\\ NGC\\,777  & 02 00 14.9 & $+$31 25 46 & 66.5 & S2/L2:: & E1     \\\\ NGC\\,841  & 02 11 17.3 & $+$37 29 50 & 59.5 & L1.9: & (R')SAB(s)ab \\\\ NGC\\,1169 & 03 03 34.7 & $+$46 23 09 & 33.7   & L2   & SAB(r)b \\\\ NGC\\,1961 & 05 42 04.8 & $+$69 22 43 & 53.1   & L2   & SAB(rs)c\\\\ NGC\\,2336 & 07 27 03.7 & $+$80 10 42 & 2.9 & L2/S2 & SAB(r)bc\\\\ NGC\\,2681 & 08 53 32.8 & $+$51 18 50 & 13.3 & L1.9 & (R')SAB(rs)0/a \\\\ NGC\\,2768 & 09 11 37.5 & $+$60 02 15 & 23.7 & L2 & E6: \\\\ NGC\\,2832 & 09 19 46.9 & $+$33 44 59 & 91.6 & L2:: & E$+$2: \\\\ NGC\\,2841 & 09 22 02.6 & $+$50 58 35 & 12.0 & L2 & SA(r)b: \\\\ NGC\\,2985 & 09 50 21.6 & $+$72 16 44 & 22.4 & T1.9 & (R')SA(rs)ab   \\\\ NGC\\,3166 & 10 13 45.6 & $+$03 25 32 & 22.0 & L2 & SAB(rs)0/a     \\\\ NGC\\,3169 & 10 14 15.0 & $+$03 27 57 & 19.7 & L2 & SA(s)a pec    \\\\ NGC\\,3190 & 10 18 05.8 & $+$21 49 56 & 22.4 & L2 & SA(s)a pec spin \\\\ NGC\\,3226 & 10 23 27.0 & $+$19 53 54 & 23.4 & L1.9 & E2: pec    \\\\ NGC\\,3301 & 10 36 55.8 & $+$21 52 55 & 23.3 & L2 & (R')SB(rs)0/a     \\\\ NGC\\,3368 (M\\,96) & 10 46 45.7 & $+$11 49 12 & 8.1 & L2 & SAB(rs)ab  \\\\ NGC\\,3414 & 10 51 16.2 & $+$27 58 30 & 24.9 & L2 & S0 pec   \\\\ NGC\\,3433 & 10 52 03.9 & $+$10 08 54 & 39.5 & L2/T2:: & SA(s)c \\\\ NGC\\,3507 & 11 03 25.4 & $+$18 08 12 & 19.8 & L2 & SB(s)b      \\\\ NGC\\,3607 & 11 16 54.3 & $+$18 03 10 & 19.9 & L2 & SA(s)0: \\\\ NGC\\,3623 (M\\,65) & 11 18 55.9 & $+$13 05 32 & 7.3 & L2: & SAB(rs)a     \\\\ NGC\\,3626 & 11 20 03.8 & $+$18 21 25 &  26.3 & L2:  & (R)SA(rs)0$+$ \\\\ NGC\\,3627 (M\\, 66) & 11 20 15.0 & $+$12 59 30 & 6.6 & T2/S2 & SAB(s)b  \\\\ NGC\\,3628 & 11 20 17.0 & $+$13 35 22 & 7.7 & T2 & SAb pec spin \\\\ NGC\\,3646 & 11 22 14.7 & $+$20 12 31 & 56.8 & H/T2: & SB(s)a  \\\\ NGC\\,3675 & 11 26 07.9 & $+$43 35 10 & 12.8  & T2  & SA(S)b \\\\ NGC\\,3718 & 11 32 35.3 & $+$53 04 01 & 17.0 & L1.9 & SB(s)a pec    \\\\ NGC\\,3780 & 11 39 22.3 & $+$56 16 14 & 37.2 & L2:: & SA(s)c:   \\\\ NGC\\,3884$^a$ & 11 46 12.2 & $+$20 23 30 & 91.6 & L1.9 & SA(r)0/a;   \\\\ NGC\\,3898 & 11 49 15.2 & $+$56 05 04 & 21.9 & T2 & SA(s)ab;  \\\\ NGC\\,3945 & 11 53 13.6 & $+$60 40 32 & 22.5 & L2 & SB(rs)0$+$    \\\\ NGC\\,4013 & 11 58 31.3 & $+$43 56 49 & 17.0 & T2 & SAb     \\\\ NGC\\,4036 & 12 01 26.9 & $+$61 53 44 & 24.6 & L1.9 & S0$-$   \\\\ NGC\\,4143 & 12 09 36.1 & $+$42 32 03 & 17.0 & L1.9 & SAB(s)0   \\\\ NGC\\,4203 & 12 15 05.0 & $+$33 11 50 & 9.7 & L1.9 & SAB0$-$:   \\\\ NGC\\,4293 & 12 21 12.8 & $+$18 22 58 & 17.0 & L2 & (R)SB(s)0/a  \\\\ NGC\\,4321 (M\\,100) & 12 22 54.9 & $+$15 49 21 & 16.8 & T2 & SAB(s)bc  \\\\ NGC\\,4414 & 12 26 27.1 & $+$31 13 24 & 9.7 & T2: & SA(rs)c?  \\\\ NGC\\,4435 & 12 27 40.5 & $+$13 04 44 & 16.8 & T2/H:  & SB(s)0  \\\\ NGC\\,4438 & 12 27 45.6 & $+$13 00 32 & 16.8 & L1.9 & SA(s)0/a \\\\ NGC\\,4450 & 12 28 29.5 & $+$17 05 06 & 16.8 & L1.9 & SA(s)ab  \\\\ NGC\\,4457 & 12 28 59.1 & $+$03 34 14 & 17.4 & L2 & (R)SAB(s)0/a  \\\\ NGC\\,4548 (M\\,91) & 12 35 26.4 & $+$14 29 47 & 16.8 & L2 & SBb(rs)   \\\\ NGC\\,4589 & 12 37 25.0 & $+$74 11 31 & 30.0 & L2 & E2 \\\\ NGC\\,4636 & 12 42 50.0 & $+$02 41 17 & 17.0 & L1.9 & E0$+$    \\\\ NGC\\,4736 (M\\,94) & 12 50 53.0 & $+$41 07 14 & 4.3 & L2 & (R)SA(r)ab  \\\\ NGC\\,4750 & 12 50 07.1 & $+$72 52 30 & 26.1 & L1.9 & (R)SA(rs)ab   \\\\ NGC\\,4772 & 12 53 29.0 & $+$02 10 02 & 16.3 & L1.9 & SA(s)a  \\\\ NGC\\,4826 (M\\,64) & 12 56 43.7 & $+$21 40 52 & 4.1 & T2 & (R)SA(rs)ab   \\\\ NGC\\,5077$^b$ & 13 19 31.6 & $-$12 39 26 & 40.6 & L1.9 & E3$+$\\\\ NGC\\,5297 & 13 46 23.7 & $+$43 52 20 & 37.8 & L2 & SAB(s)c: spin   \\\\ NGC\\,5322 & 13 49 15.2 & $+$60 11 26 & 31.6 & L2:: & E3$+$ \\\\ NGC\\,5353 & 13 53 26.7 & $+$40 16 59 & 37.8 & L2/T2: & SA0 spin  \\\\ NGC\\,5363 & 13 56 07.1 & $+$05 15 20 & 22.4 & L2 & IA0? \\\\ NGC\\,5371 & 13 55 39.9 & $+$40 27 43 & 37.8 & L2 & SAB(rs)bc \\\\ NGC\\,5377 & 13 56 16.6 & $+$47 14 08 & 31.0 & L2 & (R)SB(s)a    \\\\ NGC\\,5448 & 14 02 49.7 & $+$49 10 21 & 32.6 & L2 & (R)SAB(r)a      \\\\ NGC\\,5678 & 14 32 05.8 & $+$57 55 17 & 35.6 & T2 & SAB(rs)b     \\\\ NGC\\,5813 & 15 01 11.2 & $+$01 42 08 & 28.5 & L2: & E1$+$         \\\\ NGC\\,5838 & 15 05 26.2 & $+$02 05 58 & 28.5 & T2:: & SA0$-$  \\\\ NGC\\,6340 & 17 10 24.9 & $+$72 18 16 & 22.0 & L2 & SA(s)0/a  \\\\ NGC\\,6501 & 17 56 03.7 & $+$18 22 23 & 39.6 & L2:: & SA0$+$:   \\\\ NGC\\,6702 & 18 46 57.6 & $+$45 42 20 & 62.8 & L2:: & E: \\\\ NGC\\,6703 & 18 47 18.8 & $+$45 33 02 & 35.9 & L2:: & SA0$-$     \\\\  \\hline  \\end{tabular}  \\smallskip  $^a$ Does not fulfill the magnitude criterion for the Palomar sample.  $^b$ Does not fulfill declination criterion for the Palomar sample.  \\end{minipage}  \\end{center}  \\end{table*}  \\normalsize   Several   important  radio   surveys  have   been  conducted   on  the magnitude-limited Palomar bright  nearby galaxy sample (Ho, Filippenko \\& Sargent 1995, 1997a, b), revealing a large fraction of radio cores, not only  in ellipticals  but also in  bulge-dominated spirals.   In a recent  Very Large  Array  (VLA) 5  and 1.4\\,GHz,  1\\arcsec~resolution survey of  the low-luminosity  Seyferts of the  Palomar sample  (Ho \\& Ulvestad 2001; HU01 hereafter), it was found  that over 80\\% of the sources harbour a radio core.  In a distance-limited sample of low- luminosity  Seyferts, LINERs  and composite  LINER/H{\\sc  ii} galaxies observed with  the VLA  at 15\\,GHz, 0\\arcsecpoint25  resolution (Nagar \\etal 2000, 2002; VLA/N00 and VLA/N02 hereafter; Nagar, Falcke \\& Wilson 2005; VLA/N05 hereafter), it was found that $\\sim$40\\% of the objects harbour subarcsecond-scale compact radio cores.  A similar study with the VLA, at  8.4\\,GHz, 2\\arcsecpoint5~resolution of  all the  composite LINER/H{\\sc ii}  galaxies in  the  Palomar  sample has  been  presented in  Filho, Barthel \\&  Ho (2000, 2002a),  revealing radio cores in  $\\sim$25\\% of the  sample sources.   However, although  the radio  core  emission in these sources  is consistent with the  presence of a  LLAGN, we cannot exclude  a  stellar origin  from  the  brightness temperature  figures (T$_{\\rm   B}$\\lax10$^{5}$~K;   Condon   1992)   obtained   at   these resolutions.   As  conclusive judgement  requires  Very Long  Baseline (VLBI)-resolution,  multi-wavelength Very  Long Baseline  Array (VLBA) and   European  Very  Long   Baseline  Interferometer   Network  (EVN) observations   have   been  obtained   for   selected  subsamples   of low-luminosity Seyferts, LINERs and composite LINER/H{\\sc ii} galaxies that showed arcsecond- or subarcsecond-scale radio cores (Falcke \\etal 2000; F00 hereafter; Nagar  \\etal 2002; VLBA/N02 hereafter; Nagar, Falcke \\& Wilson 2005; VLBA/N05 hereafter; Ulvestad \\&  Ho 2001b; Filho, Barthel  \\& Ho 2002b; Filho \\etal  2004; Anderson, Ulvestad \\& Ho  2004; AU04 hereafter).  In sources with  subarcsecond-  or  arcsecond-scale  radio  peak  emission  above 2.5~mJy,  results reveal  a 100\\%  detection rate  of  high-brightness temperature   (T$_{\\rm   B}$\\gax10$^8$~K),   compact,  flat   spectrum ($\\alpha$$<$0.5) radio cores, enforcing the LLAGN scenario for the radio emission (VLBA/N05; see also Ulvestad \\& Ho 2001b; Filho, Barthel \\& Ho 2002b; Filho \\etal 2004; AU04). Their low  radio luminosities suggest we are  probing the very faint end of the AGN population.  Unambiguously  determining the  physical nature  of the  nearby galaxy radio cores is  more than of mere phenomenological  interest.  If they truly contain  an accretion-powered nucleus, then  they obviously need to  be included  in  the AGN  population.   Their non-trivial  numbers impact on several astrophysical problems ranging from the cosmological evolution of the AGN luminosity function  to their contribution to the X-ray background.  The present paper  deals with high-resolution radio-imaging of LLAGN, carried   out  with  the   Jodrell  Bank   Multi-Element  Radio-Linked Interferometer  Network  (MERLIN), completing  the  radio census of  nuclear activity  in the Palomar galaxy survey of  Ho, Filippenko  \\& Sargent (1997a). ",
        "conclusions": "We have  undertaken a  MERLIN survey of  nearby galaxies that  did not have  available  2\\arcsecpoint5~resolution  or  better  radio  observations. Results reveal  a 21\\%  radio-detection rate  among the  sources, with fifteen radio  detections, three of  which are new AGN  candidates.  A compilation  of  radio observations  of  all low-luminosity  Seyferts, LINERs and composite LINER/H{\\sc ii} galaxies in the magnitude-limited Palomar survey reveal  a radio-detection rate of 54\\%  (or 22\\% of all bright nearby galaxies), with a more than 50\\% detection rate (or 20\\% for all bright nearby  galaxies) of low-luminosity active nuclei. The radio detection of the  Seyferts, LINERs  and composite  LINER/H{\\sc ii} sources in the Palomar survey  allow the construction of a local radio luminosity  function.  Our results show  that  the  Seyferts, LINERs  and composite LINER/H{\\sc ii} sources  form a smooth luminosity transition from higher  redshift, more  luminous AGN as  sampled by  the 2dF/NVSS survey."
    },
    "0601/astro-ph0601049_arXiv.txt": {
        "abstract": "{This paper presents a method to determine effective temperatures, angular semi-diameters and bolometric corrections for population I and II FGK type stars based on $V$ and \\MASS IR photometry. Accurate calibration is accomplished by using a sample of solar analogues, whose average temperature is assumed to be equal to the solar effective temperature of 5777 K. By taking into account all possible sources of error we estimate associated uncertainties better than 1\\% in effective temperature and in the range 1.0--2.5\\% in angular semi-diameter for unreddened stars. Comparison of our new temperatures with other determinations extracted from the literature indicates, in general, remarkably good agreement. These results suggest that the effective temperaure scale of FGK stars is currently established with an accuracy better than 0.5\\%--1\\%. The application of the method to a sample of 10999 dwarfs in the Hipparcos catalogue allows us to define temperature and bolometric correction ($K$ band) calibrations as a function of $(V-K)$, $[m/H]$ and \\logg. Bolometric corrections in the $V$ and $K$ bands as a function of \\teft, $[m/H]$ and \\logg are also given. We provide effective temperatures, angular semi-diameters, radii and bolometric corrections in the $V$ and $K$ bands for the 10999 FGK stars in our sample with the corresponding uncertainties. ",
        "introduction": "Effective temperature and luminosity are two fundamental stellar parameters that are crucial to carry out tests of theoretical models of stellar structure and evolution by comparing them with observations. The accuracy in the determination of other stellar properties, such as metallicity, age or radius, hinges on our ability to estimate the effective temperatures and luminosities.  There are several approaches in the literature to compute effective temperature and/or luminosity. Except when applied to the Sun, very few of them are {\\it direct} methods that permit an empirical measurement of these parameters. Usually, {\\it semi-empirical} or {\\it indirect} methods are based to a certain extent on stellar atmosphere models. Among the {\\it direct} approaches we find the remarkable work by \\citet{code76}, which is based on interferometric measurements of stellar angular semi-diameters ($\\theta$) and total fluxes ($F_{\\rm bol}$) at Earth, and the more recent works of \\citet{mozurkewich03} and \\citet{kervella04}, also based on interferometry. On the other hand, {\\it indirect} methods are mainly based on the use of photometry, spectroscopy, or a combination of both. In the case of the temperatures, although many of the published calibrations claim to have uncertainties of the order of several tens of degrees, values obtained by different authors can easily have discrepancies of 100 K or even larger. The reason for such nagging differences must be found somewhere in the {\\it ingredients} of the methods: atmosphere models, absolute flux calibrations, oscillator strengths, calibration stars, etc.  In this paper we present a {\\it semi-empirical} method to determine effective temperatures (\\teft) and bolometric corrections ($BC$) from 2MASS\\footnote{\\tt http://www.ipac.caltech.edu/2mass} $JHK$\\footnote {Throughout the paper, $K$ refers to $K_s$ band.} photometry \\citep{cutri03b} that is applicable to FGK type stars. As all others, our method is susceptible to problems derived from the uncertainties in the ingredients mentioned above. However, our approach benefits from two major features: First, it provides a way to evaluate realistic individual uncertainties in \\teft, $\\theta$ and luminosity by considering all the involved errors; and second, as it is calibrated to use the \\MASS photometry, it allows the calculation of consistent and homogeneous \\tefts and $BC$ for several million stars in the \\MASS catalogue. This paper also provides \\teft, angular semi-diameters, radii and $BC$s for 10999 dwarfs and subdwarfs in the Hipparcos catalogue \\citet{ESA97}. Such large sample has allowed us to construct simple parametric calibrations as a function of $(V-K)_0$, \\feh and \\logG. Note that a preliminary version of the method presented here was already successfully applied to the characterization of the properties of planet-hosting stars \\citep{iribas03a}.  The present paper is organized as follows. Section \\ref{sedf} presents the method and explains in detail the procedure to obtain \\tefts and angular semi-diameters, including the fitting algorithm, zero point corrections and error estimates. The comparison of our temperatures with several previous works, both based on photometric and spectroscopic techniques, is described in Sect. \\ref{comparacions}. In Sect. \\ref{calib} we present simple parametric calibrations of \\tefts and $BC$ as a function of $(V-K)_0$, $[m/H]$ and \\logg valid for dwarf and subdwarf stars. The sample of 10999 stars used to build the calibrations is also described in this section together with a detailed explanation of the different contributors to the final uncertainties.  Finally, the results are discussed in Sect. \\ref{disc} and the conclusions of the present work are presented in Sect. \\ref{conclu}. ",
        "conclusions": "\\label{conclu}  We have presented a method (called SEDF) to compute effective temperatures, angular semi-diameters and bolometric corrections from \\MASS photometry. We have adopted an approach based on the fit of the observed $VJHK$ magnitudes using synthetic photometry, and it yields accuracies around 1\\% in \\teft, 2\\% in $\\theta$, and 0.05 mag in $BC$, in the temperature range 4000--8000 K.  A zero point offset was added to the synthetic photometry computed from the Kurucz atmosphere models to tie in our temperature scale with the Sun's temperature through a sample of solar analogues. From the application to a large sample of FGK Hipparcos dwarfs and subdwarfs we provide parametric calibrations for both effective temperature and bolometric correction as a function of $(V-K)_0$, $[m/H]$ and \\logG. Note that the method presented here has been selected as one of the main sources of effective temperatures to characterize the primary and secondary targets of the COROT space mission \\citep {baglin00}. Also, it is being currently implemented as one of the tools offered by the Spanish Virtual Observatory \\citep{solano05}.  The resulting temperatures have been compared with several photometric and spectroscopic determinations. Although we obtained remarkably good agreement, slight systematic differences with other semi-empirical methods, such as the IRFM, are present. This is probably due to the uncertainties in the absolute flux calibration used by different techniques. It is possible that, in spite of the great effort carried out by \\citet{cohen03a} and others to construct a consistent absolute flux calibration in both the optical and the IR regions, some problems still remain, which introduce small systematic effects in the temperatures. However, these effects seem to be as small as 20--30 K and could be explained through uncertainties in the IR fluxes of about 2\\%. In conclusion, the results presented here strongly suggest that, given the small differences found between methods, the effective temperature scale of FGK stars (4000--8000 K) is currently established with a net accuracy better than 0.5--1.0\\%."
    },
    "0601/astro-ph0601519_arXiv.txt": {
        "abstract": "We present an efficient second order accurate scheme to treat stiff source terms within the framework of higher order Godunov's methods. We employ Duhamel's formula to devise a modified predictor step which accounts for the effects of stiff source terms on the conservative fluxes and recovers the correct isothermal behavior in the limit of an infinite cooling/reaction rate. Source term effects on the conservative quantities are fully accounted for by means of a one-step, second order accurate semi-implicit corrector scheme based on the deferred correction method of Dutt et. al. We demonstrate the accurate, stable and convergent results of the proposed method through a set of benchmark problems for a variety of stiffness conditions and source types. ",
        "introduction": "We wish to solve the following system of partial differential equations describing a hydrodynamic flow with a stiff (energy) source term \\begin{equation} \\label{hypsys:eq} \\frac{\\partial U}{\\partial t} + \\sum _{d=1}^{D} \\frac{\\partial F_d(U)}{\\partial x_d}  = S(U) \\end{equation} where $D$ is the dimensionality of the problem, $U,~F(U),~S(U)$ are the conservative variables, the conservative fluxes and the source term respectively, given by \\begin{equation} U= \\begin{pmatrix} \\rho \\\\ \\rho u_1 \\\\ \\vdots \\\\ \\rho u_D \\\\ \\rho E \\end{pmatrix} ; \\quad F_d(U)= \\begin{pmatrix} \\rho u_d \\\\ \\rho u_1 u_d + p \\,\\delta_{1d} \\\\ \\vdots \\\\ \\rho u_D u_d + p \\,\\delta_{Dd} \\\\ (\\rho E + p) u_d \\end{pmatrix}; \\quad S(U)= \\begin{pmatrix} 0 \\\\ 0 \\\\ \\vdots \\\\ 0 \\\\ \\rho\\Lambda(e,\\rho) \\end{pmatrix}. \\label{eq:ugs} \\end{equation} In the above equations, $\\rho$ is the density, $u_d$ the velocity in the $d$ direction, $E=e+ \\sum_{d=1}^Du_d^2/2$ is the total specific energy with, $e$, the specific internal energy.  $\\Lambda(e,\\rho)$ is the term describing the source of specific internal energy.  In the following we consider the case of a stiff source term corresponding to an endothermic process, such as occurs in radiative losses.  In addition, we restrict our analysis to source types that, at least near equilibrium, behave as a relaxation law.  In the stiff case, the characteristic relaxation time scale for $S$ may be much smaller than the CFL time step for the hydrodynamic waves. For that reason, we would like to use a semi-implicit method, treating the stiff source term implicitly, while using an explicit method for the hyperbolic terms. However, the classical analysis of such fast endothermic processes shows that, in the limit as the relaxation time goes to zero, the gas can be described by the compressible flow equations with an isothermal equation of state \\cite{vincentiKruger65}. Pember \\cite{pember93} showed that the use of formally second-order accurate semi-implicit methods such as Strang splitting, or a second-order Godunov predictor-corrector method could lead to a substantial loss of accuracy, due to inconsistencies between the the characteristic tracing step without sources and the effective limiting isothermal behavior. Such inconsistencies between the flux calculation with the limiting isothermal equation of state can lead to dramatic errors particularly at sonic points. Pember proposed various approaches to the problem based on classical relaxation theory.  Roe and Hittinger \\cite{roehit01} also addressed the issues raised here in relation to Godunov's method with stiff relaxation. In their approach they split the equations based on a splitting of state space into stiff and non-stiff subspaces of the linearized source term to obtain in the stiff limit formulations similar to ours. However, neither Pember nor Roe and Hittinger did present a complete method that is second-order accurate in both the stiff and non-stiff limits, nor did they discuss the extension to more than one dimension.  The problem of hyperbolic system with stiff relaxation has also been considered by other authors in the past mostly for one-dimensional systems and within the framework of Runge-Kutta based methods of lines.  In particular Jin \\cite{jin95} designed second order Runge-Kutta type splitting methods with the correct asymptotic limit. Jin and Levermore's \\cite{jinlevermore96} developed a semidiscrete high resolution method which, in order to ensure the correct asymptotic behavior, employes a linear combination of the conservative fluxes for the homogeneous (i.e. without the relaxation term) and equilibrium system.  The fluxes are computed with a higher order Godunov's method and the scheme allows for a rapid transition between the stiff and nonstiff regimes.  As the authors point out, however, the upwind property of the scheme is not strictly guaranteed for all stiffness conditions.  Finally, Caflisch, Jin and Russo \\cite{caflisch97} developed a scheme for hyperbolic systems with relaxation that is uniformly accurate for various ranges of stiffness conditions (see also Ref.~\\cite{jinpareschitoscani98,pareschirusso05}).  The aim of this paper is to build a higher order Godunov's method that preserves the properties of robustness and accuracy across a variety of stiffness conditions thus avoiding the problems described in \\cite{pember93}.  In particular, in order to preserve higher order accuracy, we aim for a semi-implicit method that corresponds to a standard second-order Godunov method of the appropriate hyperbolic problem for the stiff or non-stiff limits. To this end, we use second-order accurate deferred corrections method of a type presented in \\cite{dugrro00}, to obtain a semi-implicit corrector that is  a special case of the algorithms described in \\cite{minion03}, although any implicit L-stable second-order one-step method would be acceptable. The main new idea in our work is contained in our treatment of the predictor step for computing the hyperbolic fluxes, based on the derivation of a local effective dynamics using Duhamel's formula. This leads to an explicit predictor step that corresponds to that for a conventional second-order Godunov method for Eq.~(\\ref{hypsys:eq}) in the limit where the relaxation time is comparable to or greater than the hydrodynamic CFL time step; and to a second-order Godunov method for the isothermal equations in the limit where the relaxation time is much smaller than the hydrodynamic time step.  Our approach is similar to that used in \\cite{tcm05} for obtaining a well-behaved numerical method for incompressible viscoelastic flows in both the viscous and elastic limits; however, the details there are quite different than those for the present setting.  The paper is organized as follows. In section~\\ref{sdc:se} we describe a second order accurate, semi-implicit corrector method based on the deferred corrections ideas presented in~\\cite{dugrro00,minion03} to be used for the final source term update.  In section~\\ref{mgf:se}, based on Duhamel's formula, we work out a modified formulation of Godunov's predictor step and flux calculation suitable for the case of stiff source terms.  In section~\\ref{sa:se} we discuss stability issues for our approach, and Section~\\ref{lrnz:se} contains the extension of the method to the case in which the source term depends both on the gas density as well as the internal energy.  In section~\\ref{test:se} we test the performance of the code and demonstrate the accuracy of the method in various stiffness conditions.  The paper concludes with section~\\ref{concl:se} where the main results of the paper are summarized. ",
        "conclusions": "\\label{concl:se}  We have presented a second order accurate semi-implicit predictor-corrector scheme to treat stiff source terms within the framework of higher order Godunov's methods.  Our treatment of the predictor step for computing the hyperbolic fluxes, is based on the derivation of a local effective dynamics using Duhamel's formula. This leads to a conventional second-order Godunov method when the system relaxation time is larger than the time step and to a second-order Godunov method for the isothermal equations in the limit of a stiff source term.  Finally, we obtain a semi implicit corrector using a one-step second-order accurate deferred corrections method as suggested in \\cite{dugrro00,minion03}.  Our tests indicate that the proposed method is stable, robust and its second order accuracy preserved across a variety of stiffness conditions.  We have also discussed the case of a general source term which depends both on $e$ and $\\rho$ and shown that the method is applicable provided that the flow is thermally stable or the non-stiff part of the source term is resolved in time.  The additional cost involved in the formulation of our scheme is minimal; all it requires is an estimate of the term $\\Lambda_e$ which in a purely relaxation case is trivial and for a more complicated source term (such as the case of radiative losses) is still minor compared to the estimate of the source term itself. In our implementation the factor $\\alpha(\\gamma-1)+1$ is stored as an additional primitive variable and used as polytropic index in the characteristic analysis in stead of $\\gamma$.   \\vskip 1truecm \\leftline{\\bf Acknowledgment} FM is grateful to the Lawrence Berkeley National Laboratory for its hospitality and acknowledges support by the Swiss Institute of Technology through a Zwicky Prize Fellowship. PC was supported by the Mathematical, Information, and Computing Sciences Division of the United States Department of Energy Office of Science under contract number DE-AC02-05CH11231."
    },
    "0601/astro-ph0601033_arXiv.txt": {
        "abstract": "We determine the disk mass distribution around 336 stars in the young  ($\\sim 1$ Myr) Orion Nebula cluster by imaging a $2\\rlap{.}'5 \\times 2\\rlap{.}'5$ region in $\\lambda$3 mm continuum emission with the Owens Valley Millimeter Array.  For this sample of 336 stars, we observe 3 mm emission above the  3$\\sigma$ noise level toward ten sources, six of which have also been detected optically in silhouette against the bright nebular background.  In addition, we detect 20 objects in 3 mm continuum emission that do not correspond to known near-IR cluster members. Comparisons of our measured fluxes with longer wavelength observations enable rough separation of dust emission from thermal free-free emission, and we find substantial dust emission toward most objects.  For the sample of ten objects detected at both 3 mm and near-IR wavelengths, eight exhibit substantial dust emission. Excluding the two high-mass stars ($\\theta^1$OriA and the BN object) and assuming a gas-to-dust ratio of 100, we estimate circumstellar masses ranging from 0.13 to 0.39 M$_\\odot$. For the cluster members not detected at 3 mm, images of individual objects are stacked to constrain the mean 3 mm flux of the ensemble.  The average flux is detected at the 3$\\sigma$ confidence level, and implies an average disk mass of 0.005 M$_{\\odot}$, comparable to the minimum mass solar nebula. The percentage of stars in Orion surrounded by disks more massive than $\\sim 0.1$ M$_{\\odot}$ is consistent with the disk mass distribution in Taurus, and we argue that massive disks in Orion do not appear to be truncated through close encounters with high-mass stars. Comparison of the average disk mass and number of massive dusty structures in Orion with similar surveys of the NGC 2024 and IC 348 clusters is used to constrain the evolutionary timescales of massive circumstellar disks in clustered environments. ",
        "introduction": "Over the last two decades, high resolution millimeter, infrared, and optical images have provided evidence for the existence of circumstellar disks on scales of $\\sim 0.1$--1000 AU around young stars \\citep[e.g.,][]{KS95,DUTREY+96, PADGETT+99,OW96,EISNER+04}.  Circumstellar disks are the likely birth-sites for planetary systems, and determining their ubiquity, properties, and lifetimes is crucial for constraining the timescales and mechanisms of planet formation. The mass distribution of protoplanetary disks is especially important since disk mass is related to the mass of planets that may potentially form. The minimum-mass protosolar nebula, $\\sim 0.01$ M$_{\\odot}$ \\citep{WEID+77,HAYASHI81}, is an informative benchmark against which to compare disk masses around young stars, and such comparisons can constrain the number of protoplanetary disks with the potential to form planetary systems like our own solar system.  While direct imaging at optical through near-IR wavelengths has provided concrete evidence for a limited number of circumstellar disks \\citep[e.g.,][]{OW96}, and observations of near-IR excess emission have shown statistically that most young stars with ages less than a few million years still possess inner circumstellar disks \\citep[e.g.,][]{STROM+89,HLL01b}, these studies did not constrain the disk mass distribution. To probe the bulk of the disk mass, which resides in cooler, outer disk regions, observations of optically-thin millimeter emission are needed.  Several investigators have carried out comprehensive single-dish mm and sub-mm continuum surveys toward regions of star formation comprising loose aggregates of stars: Taurus \\citep{BECKWITH+90,OB95,MA01,AW05}, $\\rho$ Ophiuchi \\citep{AM94,NUERNBERGER+98, MAN98}, Lupus \\citep{NCZ97}, Chamaeleon I \\citep{HENNING+93}, Serpens \\citep{TS98}, and MBM 12 \\citep{ITOH+03,HOGERHEIJDE+02}. In Taurus, 37\\% of stars appear to possess disks more massive than $\\sim 0.01$ M$_{\\odot}$, and the median disk mass is $5 \\times 10^{-3}$ M$_{\\odot}$ \\citep{AW05}.  The fraction of massive disks and the median disk mass is comparable in $\\rho$ Ophiuchi \\citep{AM94}.  Expanding millimeter continuum surveys to include rich clusters allows the determination of accurate statistics on the frequency and evolution of massive disks as a function of both stellar mass and age. Also, since most stars in the Galaxy form in rich clusters \\citep{LADA+91,LSM93,CARPENTER00,LL03}, understanding disk formation and evolution in cluster environments is a vital component in our general understanding of how stars and planets form. The main challenge to observing rich clusters at mm-wavelengths is that very high angular resolution is required to resolve individual sources and to distinguish compact disk emission from the more extended emission of the molecular cloud.  Single-aperture mm-wavelength telescopes lack sufficient angular resolution, and to date, only three rich clusters have been observed with mm-wavelength interferometers: the Orion Nebula cluster \\citep{MLL95,BALLY+98,WAW05}, IC 348 \\citep{CARPENTER02}, and NGC 2024 \\citep{EC03}.  The enhanced sensitivity of the most recent observations of Orion has enabled the detection of several massive ($\\ga 0.01$ M$_{\\odot}$) disks \\citep{WAW05}, while upper limits from other surveys range from $\\sim 0.025$--0.17 M$_{\\odot}$ \\citep{MLL95,BALLY+98}.  Moreover, extending detection limits by considering as ensembles the large numbers ($\\ga 100$) of young stars included in the surveys of IC 348 and NGC 2024 allowed estimates of mean disk masses of $\\sim 0.002$ and 0.005 M$_\\odot$, respectively \\citep{CARPENTER00,EC03}. While it appears that an average star aged $\\la 1$ Myr still possesses a massive circumstellar disk, more sensitive observations are necessary to detect directly large numbers of massive disks at a range of ages, and thereby constrain the mass distribution and evolutionary timescales.  Here, we present a new mm-wavelength interferometric survey of the Orion Nebula cluster (ONC), a young, deeply embedded stellar cluster that includes the bright, massive Trapezium stars. The Trapezium region contains hundreds of stars within several arcminutes, and pre-main-sequence evolutionary models \\citep[e.g.,][]{DM94} fitted to spectroscopic and/or photometric data indicate that most stars are less than approximately one million years old \\citep[e.g.,][]{PROSSER+94,HILLENBRAND97}. Moreover, the standard deviation in the distribution of inferred stellar ages is $\\la 1$ Myr \\citep{HILLENBRAND97}. Our observations thus provide a snapshot of millimeter emission around a large number of roughly coeval young stars.  With the large number of stars in the ONC, we can begin to investigate the correlation of disk properties with stellar and/or environmental properties.  Previous investigations of near-IR excess emission have explored the dependence of inner disk properties on stellar mass, age, and environment \\citep[e.g.,][]{HILLENBRAND+98,LADA+00}. For example, the fraction of stars in Orion exhibiting near-IR excess emission seems largely independent of stellar age and mass, although there are indications of a paucity of disks around very massive stars \\citep{HILLENBRAND+98,LADA+00}.  In addition, the inner disk fraction may decrease at larger cluster radii \\citep{HILLENBRAND+98}. To explore how the properties of the outer disk component correlate with such stellar and environmental properties, millimeter observations of cool, optically-thin dust emission are necessary.  Our observations represent an improvement over previous work because our mosaicked image encompasses more than three times as many sources as the previous surveys, enabling an improvement of $\\sqrt{3}$ in estimates of frequency and mean mass of circumstellar disks.  Moreover, the comparable sensitivity of this survey with previous observations of IC 348 and NGC 2024 allows a more direct comparison between relatively young (NGC 2024: 0.3 Myr; Meyer 1996; Ali 1996), intermediate (ONC: 1 Myr; Hillenbrand 1997), and old (IC 348: 2 Myr; Luhman et al. 1998; Luhman 1999) clusters, providing constraints on timescales for disk evolution. ",
        "conclusions": "} We observed a $2\\rlap{.}'5 \\times 2\\rlap{.}'5$ region of the Orion Nebula cluster in $\\lambda$3 mm continuum emission with the Owens Valley Millimeter Array.  The mosaic encompassed 336 young stars in the vicinity of the massive Trapezium stars, and constrained the disk mass distribution for this large number of cluster members.  We detected 30 objects in 3 mm continuum emission above the  3$\\sigma$ limit, 10 of which correspond with near-IR cluster members.  Six of these, in turn, also correspond with optically-detected proplyds. Comparison of our measured fluxes with longer wavelength observations enabled rough separation of emission due to dust and that due to thermal free-free emission, and we found that the 3 mm emission toward 8 objects likely arises, in part, from dust.  We argued that with the exception of the massive stars $\\theta^1$OriA and the BN object, sources detected in both 3 mm and near-IR emission are probably young stars surrounded by disks, and we computed circumstellar masses of $0.13$-$0.39$ M$_{\\odot}$ based on observed 3 mm fluxes (these masses are uncertain by at least a factor of three due to uncertainties in converting flux into mass). Since the vast majority ($\\ga 98\\%$) of near-IR cluster members do not possess disks more massive than $\\sim 0.1$ M$_{\\odot}$, we placed constraints on lower-mass disks by considering the ensemble of 326 non-detected, predominantly low-mass stars.  For the ensemble, we computed an average disk mass of 0.005 M$_{\\odot}$, which has a statistical significance of 3$\\sigma$ (although the absolute value of the mass contains larger uncertainties related to converting flux into mass).  The average disk in the ONC is thus comparable to the minimum mass solar nebula, suggesting that most young ($\\la 1$ Myr) stars in richly clustered environments (and by extension, most stars in the Galaxy) are surrounded by sufficient circumstellar masses to form solar systems like our own.  Furthermore, we found that the frequency of massive disks in the ONC appears similar to that in Taurus, perhaps refuting suggestions that disks in Orion are truncated due to close encounters with the massive Trapezium stars. Finally, we compared the disk mass distributions in three clusters of different ages to begin to constrain the evolutionary timescales of massive disks.  Although substantial uncertainties remain, it appears that massive disks may evolve significantly on 1-2 Myr timescales.   \\noindent{\\bf Acknowledgments.}  JAE is currently supported by a Miller Research Fellowship, and acknowledges past support from a Michelson Graduate Research Fellowship. JMC acknowledges support from the Owens Valley Radio Observatory, which is supported by the National Science Foundation through grant AST-9981546.  The authors also wish to thank John Bally for providing HST images of the Orion proplyds."
    },
    "0601/astro-ph0601633.txt": {
        "abstract": "\\spitzer\\ IRAC and MIPS images of the Trifid Nebula (M20)  reveal its spectacular appearance in infrared light, highlighting the nebula's special evolutionary stage. The images feature recently-formed massive protostars and numerous young stellar objects, and a single O star that illuminates the  surrounding molecular cloud from which it formed, and unveil large-scale, filamentary dark clouds. The  hot dust grains show contrasting infrared colors in  shells,  arcs, bow-shocks and dark cores.  Multiple protostars are detected in the infrared, within the cold dust cores of TC3 and TC4, which were previously defined as Class 0. %because they are the cold dust condensations and presence of %molecular outflows. The cold dust continuum cores of TC1 and TC2 contain only one protostar each; the newly discovered  infrared source in TC2 is   the driving source of the HH399 jet.   The \\spitzer\\ color-color diagram allowed us to identify $\\sim$160 young stellar objects (YSOs) and classify them into different evolutionary stages. The diagram also revealed a unique group of YSOs which are bright at 24$\\mu$m but have the spectral energy distribution peaking at 5-8$\\mu$m.   Despite expectation that Class 0 sources would be ``starless\" cores,   the \\spitzer\\ images, with unprecedented sensitivity,  uncover mid-infrared emission from these Class 0 protostars.  The mid-infrared detections of Class 0 protostars show that the emission  escapes the dense, cold envelope of young protostars. The mid-infrared emission of the  protostars can be fit by two temperatures of 150 and 400 K; the hot core region is probably optically thin in the mid-infrared regime, and  the size of hot core is much smaller than that of the cold envelope. %the mid-infrared emission %cannot arise from the same location as the mm-wave emission, and %instead must arise from a much smaller region with less intervening %extinction to the central object which drives the bipolar outflows. The presence of multiple protostars within the cold cores of Class 0 objects implies that clustering occurs at this early stage of star formation. The TC3 cluster shows that the most massive star is located at the center of the cluster and at the bottom of the gravitational-potential well. ",
        "introduction": "The Trifid Nebula (M20), is one of the best-known astrophysical objects: a classical nebula of ionized gas from an O7 star (HD 164492). The nebula glows  in red light, trisected by obscuring  dust lanes, with a reflection nebula in the north.  The Trifid Nebula is a very young  \\ion{H}{2} region with an age of $\\sim 3\\times 10^{5}$ years. The Infrared Space Observatory (ISO) and the Hubble Space Telescope (HST) (Cernicharo et al.\\ 1998; Lefloch \\& Cernicharo 2000; Hester et al.\\ 1999) show the Trifid to be a dynamic, ``pre-Orion\" star forming region containing young stellar objects (YSOs) undergoing episodes of violent mass ejection, and protostars (like HH399) losing mass and energy to the nebula in jets. Four massive (17--60 M$_\\odot$) protostellar cores were  discovered, with millimeter-wave observations, in the Trifid (Lefloch \\& Cernicharo 2000). A number of YSO candidates were identified  using near-infrared  (Rho et al. 2001) and X-ray observations (Rho et al. 2004). We adopted a distance of 1.68 kpc for the Trifid Nebula, which was measured by Lynds et al. (1986). There have been two primary tools for studying massive star formation: molecular line emission which traces dense material and outflows  (e.g. Shirley et al. 2003), and dust continuum emission which traces cold, dense material (e.g. William et al. 2004).  The youngest objects, Class 0 protostars, are those that are detected as condensations in dust continuum maps characterized by very low values of the ratio L$_{bol}$/L$_{submm}$, and show collimated CO outflows or internal-heating sources (Andr\\'e 1994). The Class 0 protostars were believed to be ``starless\" cores,  with neither near-infrared nor mid-infrared ($<$20$\\mu$m) emission. In this paper, we present spectacular \\spitzer\\ images of the Trifid Nebula and report $\\sim$160 newly identified young stellar objects (YSOs)  using their infrared colors. We  found  multiple protostars within cold cores and many evolved YSOs are located along the ionization fronts. We illustrate that  \\spitzer\\ infrared images  provide new, excellent tools for studying massive-star formation in ways  that were not previously available. ",
        "conclusions": "The detection of multiple mid-infrared protostars from the cold dust cores is a new result. How does the mid-infrared emission leak through  thick envelopes of material infalling onto the stars? Isn't the mid-infrared emission absorbed by the cold and  thick envelopes?  We can address where the emission arises, how the dust is heated, and how the emission escapes cold cores.  The cold envelope sizes of the TC3, TC4, and TC4b cores are  0.2, 0.2, and 0.16 pc, respectively , which were  determined  by \\citet{lef00}, using the 50\\% contour in the millimeter flux map. The estimated optical depth is: \\begin{equation} \\tau_{1250} = {F_{1250} \\over {\\Omega \\, B_{\\nu}(T)}} = 1.08\\times 10^{-4}  F_{1250} {d_{1.68kpc}^2 \\over  R^{2} } \\label{eq1} \\end{equation}  \\noindent where F$_{1250}$  is the mm-wave flux, $\\Omega$ is a solid angle of  the source size,   B$_{\\nu}$(T) is a black body Planck function at a temperature of T,  and R is a source radius in unit of pc.  The optical depths of the TC3, TC4, and TC4b cores are between  0.7-2.5 $\\times 10^{-3}$ (for a temperature of 22 K). The corresponding optical depths at 8$\\mu$m were determined using $\\tau_{8}$ = $\\tau_{1250} \\times (1250/8)^{\\alpha_1}$ $(8/100)^{\\alpha_2}$, where $\\alpha_1=2$ is the extinction slope between 100-3000$\\mu$m and $\\alpha_2=1$  is the extinction slope between 10-100$\\mu$m (Li \\& Draine 2001); this yields $\\tau_{8 \\mu m}$ $\\sim$ 1--7. The corresponding visual extinction was determined using A$_{8{\\mu}m}$/A$_K$ = 0.5 (Draine 2003); this yields a total extinction of A$_v$  $\\sim$ 20--150.   Despite the high extinction of the envelope and the high optical depth of the protostars (estimated using a temperature of 22K) at 8$\\mu$m, apparently we observed mid-IR emission from the early protostars. We note that the optical depths of 150 K and 400 K components are much smaller, more than one magnitude smaller than that of the cold (22 K) envelope. From the SED fits (see  Fig. \\ref{sedtc4a}) we estimated that  the optical depth of the 8$\\mu$m emitting region is $< 10^{-3}$ for the warm and hot temperature components, which suggests that the accretion region is probably optically thin if it is larger than $\\sim 10$ AU. In addition, the \\spitzer\\ images revealed that the infrared sources of the protostars were point-sources, inferring that the regions emitting 400 K and 150 K would be smaller than  the PSF radii of the IRAC and MIPS 24$\\mu$m images, which are 0.007 pc (= 1500 AU) and 0.025 pc (= 5$\\times 10^3$ AU), respectively. %which are 0.014 pc (=3000 AU) and 0.05 pc, respectively. This is consistent with the models of the massive protostars \\citep{oso99}; the cold component (responsible for far-infrared and 1300$\\mu$m emission) would be emitted from a region $>$ $10^4$ AU  from the  central source, while the 400 K component (responsible for mid-infrared emission) would be  from $<$ 1500 AU region. The mid-infrared emission can escape the dense envelope; it suffers a modest mid-infrared extinction  but traces the emission directly from the accretion region.  Previous ISO observations showed similarly  mid-infrared counterparts of a few low mass Class 0 objects in other star forming regions \\citep{cer00}. % although the ISO spatial resolution was limited. New \\spitzer\\ observations of Class 0 objects also detected mid-infrared protostars in L1014 (Young et al. 2004) and  Cepheus E \\citep{nor04}. The detections infer that  the region of hot core in early protostars is  optical thin at  the mid-infrared regime, and the size of hot core is smaller than the cold envelope.   The mid-infrared emission is powered by accretion of envelope material onto the protostar, and  most of the bolometric luminosity at the early stage of the protostar is from accretion. The accretion rate is given by \\cite{mck02,mck03}:   $$\\dot{m_{*}}= 4.75\\times10^{-4}\\,\\, \\epsilon_{core}^{1/4} (f_{gas} \\,\\phi_P \\,\\alpha_{vir})^{3/8} \\,\\,  \\biggl({m_{*f} \\over {30 M_\\odot}}\\biggr)^{3/4} \\Sigma_{cl}^{3/4} \\biggl({m_{*} \\over m_{*f} }\\biggr)^{1/2} \\,\\, M_{\\odot} \\, yr^{-1} $$ \\begin{equation} \\sim 4.6\\times 10^{-4}  \\biggl({m_{*f} \\over {30 M_\\odot}}\\biggr)^{3/4} \\Sigma_{cl}^{3/4}  \\biggl({m_{*} \\over m_{*f} }\\biggr)^{1/2} \\,\\, M_{\\odot} \\, yr^{-1} \\label{eq2} \\end{equation}  where $\\epsilon_{core}$ is the fraction of the total core mass (M$_{core}$),   $\\phi_P$ is the ratio of the core's surface pressure to the mean pressure in the clumps, $f_{gas}$ is the fraction of the cloud's mass that is in gas, as opposed to stars, $\\Sigma_{cl}$ is a mean mass column density, $m_{*}$ (=$\\epsilon_{core}$M$_{core}$, where M$_{core}$ is the total core mass) is  the instantaneous mass (the  current mass), and $m_{*f}$ is the final mass of the star. We used the values of $\\Sigma_{cl}$ =1, $\\epsilon_{core}$=0.5  and $\\phi_P$=0.663 (see \\cite{mck02} for details). The bolometric luminosity is the sum of the internal and accretion luminosity (L$_{bol}$ = L$_{int}$ + L$_{acc}$). The accretion luminosity is given by \\cite{mck03}: \\begin{equation} L_{acc} \\sim 3.0\\times10^4 \\,\\, \\biggl({f_{acc} \\over 0.5}\\biggr) \\biggl({m_{*f} \\over {30 M_\\odot}}\\biggr)^{1.2}  \\Sigma_{cl}^{3/4} \\biggl({m_{*} \\over m_{*f} }\\biggr)^{0.95} \\,\\, \\Sigma_{cl}^{3/4}  \\,\\, L_{\\odot} \\end{equation} where ${f_{acc}}$ ($<<$1) is a factor that accounts for the energy advected into the star or used to drive protostellar outflows.  The internal luminosity, L$_{int}$, is equal to the luminosity transported by radiation and is determined mostly by the stellar mass (L$_{int}$ $\\propto$ M$_*^{5.5}$ R$_*^{-0.5}$ $\\propto$ M$_{*}$). A simple formula cannot be used for a wide range of masses of stars, because as the mass increases, the relative contribution of the Kramers and electron scattering opacities changes \\citep{nak00}. Therefore, in order to estimate the mass of star from the bolometric luminosity, we directly used  Figure 2 of \\citet{mck02}. % and Figure 6 of \\citet{mck03}. The instantaneous masses of TC3A and TC4A are 3-8 M$_{\\odot}$ and 2-5 M$_{\\odot}$ for the bolometric luminosities of $\\sim$ 1700 and 1200 L$_{\\odot}$, respectively, assuming a significant amount of the 1.3mm flux is from Component A.  The accreting properties of the TC3 core are similar to those of G34.24+0.1 \\citep{mck03}.  We estimated accretion rates using the equation 2 for a grid of final masses of 7.5, 30 and 120 M$_\\odot$.  The accretion rates for TC3 and TC4 are 1-5$\\times$10$^{-4}$ M$_{\\odot}$ yr$^{-1}$ and 0.9-4$\\times$10$^{-4}$ M$_{\\odot}$ yr$^{-1}$, respectively. We also estimated the current masses and the accretion rates of TC1 and TC2 using L$_{bol}$ of 600 L$_{\\odot}$; the current masses are 1-3 M$_{\\odot}$ and the  accretion rates are 0.8-1.7$\\times$10$^{-4}$ M$_{\\odot}$ yr$^{-1}$. The estimated accretion luminosities infer that high portions of the bolometric luminosities are from accretion luminosities in the early stage of massive-star formation. The timescale of the infalling stage is determined mainly by  the conditions in the stars' natal cloud and weakly depends on the mass of stars.     In the study of massive-star formation, a fundamental open question is how clusters are formed.  Do massive dense cores have internal substructures? Do clumps evolve independently to produce stars, or do they share a common evolutionary process?  Are apparent-single stars born single, or are they born in groups and subsequently ejected? While low-mass stars are believed to be produced mainly through accretion (Lada 1991), there are  two main scenarios to explain  high-mass star formation.  One scenario is through accretion,  like that of low-mass stars but with higher accretion rates; the other scenario is through formation of  high-mass stars through coalescence of lower-mass protostars. We found  5 protostars for each of TC3 and TC4 cores, and the brightest and the most massive star in each cluster is located in the dust continuum peak (see Table 2), implying these systems are possible protoclusters. The brightest protostar appears near the center of the dust continuum peak. The brightest 24$\\mu$m source in the TC3, TC3A, is located in the millimeter peak and at the center of the core.  Within the cluster, as the separation of TC3B-TC3C from the largest star of each Component A increases, the mid-infrared luminosity decreases in the case of the  TC3 core, as shown in Table 2. This suggests that the brightest and most massive star among 4 protostars  is  located at the bottom of the gravitational potential well and is located where it formed,  although the mid-infrared luminosities may not be directly correlated with the total bolometric luminosities.  Accretion from gas or nearby  fragmented clumps is found through the accretion of residual gas onto relatively low-mass cores. This picture may be consistent with  the mass segregation, in young stellar clusters,  and competitive accretion proposed by Bonnell \\& Davies (1998) and Bonnell et al. (2004). However, the brightest 24$\\mu$m source in the TC4 core, TC4A, is slightly off from the center of the millimeter peak, although TC4A is still the largest star in the cluster and the closest star to the dust continuum peak. %As the separation %of TC4B from the largest star, Component A, increases, The mid-infrared luminosity of TC4B is also smaller than that of the TC4A. %the star does not directly decrease in the case of the core of %TC4. It is unknown how these mid-infrared luminosities correlate with the total bolometric luminosities for individual sources. Comparable resolution far-infrared or millimeter observations are required in order to answer this question. Unfortunately, even through the Trifid Nebula is one of the nearest massive-star-forming regions with a distance of only 1.68 kpc, protostars in the Trifid as well as other comparable massive-star-forming regions are not bright enough for such observatories as the submillimeter Array (SMA). Future far-infrared observations with comparable spatial resolution are needed to answer this question."
    },
    "0601/astro-ph0601205_arXiv.txt": {
        "abstract": "{We present optical and near-infrared observations with Keck of the binary millisecond pulsar PSR~J0751+1807. We detect a faint, red object -- with $R=25.08\\pm0.07$, $B-R=2.5\\pm0.3$, and $R-I=0.90\\pm0.10$ -- at the celestial position of the pulsar and argue that it is the white dwarf companion of the pulsar. The colours are the reddest among all known white dwarfs, and indicate a very low temperature, $T_\\mathrm{eff}\\approx4000$\\,K. This implies that the white dwarf cannot have the relatively thick hydrogen envelope that is expected on evolutionary grounds. Our observations pose two puzzles. First, while the atmosphere was expected to be pure hydrogen, the colours are inconsistent with this composition. Second, given the low temperature, irradiation by the pulsar should be important, but we see no evidence for it. We discuss possible solutions to these puzzles. ",
        "introduction": "\\label{sec:intro} Among the pulsars in binaries, the largest group, the low-mass binary pulsars, has low-mass white-dwarf companions. Before the companions became white dwarfs, their progenitors filled their Roche lobe and mass was transferred to the neutron stars, thereby spinning them up and decreasing their magnetic fields.  Considerations of the end of this stage, where the white dwarf progenitor's envelope becomes too tenuous to be supported further, allow one to make predictions for relations between the orbital period and white dwarf mass, and orbital period and eccentricity (for a review, e.g., \\citealt{pk94,sta04}). Furthermore, after the cessation of mass transfer, two clocks will start ticking at the same time: the neutron star, now visible as a millisecond pulsar, will spin down, while the secondary will contract to a white dwarf and start to cool.  Consequently, the spin-down age of the pulsar should equal the cooling age of the white dwarf.  From optical observations of white-dwarf companions to millisecond pulsars one can estimate the white-dwarf cooling age and compare it with the pulsar spin-down age. Initial attempts to do this \\citep{hp98a,hp98b,sdb00} revealed a dichotomy in the cooling properties of white dwarfs in the sense that some white dwarf companions to older pulsars have cooled less than those of younger pulsars. In particular, the companions of PSR~J0437$-$4715 \\citep{dbv93,sbb+01} and PSR~B1855+09 \\citep{kbkk00,rt91} have temperatures of about 4000--5000\\,K, with characteristic pulsar ages of 5\\,Gyr. This is in contrast to the companion of PSR~J1012+5307 \\citep{llfn95,kbk96,cgk98}, which has a higher temperature (8600\\,K), while it orbits an older pulsar (8.9\\,Gyr).  A likely cause for this dichotomy is the difference in the thickness of the envelope of hydrogen surrounding the helium core of the white dwarf \\citep{ashp96}. After the cessation of mass transfer, the white dwarfs have relatively thick ($\\sim\\!10^{-2}$\\,M$_\\odot$) hydrogen envelopes which are able to sustain residual hydrogen shell-burning, keeping the white dwarf hot and thereby slowing the cooling \\citep{dsbh98}.  The shell burning, however, can become unstable and lead to thermal flashes which can reduce the mass of the envelope. White dwarfs with such reduced, relatively thin ($\\la10^{-3}$\\,M$_\\odot$) hydrogen envelopes cannot burn hydrogen and, as a result, cool faster. The transition between thick and thin hydrogen envelopes was predicted to lie near 0.18--0.20\\,M$_\\odot$ (where heavier white dwarfs have thin envelopes; \\citealt{ashp96,seg00,asb01}).  Until recently, PSR~J1012+5307, with an orbital period $P_\\mathrm{b}=0.60$\\,d, was the only system for which a thick hydrogen envelope was required to match the two timescales. Given the relation between the white dwarf mass and the orbital period \\citep{jrl87,rpj+95,ts99}, companions in similar or closer orbits should have similar or lower mass, and thus have thick hydrogen envelopes as well. This was confirmed by the recent discovery of two new, nearby, binary millisecond pulsars with orbital periods similar to that of PSR~J1012+5307; PSR~J1909$-$3744 (1.53\\,d, \\citealt{jhb+05}) and PSR~J1738+0333 (0.354\\,d, Jacoby et al., in prep.; see \\citealt{kbjj05} for preliminary results). For both, the temperatures and characteristic ages are similar to those of PSR J1012+5307, and thus one is led to the same need for a thick hydrogen envelope. These discoveries, combined with the thin envelopes inferred for PSR~J0034$-$0534 (1.59\\,d) and binaries with longer periods, suggest that the transition occurs at a mass that corresponds to an orbital period just over 1.5\\,d (\\citealt{kbjj05}). All systems with shorter orbital periods should have thick hydrogen envelopes.  The two known millisecond pulsars with white dwarf companions that have shorter orbital periods than PSR~J1012+5307 but do not have optical counterparts, are PSR~J0613$-$0200, with a 1.20\\,d period, and PSR~J0751+1807, which has the shortest orbital period of all binary millisecond pulsars with $M_\\mathrm{c}>0.1$\\,M$_\\odot$ companions, 0.26\\,d \\citep{lzc95}. The latter system is of particular interest because the companion mass has been determined from pulse timing ($M_\\mathrm{WD}=0.19\\pm0.03$\\,M$_\\odot$ at 95\\% confidence; \\citealt{nss+05}), so that one does not have to rely on the theoretical period-mass relationship. Intriguingly, for PSR J0751+1807, optical observations from \\citet{lcf+96} set a limit to the temperature of 9000\\,K, which is only marginally consistent with it having a thick hydrogen envelope. Based on this, \\citet{esa01}, suggested the hydrogen envelope may have been partially lost due to irradiation by the pulsar.  The faintness of the companion to PSR~J0751+1807 aroused our curiosity and motivated us to obtain deep observations to test the theoretical ideas discussed above. We describe our observations in Sect.~\\ref{sec:observations}, and use these to determine the temperature, radius and cooling history in Sect.~\\ref{sec:tandr}. In Sect.~\\ref{sec:irradiation}, we investigate irradiation by the pulsar, finding a surprising lack of evidence for any heating. We discuss our results in Sect.~\\ref{sec:discussion}. ",
        "conclusions": "\\label{sec:discussion} We have optically identified the white dwarf companion of the binary millisecond pulsar PSR~J0751+1807. We find that the companion has the reddest colours of all known millisecond pulsar companions and white dwarfs. These colours indicate that the companion has a very low (ultra-cool) temperature of $T_\\mathrm{eff}\\sim\\!3500-4300$\\,K. Furthermore, the colours suggest that the white dwarf has a pure helium atmosphere, or a helium atmosphere with some hydrogen mixed in, as invoked for the field white dwarf WD~0346+246 which has similar colours \\citep{osh+01,ber01}.  Our observations are inconsistent with evolutionary models, from which one would expect a pure hydrogen atmosphere. Indeed, as for other short-period systems, the hydrogen envelope is expected to be thick enough to sustain significant residual hydrogen burning, leading to temperatures far in excess of those observed. It may be that the mass of the envelope was reduced due to shell flashes or irradiation by the pulsar, as was proposed by \\citet{esa01}.  However, we see no evidence for irradiation, despite the fact that the pulsar spin-down flux inpinging on the white dwarf is roughly double the observed thermal flux. Clues to what happens might be found from more detailed studies of the spectral energy distribution, or more accurate phase-resolved photometry.  Finally, a deeper observation at infrared wavelengths would allow one to distinguish between the different atmosphere compositions for the companion: for a pure helium atmosphere, black-body like colours are expected, while if any hydrogen is present, the infrared flux would be strongly depressed (as is seen for WD 0346+246). With adaptive optics instruments, such observations should be feasible."
    },
    "0601/astro-ph0601496_arXiv.txt": {
        "abstract": "{ Some common properties of helical magnetic fields in decaying and driven turbulence are discussed. These include mainly the inverse cascade that produces fields on progressively larger scales. Magnetic helicity also restricts the evolution of the large-scale field: the field decays less rapidly than a non-helical field, but it also saturates more slowly, i.e.\\ on a resistive time scale if there are no magnetic helicity fluxes. The former effect is utilized in primordial field scenarios, while the latter is important for successfully explaining astrophysical dynamos that saturate faster than resistively. Dynamo action is argued to be important not only in the galactic dynamo, but also in accretion discs in active galactic nuclei and around protostars, both of which contribute to producing a strong enough seed magnetic field. Although primordial magnetic fields may be too weak to compete with these astrophysical mechanisms, such fields could perhaps still be important in producing polarization effects in the cosmic background radiation. ",
        "introduction": "Magnetic helicity plays a fundamental role both in primordial theories of galactic magnetism as well as in dynamo theories amplifying and sustaining contemporary galactic fields. Both issues have been reviewed in recent years (Grasso \\& Rubinstein 2001; Widrow 2002; Giovannini 2004; Brandenburg \\& Subramanian 2005a). We will therefore only try to collect the main points relevant to the issues concerning magnetic helicity in galactic and protogalactic magnetism.  The main reason magnetic helicity is at all of concern to us is that even in the resistive case the rate of magnetic helicity dissipation asymptotes to zero as the magnetic Reynolds number goes to infinity. This is not the case with magnetic energy dissipation, which remains always important, and does not decrease with increasing magnetic Reynolds number (Galsgaard \\& Nordlund 1996). Therefore the magnetic helicity is nearly conserved at all times. This has serious consequences for the evolution of magnetohydrodynamic (MHD) turbulence, as has been demonstrated by a number of recent simulations when the resolution has been large enough (Brandenburg 2001a; Mininni et al.\\ 2005).  At a more descriptive level, magnetic helicity characterizes the degree of field line linkage. As the magnetic field relaxes, its energy decreases, but the linkage stays, at least as much as possible. The field's inability to undo its knots implies also that the field cannot decay freely. This slows down the decay, which is important if a primordial field is to be of any significance at the time of recombination. In the driven case, on the other hand, magnetic helicity is better pictured in terms of writhe and twist (e.g.\\ Longcope \\& Klapper 1997; D\\'emoulin et al.\\ 2002). Writhe refers to the tilt of a flux tube, and we use both expressions synonymously. A cyclonic event tilts individual flux tubes, but as it does so, a corresponding amount of internal twist is necessarily introduced in the tube (Blackman \\& Brandenburg 2003). This is what saturates the dynamo, and this can be a very powerful effect if the small-scale internal twist cannot escape. In this review we discuss both decaying and driven turbulence. The former is relevant for prolonging the decay of a primordial field, while the latter is relevant for understanding how the galactic dynamo saturates and how to enable it to do so faster. ",
        "conclusions": "Not all magnetic fields will be helical, but if they are, this can have dramatic consequences for their evolution. The effects can be equally dramatic both in decaying and in driven turbulence, as has been highlighted in this review. Although we have not discussed this in the present paper, it should be emphasized that helical large-scale magnetic fields can also be generated in non-stratified shear flows where there is no $\\alpha$ effect, but there can instead be the so-called shear--current of $\\meanWW\\times\\meanJJ$ effect (Rogachevskii \\& Kleeorin 2003, 2004). This effect may also explain the large-scale dynamo action seen in \\Fig{pmean_comp}, where the results without helicity are quite similar to those with helicity (Brandenburg 2005). One-dimensional mean field calculations with the $\\meanWW\\times\\meanJJ$ effect (Brandenburg \\& Subramanian 2005c) show that in this case a magnetic $\\alpha$ effect can be produced that has different signs on the two sides of the midplane. This magnetic $\\alpha$ effect thus contributes to the saturation of the dynamo even if there is no ordinary (kinetic) $\\alpha$ effect. This highlights once more the dramatic effects played by magnetic helicity.  Whether or not the primordial magnetic field was really helical remains a big question. If it was, it is likely that an inverse cascade process has produced fields of progressively larger scale. This might lead to observable effects in the cosmic microwave background. Such a field may also be important for seeding the galactic dynamo, but it is important to realize that a variety of astrophysical mechanisms may also produce seed fields just as large. Our estimate for magnetized outflows from AGNs or YSOs assumes that the source remains active for a certain period of time, and that their exhaust goes freely into the ambient medium. Partial evidence for this actually happening lies in the fact that clusters of galaxies are chemically enriched with heavier elements. Given that magnetic fields are intrinsically connected with the outflow, just like the heavier elements in it, it is quite plausible that some degree of magnetic contamination of the cluster must have occurred.  In order to produce finally the observed large-scale magnetic fields of galaxies, some more reshaping, amplification, and maintenance against magnetic decay is necessary. Roughly, we expect this to happen just like the mean field dynamo is able to amplify and maintain the field, although it must operate on an already strong enough field. This initial field will still be random and of mixed parity about the midplane (or equator), but there will be some finite degree of quadrupolar field which is the one that is dominant in many galaxies; see Krause \\& Beck (1998) and Brandenburg \\& Urpin (1998) for a related discussion about the importance of seeding the quadrupolar field component. As we have argued above, the catastrophic quenching problem of the dynamo has to be overcome, and this is likely to be the case because of various magnetic and current helicity fluxes operating within the entire dynamo domain. In the context of the solar dynamo, simulations have now begun to demonstrate the dramatic difference made by open boundary conditions, and we hope that a similar demonstration will soon be possible for the galactic dynamo as well. Corresponding mean field calculations have already been performed showing that the catastrophic quenching effect is overcome by an advective flux out of the domain along the vertical direction. In particular, it will be interesting to see whether the shedding of magnetic helicity can actually lead to directly observable effects."
    },
    "0601/astro-ph0601175_arXiv.txt": {
        "abstract": "{We present results from  Chandra observations of the 3C/FR~I sample  of low  luminosity  radio-galaxies.  We  detected a  power-law nuclear component in 12 objects out of the 18 with available data.  In 4 galaxies  we detected  nuclear X-ray absorption  at a level  of $N_H \\sim (0.2-6)  \\times 10^{22}$  cm$^{-2}$.  X-ray absorbed  sources are associated with the  presence of highly inclined dusty  disks (or dust filaments projected  onto the  nuclei) seen in  the HST  images.  This suggests  the existence  of a  flattened X-ray  absorber, but  of much lower optical depth than in classical obscuring tori.  We thus have an un-obstructed view toward most FR~I nuclei while absorption plays only a marginal role in the remaining objects.  Three pieces  of evidence support  an interpretation for a  jet origin for the  X-ray cores: i)  the presence of strong  correlations between the  nuclear  luminosities in  the  radio,  optical  and X-ray  bands, extending  over  4  orders  of  magnitude  and  with  a  much  smaller dispersion ($\\sim$ 0.3 dex) when  compared to similar trends found for other classes of AGNs, pointing to a common origin for the emission in the three bands;  ii) the close similarity of  the broad-band spectral indices with the sub-class of BL Lac objects sharing the same range of extended  radio-luminosity, in  accord with  the  FR~I/BL~Lacs unified model; iii) the presence of  a common luminosity evolution of spectral indices in both FR~I and BL~Lacs.  The low luminosities of the  X-ray nuclei, regardless of their origin, strengthens  the interpretation  of  low efficiency  accretion in  low luminosity radio-galaxies. ",
        "introduction": "By exploring the properties of low luminosity radio-galaxies (LLRG) it is possible to  extend the study of active  galactic nuclei (AGN), and in  particular of  the  radio-loud sub-population,  toward the  lowest level of  nuclear luminosity. This  offers the opportunity  to improve our  understanding of  the mechanism  of accretion  onto super-massive black holes  and of their radiative manifestations.   In this respect, it  is important  to note  that LLRG  not only  represent the  bulk of radio-loud active  galaxies, due to the steepness  of their luminosity function,  but also  a substantial  fraction of  the  overall galaxies population. In fact LLRG are associated to a fraction as high as 40 \\% of     all    bright     elliptical     and    lenticular     galaxies \\citep{sadler89,auriemma77}.   Thus  the   presence  of  a  low  power radio-galaxy  represents  by  far  the most  common  manifestation  of nuclear activity in early type galaxies.  Low  luminosity AGN  (LLAGN) also  represent a  link between  the high luminosity AGN and the population  of quiescent galaxies, as it is now widely recognized that most (if not all) galaxies host a super-massive black-hole   \\citep[e.g.][]{kormendy95}.    The   comparison  of   the different  manifestations  of  nuclear  activity  across  the  largest possible  range of luminosity  is then  a crucial  step to  unveil the connections  (and  the  diversities)  between  active  and  non-active galaxies.  Unfortunately, the  study of LLRG  has been significantly  hampered by the contamination  from host galaxy emission which,  in most observing bands, dominates the emission from the AGN.  To constrain the physical processes  at  work  in  these  objects it  is  clearly  necessary  to disentangle   the   AGN   and   host's  contributions   via   spectral decomposition  or high  resolution imaging.   In the  last  few years, thanks  in particular  to the  Hubble Space  Telescope (HST)  this has become routinely possible.  An  example  of the  insights  that  can  be obtained  following  this approach  comes  from  the   results  derived  from  the  analysis  of broad-band HST imaging of  a sample of LLRG \\citep{chiaberge:ccc}.  In the  majority of  the targets,  HST  images revealed  the presence  of unresolved optical  nuclei.  Fluxes and luminosities  of these sources show a  tight correlation  with the radio  cores, extending  over four orders  of magnitude.  This  has been  interpreted as  being due  to a common  non-thermal emission process  in the  radio and  optical band, i.e.   that we  are seeing  the optical  emission from  the base  of a relativistic jet.  The low luminosity  of LLRG optical nuclei, and the possible dominance  of the  emission related to  out-flowing material, indicates a low  level of accretion and/or a  low radiative efficiency with  respect to  classical,  more luminous  AGN.   The high  fraction ($\\sim$  85  \\%) of  objects  with  detected  optical nuclear  sources suggests  a  general  lack  of  obscuring molecular  tori,  a  further distinction with respect  to other classes of AGN.   The tenuous torus structure, when  considered together with the low  accretion rate, the low mass of  the compact emission line regions  (10 - 10$^3 M_{\\sun}$) and  the limits  to the  mass  of the  Broad Line  Region ($M_{BLR}  < 10^{-2} M_{\\sun}$),  indicates that  a general paucity  of gas  in the innermost  regions   of  LLRG  emerges  as   the  main  characterizing difference from more powerful AGN \\citep{capetti:cccriga}.  The  advent of  Chandra provides  us  with the  unique opportunity  to extend the  study of  LLRG to  the X-ray band  since its  high spatial resolution  enables  us  to  isolate  any  low  power  nuclear  source associated to  LLRG.  Furthermore, with  respect to the  optical data, the spectral  capabilities of Chandra  allow us to study  directly the spectral behaviour of  the X-ray nuclei. This can  be used to quantify the effects  of local absorption  as well as  the slope of  their high energy emission as to better  constrain the emission processes at work in  these sources,  particularly when  combined in  a multi-wavelength analysis using also optical and radio observations.  This approach was already    followed     by    several    authors     in    the    past \\citep[e.g.][]{fr1sed,hardcastle00,trussoni03,hardcastle03}         but clearly, the  superior capabilities  of Chandra warrant  to re-explore this issue in much greater depth.  Here we present results obtained from the analysis of Chandra data for the  sample of  LLRG  drawn  from the  3C  catalogue of  radio-sources \\citep{bennett62},  more  specifically   those  with  a  morphological classification as FR~I  \\citep{fanaroff74}.  Considering LLRG from the 3C  sample, the  same studied  by \\citet{chiaberge:ccc}  represents an obvious choice. It  is the best studied sample  of radio-loud galaxies in existence with a vast suite of ground and spaced based observations for  comparison  at  essentially  all accessible  wavelengths.   Other samples  of LLRG  have been  considered  in the  literature, but,  not surprisingly, the coverage provided  by Chandra observations for these samples is  severely incomplete. For example,  only about 10  \\% of B2 sample of LLRG  \\citep{colla75} has been observed with  Chandra (to be compared to the $\\sim$ 55 \\% for the 3C/FR~I sample), and the analysis is further  obstructed by the  unknown selection biases  introduced in the choice of the individual targets.  Similarly we preferred to focus on the  Chandra data alone, to  provide the highest  possible level of uniformity in the analysis.  The paper  is organized as follows:  in Sec. \\ref{obs}  we present the observations and the data reduction that lead to the results described in Sec.   \\ref{results} and  in particular to  the detection  of X-ray nuclei in most FR~I.  In Sec.  \\ref{absorption} we explore the effects of the nuclear  X-ray absorption, by also relating  it to the presence of  dust features  seen in  the HST  images. The  origin of  the X-ray nuclei found in  LLRG is explored in Sec.  \\ref{jet}.  The results are discussed and summarized in Sec. \\ref{summary}.  We adopted $\\rm{H}_o=75 $ km s$^{-1} $Mpc$^{-1}$ and $q_0=0.5$. ",
        "conclusions": ""
    },
    "0601/astro-ph0601669_arXiv.txt": {
        "abstract": " ",
        "introduction": "The structure and evolution of dark matter halos is widely discussed in the recent years (e.g. \\cite{BM}) A large number of papers has been devoted to the investigation of spatial density profiles of dark matter halos, $\\rho(r)$. both numerical \\cite{Navarro, Moore05} and analytical \\cite{LMKW}. Recently several papers appeared that described general properties of solutions of Jeans equation \\cite{DM,AWBBD}. The Jeans equation for a sperically symmetric self-gravitating system reads as follows: \\begin{equation}  \\label{Jeans} \\frac{d}{dr} (\\rho\\,{\\sigma_r}^2) + \\frac{2\\beta}{r} (\\rho\\,{\\sigma_r}^2) + \\rho\\frac{G\\,M(r)}{r^2} = 0 \\,. \\end{equation} Here ${\\sigma_r}^2$ and ${\\sigma_t}^2$ are the radial and tangential velocity dispersion, $\\beta = 1-\\frac{{\\sigma_t}^2}{2\\,{\\sigma_r}^2}$ is the Binney's anisotropy parameter (systems with $0<\\beta<1$ are radially anisotropic, with $\\beta<0$ -- tangentially anisotropic, and $\\beta=0$ is the isotropic case).  An additional assumption is the following empirical property of dark matter halos: the generalized phase-space-{\\it like} density of dark matter particles in a halo is a power-law in radius: \\begin{equation}  \\label{gpsd} \\rho(r)/{\\sigma_r}^\\epsilon(r) \\propto r^{-\\alpha}  \\,, \\end{equation} which is established both in N-body calculations \\cite{TN} and some analytical models \\cite{AWBBD}.  For a closed set of equations one also needs to specify $\\beta$. Here two different approaches may be adopted: in \\cite{AWBBD} isotropy was assumed ($\\beta=0$), and in \\cite{DM} the following relationship between $\\beta$ and the logarithmic slope of density profile $\\gamma = d \\ln\\rho(r)/d r$ was taken: \\begin{equation}  \\label{betagammarelation} \\beta = \\beta_0 + b(\\gamma-\\gamma_0)  \\,. \\end{equation} This kind of relation has been proposed in \\cite{HM} based on empirical analysis of different numerically treated situations.  These investigations demonstrated that under these assumptions there exist only few possibilities for a self-consistent halo model. The most realistic is the set of models with density profiles having inner and outer power-law asymptotes (so-called Zhao $\\alpha\\beta\\gamma$ models \\cite{Zhao}).  However, this approach, being very fruitful, does not describe explicitly the distribution function. It operates only with its second moment with respect to velocity, i.e., the velocity dispersion. The most complete description of halo structure can be made in terms of phase-space distribution function $f(r,v)$. One attempt to deal with full velocity distribution function (VDF) was made in \\cite{Tsallis}, where an Eddington inversion formula was applied to obtain VDF for power-law density profiles. However, this study was limited to isotropic case.  An approach based on certain physical assumptions regarding the form of the distribution function seems more promising. It avoids postulating empirical assumptions like (\\ref{gpsd}, \\ref{betagammarelation}), which, nevertheless, still hold in many cases. It is convenient to express the distribution function in terms of adiabatic invariants during slow evolution of the gravitational potential. This allows to describe easily the process of adiabatic contraction of dark matter halo caused by condensation of baryonic matter in its centre, which occurs during the formation of a galaxy. This approach is adopted in the present paper.  The paper is organized as follows. In the second section we describe two models of inner halo structure and explain physical arguments for the underlying assumptions concerning the dark matter distribution function. In the third section we calculate the velocity anisotropy coefficient. Furthermore, in the fourth section we consider the halo response to the adiabatic compression and compare it with other studies. Finally, the conclusions are presented. ",
        "conclusions": " \\begin{enumerate} \\item The more radially biased is the velocity, the less compressed is the halo. \\item The shallower is initial density profile (less value of $\\gamma$), the less is the degree of compression compared to that of Blumenthal's method. Standard model of adiabatic contraction systematically overestimates the effect of contraction. \\item The model predicts a moderate enhancement (2 to 4 times) of dark matter mass in the bulge. The shallower is the initial profile, the greater the increase of dark matter mass. This result significantly increases the estimates of possible annihilation radiation flux from the center of the Galaxy \\cite{GS, Merritt, ZV}. \\end{enumerate}  I am grateful to Maxim Zelnikov for helpful discussion on the topic of the paper, to Steen Hansen for pointing out some interesting issues and to anonymous referee for valuable comments and proposals. This work was supported by Landau Foundation (For\\-schungs\\-zent\\-rum J\\\"ulich) and Russian Fund for Basic Research (project nos. 01-02-17829, 03-02-06745)."
    },
    "0601/astro-ph0601343_arXiv.txt": {
        "abstract": "{We solve a weakly singular integral equation by Laplace transformation over a finite interval of $\\mathbb{R}$. The equation is transformed into a Cauchy integral equation, whose resolution amounts to solving two Fredholm integral equations of the second kind with regular kernels. This classical scheme is used to clarify the emergence of the auxiliary functions expressing the solution of the problem. There are four such functions, two of them being classical ones. This problem is encountered while studying the propagation of light in strongly scattering media such as stellar atmospheres.\\\\   Key words: Weakly singular integral equation, finite Laplace transform, sectionally analytic function, radiation transfer theory, stellar atmospheres. ",
        "introduction": "The integral form of the equation describing the radiative transfer of energy in a static, plane-parallel stellar atmosphere is \\cite{ahuesetal2002a} \\begin{equation} \\label{eq1} S(a,b,\\tau)=S_0(a,b,\\tau)+a\\int_{0}^{b}K(\\tau-\\tau')S(a,b,\\tau')d \\tau', \\end{equation} where $S$ is the source function of the radiation field and $S_0$ describes the radiation of the primary (internal or external) sources. These functions depend on the two parameters of the problem: the albedo $a\\in]0,1[$, which characterizes the scattering properties of the stellar plasma, and the optical thickness $b>0$ of the atmosphere. They also depend on the optical depth $\\tau\\in[0,b]$, which is the space variable. Equation (\\ref{eq1}) means that the radiation field at level $\\tau$ is the sum of the {\\it{direct}} field from the primary sources, and the {\\it{diffuse}} field having scattered at least once.\\\\ In the simplest scattering process conceivable - i.e., a monochromatic and isotropic one - the kernel of the integral equation (\\ref{eq1}) is the function \\begin{equation} \\label{eq2} K(\\tau):\\,=\\frac{1}{2}E_1(\\vert \\tau \\vert)\\quad(\\tau\\in \\mathbb{R}^*), \\end{equation} where $E_1$ is the first exponential integral function \\begin{equation} \\label{eq3} E_1(\\tau):\\,=\\int_{0}^{1}\\exp(-\\tau/x)\\frac{dx}{x}\\quad(\\tau > 0). \\end{equation} Since $E_1(\\vert \\tau \\vert) \\sim - \\ln(\\vert \\tau \\vert)$ when $\\vert \\tau \\vert \\to 0^+$, the kernel of the integral equation (\\ref{eq1}) is weakly singular on its diagonal. The free term $S_0$ includes the thermal emission of the stellar plasma - of the form $(1-a)B^*(\\tau)$, where $B^*$ is a known function - and the contribution $aJ_0^{ext}(b,\\tau)$ of the external sources via the boundary conditions: see the introduction of \\cite{ahuesetal2002a}. Hence $S_0(a,b,\\tau)=(1-a)B^*(\\tau)+aJ_0^{ext}(b,\\tau)$. In a homogeneous and isothermal atmosphere assumed to be in local thermodynamic equilibrium, the function $B^*$ coincides with the Planck function $B(T)$ at the (constant) temperature $T$. Moreover, $J_0^{ext}=0$ in the absence of external sources and thus \\begin{equation} \\label{eq4} S_0(a,b,\\tau)=(1-a)B(T), \\end{equation} which shows that $S_0$ is independent of $\\tau$ in this model. The solution $S$ to the problem (\\ref{eq1}) is then \\begin{equation} \\label{eq5} S(a,b,\\tau)=(1-a)B(T)Q(a,b,\\tau), \\end{equation} where $Q$ solves the following integral equation: \\begin{equation} \\label{eq6} Q(a,b,\\tau)=1+a\\int_{0}^{b}K(\\tau-\\tau')Q(a,b,\\tau')d \\tau'. \\end{equation} It is proved in \\cite{ahuesetal2002b} that the space $C^0([0,b])$ of the continuous functions from $[0,b]$ to $\\mathbb{R}$ is invariant under the operator \\begin{equation} \\label{eq7} \\Lambda\\;:\\;f\\;\\;\\rightarrow\\;\\;\\Lambda f(\\tau):\\,= \\int_{0}^{b}K(\\tau-\\tau')f(\\tau')d\\tau', \\end{equation} with norm \\begin{equation} \\label{eq8} \\|\\Lambda\\|_{\\infty}= \\int_{0}^{b/2}E_1(\\tau)d\\tau=1-E_2(b/2). \\end{equation} Here, $E_2(\\tau):=\\int_{0}^{1}\\exp(-\\tau/x)dx$ is the exponential integral function of order 2. Equation (\\ref{eq6}), which can be written in the form \\begin{equation} \\label{eq9} Q=1+a \\Lambda Q, \\end{equation} has therefore a unique solution in $C^0([0,b])$ provided that $a<1$ or $b<+\\infty$.\\\\ This problem is a basic one in stellar atmospheres theory, and more generally in transport theory. It describes in the simplest way the multiple scattering of some type of particles (here photons) on scattering centers distributed uniformly in a slab of finite thickness, a very simple 1D-configuration. Its applications in astrophysics and neutronics - among other fields - are presented in \\cite{rutily2002}. It is important to solve this problem very accurately, thus providing a benchmark to validate the numerical solutions of integral equations of the form (\\ref{eq1}).\\\\ Physicists and astrophysicists have developed many methods for solving integral equations of the form (\\ref{eq1}) with a convolution kernel defined by (\\ref{eq2}) \\cite{rutily-bergeat1994}. The main steps for solving the prototype equation (\\ref{eq6}) are summarized in a recent article \\cite{chevallier-rutily2004}, which contains accurate tables of the function $(1-a)Q$ for different values of the parameters $a$ and $b$. While reading this paper, one is struck by the complexity of the ``classical'' solution to Eq. (\\ref{eq6}). It requires the introduction of many intricate auxiliary functions introduced in the literature over more than thirty years. The reader quickly loses the thread of the solution, which reduces his chances of exploiting it for solving problems of the more general form (\\ref{eq1}).\\\\ The aim of the present article is to get around this difficulty by solving Eq. (\\ref{eq6}) straightforwardly, introducing as few auxiliary functions as possible to express its solution. These functions are briefly studied in Appendix A. The method is based on the finite Laplace transform, which reduces the problem (\\ref{eq6}) to solving a Cauchy integral equation over $[-1, +1]$, which in turn can be transformed into two Fredholm integral equations over [0, 1]. This approach has been developed in transport theory after the publication in 1953 of the first English translation of Muskhelishvili's monograph {\\it Singular integral equations} \\cite{muskhelishvili1992}: see, e.g., \\cite{busbridge1955}, \\cite{mullikin1964} and \\cite{carlstedt-mullikin1966}. It can be considered as an extension of the Wiener-Hopf method \\cite{busbridge1960} for solving integral equations of the form (\\ref{eq1}) with $b<\\infty$. Both methods are characterized by an intensive use of the theory of (sectionally) analytic functions, which allows solution of Eq. (\\ref{eq1}) in a concise manner. This is obvious when comparing the solution we derive here to the classical solution of the particular problem (\\ref{eq6}), which does not use any calculation in the complex plane. The former method clarifies the origin and the role of the auxiliary functions expressing the solution to a problem of the form (\\ref{eq1}). Since these functions are independent of the source term $S_0$, they are ``universal'' for a given scattering kernel.\\\\ The remainder of this article is organized as follows: in Sec. 2, the finite Laplace transform of the $Q$-function is calculated on the basis of some recent developments on Cauchy integral equations \\cite{rutily-bergeat2002}-\\cite{rutily-chevallier-bergeat2004}. Then the Laplace transform is inverted and the solution achieved in \\cite{chevallier-rutily2004} is concisely retrieved with the help of the theorem of residues (Sec. 3). It involves two functions $F_+$ and $F_-$ with remarkable properties, as shown in Appendix B. The difficulties arising from the numerical evaluation of the latter functions are investigated in \\cite{chevallier-rutily2004}. ",
        "conclusions": "Using the finite Laplace transform, we derived in a concise way the expression (\\ref{eq35}) of the solution to Eq. (\\ref{eq6}), which is appropriate for numerical evaluation \\cite{chevallier-rutily2004}. Our main objective was to come directly to this expression, with emphasis on the manner in which the four auxiliary functions of the problem were generated. They are ({\\it{i}}) the function $T=T(a,z)$, which characterizes the solution to Eq. (\\ref{eq6}) in an infinite medium (the range $[0,b]$ is replaced by $\\mathbb{R}$), ({\\it{ii}}) the function $H=H(a,z)$ which expresses its solution in a half-space ($b=\\infty$), and ({\\it{iii}}) the functions $\\zeta_{\\pm}=\\zeta_{\\pm}(a,b,z)$ that complete the previous ones in the finite case ($b<\\infty$). The main properties of these functions are summarized in Appendix A. We note that the $T$-function can be expressed in terms of elementary transcendental functions [Eq. (\\ref{eqA3})], the $H$-function is defined {\\it{explicitly}} by an integral on $[0,1]$ [Eq. (\\ref{eqA10})], and the functions $\\zeta_{\\pm}$ are defined {\\it{implicitly}} as the solutions to Fredholm integral equations of the second kind [Eq. (\\ref{eqA13})]. It seems that the problem (\\ref{eq6}) has no exact closed-form solution for $b<\\infty$.\\\\ The approach of this article is appropriate for solving integral equations of the form (\\ref{eq1}), with a kernel defined by (\\ref{eq2}) and any free term. It also applies to convolution kernels defined by functions more general than (\\ref{eq2}), for example of the form \\begin{equation} \\label{eq37} K(\\tau):=\\int_{I}\\Psi(x)\\exp(-\\vert \\tau \\vert /x)dx, \\end{equation} where $I$ is an interval of $\\mathbb{R}$ and $\\Psi$ a function ensuring the existence of the integral. This class of kernels models scattering processes more complex than the one considered in this article, which corresponds to $I=\\,]0,1]$ and $\\Psi(x)=(1/2)(1/x)$.  \\appendix \\renewcommand{\\theequation}{\\Alph{section}\\arabic{equation}} \\setcounter{equation}{0} \\setcounter{section}{1}"
    },
    "0601/astro-ph0601357_arXiv.txt": {
        "abstract": "The large-scale magnetic field of our Galaxy can be probed in three dimensions using Faraday rotation of pulsar signals. We report on the determination of 223 rotation measures from polarization observations of relatively distant southern pulsars made using the Parkes radio telescope. Combined with previously published observations these data give clear evidence for large-scale counterclockwise fields (viewed from the north Galactic pole) in the spiral arms interior to the Sun and weaker evidence for a counterclockwise field in the Perseus arm. However, in interarm regions, including the Solar neighbourhood, we present evidence that suggests that large-scale fields are clockwise. We propose that the large-scale Galactic magnetic field has a bisymmetric structure with reversals on the boundaries of the spiral arms. Streaming motions associated with spiral density waves can directly generate such a structure from an initial inwardly directed radial field. Large-scale fields increase toward the Galactic Center, with a mean value of about 2~$\\mu$G in the Solar neighbourhood and 4~$\\mu$G at a Galactocentric radius of 3 kpc. ",
        "introduction": "A diffuse magnetic field exists on all scales in our Galaxy. This field can be detected through observations of Zeeman splitting of spectral lines, of polarized thermal emission from dust at mm, sub-mm or infrared wavelengths, of optical starlight polarization due to anisotropic scattering by magnetically-aligned dust grains, of radio synchrotron emission, and of Faraday rotation of polarized radio sources \\citep[see][for a review]{hw02}. The first two approaches have been used to detect respectively the line-of-sight strength and the transverse orientation of magnetic fields in molecular clouds \\citep[e.g.][]{cru99,ncr+03,fram03}. Starlight polarization can be used to delineate the orientation of the transverse magnetic field in the interstellar medium within 2 or 3 kpc of the Sun. Careful analysis of such data show that the local field is mainly parallel to the Galactic plane and follows local spiral arms \\citep[e.g.][]{hei96}. Since we live near the edge of the Galactic disk, we cannot have a face-on view of the global magnetic field structure in our Galaxy through polarized synchrotron emission, as is possible for nearby spiral galaxies \\citep[e.g.][]{bbm+96}. Polarization observations of synchrotron continuum radiation from the Galactic disk \\citep[e.g.][]{rfr+02} give the transverse direction of the field in the emission region and some indication of its strength.  Large-angular-scale features are seen emerging from the Galactic disk, for example, the North Polar Spur \\citep[e.g.][]{jfr87,dhjs97,drrf99,rfr+02}, and the vertical filaments near the Galactic center \\citep{hsg+92,dhr+98}. There are also many small-angular-scale structures resulting from diffuse polarized emission at different distances which are modified by foreground Faraday screens \\citep{gdm+01,ul02,hkd03a}.  Faraday rotation of linearly polarized radiation from pulsars and extragalactic radio sources is a powerful probe of the diffuse magnetic field in the Galaxy \\citep[e.g.,][]{sk80,sf83,ls89,rk89,hq94,hmbb97,id98,fsss01}. Faraday rotation gives a measure of the line-of-sight component of the magnetic field. Extragalactic sources have the advantage of large numbers but pulsars have the advantage of being spread through the Galaxy at approximately known distances, allowing direct three-dimensional mapping of the field.  Pulsars also give a direct estimate of the strength of the field through normalisation by the dispersion measure (DM). The rotation measure (RM) is defined by \\begin{equation} \\phi = {\\rm RM}\\; \\lambda^2 \\end{equation} where $\\phi$ is the position angle in radians of linearly polarised radiation relative to its infinite-frequency ($\\lambda = 0$) value and $\\lambda$ is its wavelength (in m). For a pulsar at distance $D$ (in pc), the RM (in rad~m$^{-2}$) is given by \\begin{equation} {\\rm RM} = 0.810 \\int_{0}^{D} n_e {\\bf B} \\cdot d{\\bf l}, \\end{equation} where $n_e$ is the electron density in cm$^{-3}$, ${\\bf B}$ is the vector magnetic field in $\\mu$G and $d {\\bf l}$ is an elemental vector along the line of sight toward us (positive RMs correspond to fields directed toward us) in pc. With \\begin{equation} {\\rm DM}=\\int_{0}^{D} n_e d l, \\end{equation} we obtain a direct estimate of the field strength weighted by the local free electron density \\begin{equation}\\label{eq_B} \\langle B_{||} \\rangle  = \\frac{\\int_{0}^{D} n_e {\\bf B} \\cdot d{\\bf l} }{\\int_{0}^{D} n_e d l } = 1.232 \\;  \\frac{\\rm RM}{\\rm DM}. \\label{eq-B} \\end{equation} where RM and DM are in their usual units of rad m$^{-2}$ and cm$^{-3}$ pc and $B_{||}$ is in $\\mu$G.  \\citet{man72,man74} was first to systematically measure a number of pulsar RMs and to use them to investigate the large-scale Galactic magnetic field. He concluded that the local uniform field is directed toward Galactic longitude $l\\sim90\\degr$. \\citet{tn80} modeled the magnetic field configuration to fit the 48 pulsar RMs available that time, and they confirmed the local field direction and also found evidence for a field reversal near the Carina--Sagittarius arm. After the large pulsar RM dataset was published by \\citet{hl87}, \\citet{ls89} used 185 pulsar RMs to study the Galactic magnetic field and confirmed the field reversal found by \\citet{tn80}.  \\citet{rl94} observed 27 RMs of distant pulsars in the first Galactic quadrant and provided evidence for a clockwise field (viewed from the north Galactic pole) near the Crux--Scutum arm. These field directions were recently re-examined by \\citet{wck+04} using revised pulsar distances from the NE2001 electron density model \\citep{cl02} together with their 17 new RMs, finding evidence for several field reversals both exterior and interior to the Solar circle. \\citet{hmq99} observed 54 RMs and tentatively identified a counterclockwise field near the Norma arm, which was later confirmed by \\citet{hmlq02} using the RM data discussed in this paper. More recently, \\citet{val05} has reanalysed the available pulsar RM data and interpreted it in terms of an overall clockwise field with a counterclockwise ring of width $\\sim 1$ kpc and radius $\\sim 5$ kpc centered on the Galactic Center.  Pulsar RMs have also been used to study the small-scale random magnetic fields in the Galaxy. Some pairs of pulsars, which are close in sky position and have similar DMs, have very different RMs, indicating an irregular field structure on scales about 100~pc \\citep{ls89}. Some of these irregularities may result from HII regions in the line of sight to the pulsar \\citep{mwkj03}. \\citet{rk89} fitted the single-cell-size model for the residuals of pulsar RMs after the RM contribution of the proposed large-scale ring-field structure was subtracted and obtained a strength for the random field $B_r\\sim5\\mu$G.  \\citet{os93} analyzed the difference of RMs and DMs of pulsar pairs and concluded that $B_r\\sim 4-6\\mu$G independent of cell-size in the range of 10 -- 100~pc. In fact the random fields exist on all scales. \\citet{hfm04} have found the power-law distribution for magnetic field fluctuations as $E_B(k)= C \\ (k / {\\rm kpc^{-1}})^{-0.37\\pm0.10}$ at scales from $1/k=$ 0.5~kpc to 15~kpc,with $C= (6.8\\pm0.8)\\ 10^{-13} {\\rm erg \\ cm^{-3} \\ kpc}$, corresponding to an rms field of $\\sim6\\mu$G in the scale range. The interstellar magnetic field is stronger at smaller scales and may be strongest at the scales of energy injection by supernova explosions and stellar winds (1 -- 10~pc).  The Parkes multibeam survey has discovered a large number of low-latitude and relatively distant pulsars \\citep{mlc+01,mhl+02,kbm+03,hfs+04}, providing a unique opportunity to probe the diffuse magnetic field in a substantial fraction of the Galactic disk with much improved spatial resolution. In addition, improved estimates of pulsar distances are available from the NE2001 electron density model \\citep{cl02}. In this paper, we adopt a distance of the Sun from the Galactic center of $R_{\\odot}=8.5$~kpc.  We have used the Parkes 64-m telescope of the Australia Telescope National Facility to observe the polarization properties of 270 pulsars, most of which were discovered in the Parkes Multibeam Pulsar Survey. After processing, we obtained 223 pulsar RMs which we present in \\S2. All available pulsar RM data have been used to reveal magnetic field directions along the spiral arms and in interarm regions, as presented in \\S3. The field strength and its Galactocentric radial dependence are analysed in \\S4. Our model for the large-scale Galactic field is discussed in \\S5 and concluding remarks are in \\S6. ",
        "conclusions": "Pulsars have unique advantages as probes of the large-scale Galactic magnetic field. Their distribution throughout the Galaxy at approximately known distances allows a true three-dimensional mapping of the large-scale field structure. Furthermore, combined with the measured DMs, pulsar RMs give us a direct measure of the mean line-of-sight field strength along the path, weighted by the local electron density.  We have measured RMs for 223 pulsars, most of which lie in the fourth and first Galactic quadrants and are relatively distant. These new measurements enable us to investigate the structure of the Galactic magnetic field over a much larger region than was previously possible. Clear evidence is found for coherent large-scale fields aligned with the spiral-arm structure of the Galaxy. In all of the inner arms (Norma, Crux-Scutum, Carina-Sagittarius) there is strong evidence that the large-scale fields are counterclockwise when viewed from the north Galactic pole. Weaker evidence also suggests that a counterclockwise field is present in the Perseus arm. On the other hand, at least in the local region and in the inner Galaxy in the fourth quadrant, there is good evidence that the fields in interarm regions are similarly coherent, but clockwise in orientation. Evidence is also presented that large-scale magnetic fields are stronger in the inner part of the Galaxy and that fields in interarm regions are weaker than those in spiral arms.  We therefore propose a bisymmetric model for the large-scale Galactic magnetic field with reversals on arm-interarm boundaries so that all arm fields are counterclockwise and all interarm fields are clockwise. This model for the Galactic magnetic field is appealing in its simplicity. It receives strong support from an objective analysis of mean line-of-sight fields near tangential points of an equiangular spiral.  However, it clearly needs to be backed up by further observational work and by modelling of the effects of streaming motions on various initial field configurations. It is possible that the mode of initial field formation is not critical in determining the dominant present-day structure. Further rotation measure observations, especially for interarm regions and especially in the first Galactic quadrant, would be especially valuable."
    },
    "0601/astro-ph0601161_arXiv.txt": {
        "abstract": "Resonant relaxation (RR) of orbital angular momenta occurs near massive black holes (MBHs) where the potential is spherical and stellar orbits are nearly Keplerian and so do not precess significantly. The resulting \\textit{coherent} torques efficiently change the magnitude of the angular momenta and rotate the orbital inclination in all directions. As a result, many of the tightly bound stars very near the MBH are rapidly destroyed by falling into the MBH on low-angular momentum orbits, while the orbits of the remaining stars are efficiently randomized. We solve numerically the Fokker-Planck equation in energy for the steady state distribution of a single mass population with a RR sink term. We find that the steady state current of stars, which sustains the accelerated drainage close to the MBH, can be $\\lesssim\\!10$ larger than that due to \\textit{non-coherent} 2-body relaxation alone. RR mostly affects tightly bound stars, and so it increases only moderately the total tidal disruption rate, which is dominated by stars originating from less bound orbits farther away.  We show that the event rate of gravitational wave (GW) emission from inspiraling stars, originating much closer to the MBH, is dominated by RR dynamics. The GW event rate depends on the uncertain efficiency of RR. The efficiency indicated by the few available simulations implies rates $\\lesssim\\!10$ times higher than those predicted by 2-body relaxation, which would improve the prospects of detecting such events by future GW detectors, such as \\textit{LISA.} However, a higher, but still plausible RR efficiency can lead to the drainage of all tightly bound stars and strong suppression of GW events from inspiraling stars. We apply our results to the Galactic MBH, and show that the observed dynamical properties of stars there are consistent with RR. ",
        "introduction": "Galactic nuclei with massive black holes (MBHs) are stellar systems with relaxation times that are often shorter than the age of the Universe. In that case the distribution function (DF) of the system may approach a steady state. This steady state is determined by the boundary conditions (an inner sink, such as the tidal radius of the MBH, and an outer source at the interface with the host galaxy) and by the mutual interactions between the stars themselves. The nature of the {}``microscopic'' interactions between the stars determines the rate at which the system relaxes to its steady state.  With the notable exception of $N$-body simulations (e.g. Baumgardt, Makino \\& Ebisuzaki \\cite{Baum04a}, \\cite{Baum04b}; Preto, Merritt \\& Spurzem \\cite{P04}; Merritt \\& Szell \\cite{MS05}), analyses of the evolution of the DF near a MBH have almost exclusively relied on the assumption that the mechanism through which stars exchange angular momentum and energy is dominated by \\textit{uncorrelated two-body interactions} (Chandrasekhar \\cite{Ch43}; see Binney \\& Tremaine \\cite{BT87} for a more recent discussion). This assumption is made in Fokker-Planck models (e.g. Bahcall \\& Wolf \\cite{BW76} {[}hereafter BW76{]}; Bahcall \\& Wolf \\cite{BW77} {[}hereafter BW77{]}; Cohn \\& Kulsrud \\cite{CK78}; Murphy, Cohn \\& Durisen \\cite{MCD91}), where the microscopic interactions are expressed by the diffusion coefficients, and in Monte Carlo simulations (e.g. Shapiro \\& Marchant \\cite{SM79}; Marchant \\& Shapiro \\cite{MS79}, \\cite{MS80}; Freitag \\& Benz \\cite{FB01}, \\cite{FB02}). Stars around MBHs are described as moving in the smooth average potential of the MBH and the stars, and the scattering by the fluctuating part of the potential is modeled as a hyperbolic Keplerian interaction between a passing star and a test star. The scattering effects accumulate non-coherently in a random-walk fashion.  The (non-resonant) relaxation time $T_{\\mathrm{NR}}$ can be defined as the time $T_{\\mE}$ it takes for the negative specific energy $\\mE\\!\\equiv\\!-v^{2}/2-\\varphi$ of a typical star (hereafter {}``energy'') to change by order unity. This is also the time $T_{J}$ it takes for its specific angular momentum $J$ (hereafter {}``angular momentum'') to change by an amount of order $J_{c}(\\mE)$, the maximal (circular orbit) angular momentum for that energy% \\footnote{\\label{ft:TJ}Throughout this paper, {}``angular momentum relaxation time'' means the time it takes until $J$ is changed by order $J_{c}$, rather than by order $J$. The time-scale for changes by order $J<J_{c}$ is shorter than the angular momentum relaxation time by a factor $(J/J_{c})^{2}$.% }. In a spherical potential $J_{c}\\!=\\! G[\\Mbh+N(>\\!\\mE)\\Ms]/\\sqrt{2\\mE}$, where $\\Mbh$ is the MBH mass, $\\Ms$ is the stellar mass and $N(>\\!\\mE)$ is the number of stars with orbital energies above $\\mE$ (more bound than $\\mE$). On Keplerian orbits $J_{c}\\!=\\!\\sqrt{G\\Mbh a}$, where $a$ is the semi-major axis. When relaxation is dominated by uncorrelated two-body interactions, the {}``non-resonant'' relaxation time $T_{\\textrm{NR}}$ of stars of mass $\\Ms$ can be written in the Keplerian regime as  \\begin{equation} T_{\\textrm{NR}}=A_{\\Lambda}\\left({\\frac{\\Mbh}{\\Ms}}\\right)^{2}{\\frac{P(a)}{N(<a)}}\\qquad(\\Mbh\\!\\gg\\!\\Ms)\\,,\\label{e:tr}\\end{equation} where $P\\!=\\!2\\pi\\sqrt{a^{3}/(G\\Mbh)}$ is the orbital period and $A_{\\Lambda}$ is a dimensionless constant which includes the Coulomb logarithm. We assume throughout a single mass stellar population; for numerical estimates we assume $\\Ms\\!=\\!1\\,\\Mo$.  The assumption of uncorrelated two-body interactions is well-justified in many systems, such as globular clusters. However, Rauch \\& Tremaine (\\cite{RT96}, hereafter RT96) showed that this does not hold for motion in potentials with certain symmetries, in particular for Keplerian motion in the potential of a point mass where the orbits do not precess because of the $1\\!:\\!1$ resonance between the radial and azimuthal frequencies. This is to a good approximation the case for stars orbiting close, but not too close to a MBH, since wider orbits precess due to the potential of the enclosed stellar mass, while tighter orbits precess due to General Relativity (GR). As long as the precession timescale $\\tp$ (the time to change the argument of the periapse $\\omega$ by $\\pi$) is much longer than the orbital period, the orbits effectively remain fixed in space over times $P\\!\\ll\\! t\\!\\ll\\!\\tp$ and the interactions are correlated. The orbiting stars can then be represented by one-dimensional elliptical {}``wires'' whose mass density varies along the wire in proportion to the time spent there (RT96). In this picture the interaction between two stars can be described as the torque between two wires. This results in mutual changes in \\emph{both} the direction and the magnitude of the angular momenta.  Following RT96, we denote such RR, which also changes the magnitude of $\\mathbf{J}$\\emph{, scalar RR}. The change grows coherently ($\\propto\\! t/\\tp$) on timescales $t\\!\\ll\\!\\tp$ and non-coherently ($\\propto\\!\\sqrt{t/T_{\\mathrm{RR}}}$) on longer timescales $t\\!\\gg\\!\\tp$, where $T_{\\mathrm{RR}}$ is the {}``resonant relaxation'' (RR) timescale for the angular momentum (Eq. \\ref{e:TRR}). Generally, $T_{\\mathrm{RR}}\\ll\\! T_{\\mathrm{NR}}$ when $\\Ms N(<\\! a)\\!\\ll\\!\\Mbh$ (see Eq. \\ref{e:TRRM}). Since the potential of the wires is stationary on time-scales $\\ll\\!\\tp$, they do not exchange energy, and the energy relaxation timescale remains long, $T_{\\mE}\\!\\sim\\! T_{\\textrm{NR}}$, even in the resonant regime. The mechanism of RR is reminiscent of the Kozai mechanism in triple stars (Kozai \\cite{KO62}).  A more restricted type of RR occurs in any spherical potential, where the vector $\\mathbf{J}$ is conserved, but $\\omega$ precesses. In that case the orbital rosette can be represented by a mass disk extending between the orbital periapse and apo-apse. The mutual torques exerted by such azimuthally symmetric disks change \\emph{only} the direction of $\\mathbf{J}$, but not its magnitude. This type of RR persists even in the presence of GR precession. Following RT96, we denote such RR, which changes only the direction of $\\mathbf{J}$, \\emph{vector RR}. The significance of the distinction between scalar and vector RR is that scalar RR can deflect stars into {}``loss-cone'' orbits ($J\\!\\lesssim\\! J_{c}\\rightarrow J\\!\\sim\\!0$) where they fall into the MBH, either directly ({}``infall'') or gradually ({}``inspiral''), whereas vector RR can only randomize the orbital orientation, but cannot affect the loss rate.  In this paper we explore the consequences of RR on the stellar DF and on the infall and inspiral rates. The mechanism of RR is briefly reviewed in \\S \\ref{s:RR}. The Galactic Center (GC) model is defined and the assumptions and approximations are discussed in \\S \\ref{s:model}. In \\S\\ref{s:dens} We solve the Fokker-Planck equation in energy and analyze the effects of RR on the DF (\\S \\ref{ss:DF}) and on the stellar current (\\S \\ref{ss:flow}). In \\S \\ref{s:implications} we discuss the observational consequences of RR, such as the event rates of tidal disruption (\\S\\ref{sss:tiddisr}) and gravitational wave (GW) emission (\\S\\ref{sss:GW}). We discuss our results in the context of the young, non-relaxed stars in the GC (\\S \\ref{ss:GC}). We discuss and summarize our results in \\S \\ref{s:summary}. ",
        "conclusions": "\\label{s:summary}  RR is a coherent relaxation process that operates when symmetries in the gravitational potential restrict the evolution of the orbits (R96; RI98; Tremaine \\cite{T05}). As a result, the orbits maintain their orientation over many orbital periods, and over that time the stars exert coherent mutual torques. These efficiently change the orbital angular momentum $\\mathbf{J}$ (energy relaxation continues on the slow NR timescale). RR operates only under specific conditions\\@. However, when these are met, it can be orders of magnitude more efficient than NR. Here we consider \\emph{scalar RR}, which affects both the direction and magnitude of $J$, in the regime around a MBH where the orbits are Keplerian (where both mass and GR precession are negligible), and \\emph{vector RR}, which affects the direction but not the magnitude of $J$, in the regime around a MBH where the potential is spherical. Scalar RR can deflect stars into the MBH, whereas vector RR can only change their orbital orientation.  In this paper we studied the effect RR has on the energy DF of stars near a MBH. We explored the consequences for the disruption and capture of stars by the MBH, in particular through inspiral by GW emission, and used our results to interpret the properties of the observed dynamical components in the GC. The complex full 2+1 $(\\mE,J;t)$ problem was reduced to an approximate 1+1 $(\\mE;t)$ Fokker-Planck equation, which we solved numerically. We also carried out Monte Carlo experiments to test some of our assumptions.  Our results are as follows. (1) RR leads to the depletion of the high-energy end of the DF (stars tightly bound to the MBH), accompanied by an enhanced current of stars to high energies. The exact extent of the depletion and the effective high-energy cutoff depend on the poorly determined efficiency of RR. The currently available estimates of the RR efficiency factor indicate that $\\chi\\!\\sim\\!1$. (2) We confirm the result of RT96 that for $\\chi\\!=\\!1$, the direct tidal disruption rate is only modestly enhanced by RR. This is because scalar RR is most efficient close to the MBH where the orbits are Keplerian, whereas most of the tidally disrupted stars originate farther away from the MBH. A higher RR efficiency than assumed here could substantially \\emph{increase} the tidal disruption rate. (3) We show, in contrast, that the GW inspiral rate is dominated by RR dynamics, and is likely increased by almost an order of magnitude relative to the rate predicted assuming NR only. This is because stars undergoing GW inspiral originate very close to the MBH. A higher RR efficiency than assumed here could substantially \\emph{decrease} the GW inspiral rate by completely depleting the tightly bound stars. (4) We apply our results to the GC and show that vector RR can naturally explain the inner cutoff of the ordered star disks at $0.04$ pc and the transition to the randomized inner S-star cluster, as well as the randomized state of the old red giants in the inner $\\sim\\!1$ pc. We also show that scalar RR is consistent with the presence of the disks, as it is too slow to disrupt them. If the S-stars were born in the disk on circular orbits, then scalar RR may explain may explain their present non-zero eccentricities.  We note that there are additional processes that could be affected by RR, which we did not study here. Stars that undergo tidal capture and subsequent tidal heating ({}``squeezars'', Alexander \\& Morris \\cite{AM03}) also originate close to the MBH, where RR dominates the dynamics. Squeezars could be directly observed in the GC if RR substantially enhances their rates. A similar process is the tidal capture of a binary companion by an intermediate mass black hole in a young cluster and the subsequent Roche-lobe feeding that could power an ultra-luminous X-ray source (Hopman, Portegies Zwart \\& Alexander \\cite{HPZA04}; Hopman \\& Portegies Zwart \\cite{HPZ05}; Baumgardt et al. \\cite{Bea05}). The role of RR in tidal capture is still unclear, as there is no RR quenching by GR precession at the tidal capture radius that can stop the star from being rapidly destroyed.  There are several limitations and uncertainties in our results that will have to be addressed by future studies. Our treatment of the problem in $\\mE$-space incorporates the RR sink terms in approximate form only. The efficiency of RR is poorly determined as it has been calibrated based on a restricted set of $N$-body simulations and its dependence on the parameters of the stellar system has not been investigated in full. This uncertainty is significant, as it could potentially reverse the sign of the effect of RR on the GW inspiral rate. While the rate is increased by an order of magnitude if $\\chi\\!=\\!1$, it is completely suppressed if $\\chi\\!\\gtrsim\\!10$ (Fig. \\ref{f:Chi}).  An important omission in our treatment is the assumption of a single mass stellar population. A multi-mass population will induce mass segregation. This will modify the DF directly (BW77), and likely also change the RR efficiency. The stellar DF also depends on processes that destroy one type of star while not affecting others, such as stellar collisions (e.g. Freitag \\& Benz \\cite{FB01}; \\cite{FB05}).  Progress on these issues will likely require large scale numerical simulations. The effects of RR on the DF of stellar systems with MBHs have not yet been studied by $N$-body simulations. The recent fully self-consistent $N$-body simulations of Baumgardt et al. (\\cite{Baum04a}, \\cite{Baum04b}), Preto et al. (\\cite{P04}) and Merritt \\& Szell (\\cite{MS05}), indicate that such studies are almost within reach of current hardware. Such simulations are particularly important for predicting GW inspiral rates, as these depend on a combination of complex mechanisms, including mass-segregation and RR, which are hard to assess with (semi-) analytical methods (e.g. Baumgardt et al. \\cite{Bea05}). Our analysis shows that it may be of considerable importance to use a relativistic potential in such simulations (see e.g. Rauch \\cite{R99}), since otherwise it is unclear whether RR is quenched, and a large number of stars may fall directly into the MBH without emitting an observable GW signal. Yet larger-scale modeling will require abandoning the exact $N$-body approach in favor of approximate methods such as Monte Carlo simulations with RR (Freitag \\& Benz \\cite{FB01}, \\cite{FB02}) or numerical 2+1 Fokker-Planck models.  We note that throughout we assumed a spherically symmetric DF. Deviations from spherical symmetry may affect both RR and the loss-cone structure in phase-space (see e.g. Magorrian \\& Tremaine [\\cite{MT99}], who discuss the possibility of a loss 'wedge' in case of a triaxial DF). Typically this enhances the event rates.  Finally, we note that the observed dynamics of the GC provide a promising empirical basis for calibrating and cross-checking theoretical and numerical studies of RR."
    },
    "0601/astro-ph0601482_arXiv.txt": {
        "abstract": "We present a collation of the available data on the opening angles of jets in X-ray binaries, which in most cases are small ($\\lesssim10$\\degree).  Under the assumption of no confinement, we calculate the Lorentz factors required to produce such small opening angles via the transverse relativistic Doppler effect.  The derived Lorentz factors, which are in most cases lower limits, are found to be large, with a mean $>10$, comparable to those estimated for AGN and much higher than the commonly-assumed values for X-ray binaries of 2--5.  Jet power constraints do not in most cases rule out such high Lorentz factors.  The upper limits on the opening angles show no evidence for smaller Lorentz factors in the steady jets of Cygnus X-1 and GRS\\,1915+105.  In those sources in which deceleration has been observed (notably XTE\\,J1550--564 and Cygnus X-3), some confinement of the jets must be occurring, and we briefly discuss possible confinement mechanisms.  It is however possible that all the jets could be confined, in which case the requirement for high bulk Lorentz factors can be relaxed. ",
        "introduction": "Proper motions of X-ray binary (XRB) jets have often been used to place limits on the jet Lorentz factors.  \\citet{Fen03} recently argued that it was in fact impossible to do more than place a lower limit on the Lorentz factors of the jets from two-sided jet proper motions.  For the persistent, continuous jets observed to exist in the low/hard X-ray states of black hole candidates, Gallo, Fender \\& Pooley (2003) found a universal correlation between the X-ray and radio fluxes of the sources, and used the scatter about this relation to constrain the Lorentz factors of such jets to $\\lesssim 2$. However, \\citet{Hei04} argued that the scatter about such a relation could not be used to constrain the mean Lorentz factor of the jets, but rather only the width of the Lorentz factor distribution.  Other arguments, such as those based on jet power requirements, must be used to determine the absolute values of the jet Lorentz factors.  Bulk jet flow velocities close to $c$, the speed of light, have been inferred in many XRB systems \\citep[e.g.][]{Mir94,Hje95,Fen04}, whereas the transverse expansion speeds have not yet been reliably measured.  To date, there are few reported detections of XRB jets resolved perpendicular to the jet axis.  This places strong upper limits on the opening angles of the jets (often less than a few degrees; see Table \\ref{tab:obs}), and hence on the transverse expansion speeds.  While jets can in principle undergo transverse expansion at a significant fraction of $c$, time dilation effects associated with the bulk motion would reduce the apparent opening angle in the observer's frame.  The magnitude of this effect would be determined by the bulk Lorentz factor of the flow.  This raises the possibility of using the observed opening angle of a freely-expanding jet to constrain its Lorentz factor.  Alternatively, if the Lorentz factors thus derived were incompatible with values deduced from independent methods, a strong argument could be made for jet confinement out to large (parsec) scales in such Galactic sources.  In this paper, we first develop the formalism for deducing the Lorentz factor of a freely-expanding jet given a measurement of the opening angle and the inclination angle of the jet axis to the line of sight. In \\S\\,\\ref{sec:constraints} we compare constraints on the Lorentz factors derived from opening angle considerations with those from other methods.  Transient and steady jets are compared in \\S\\,\\ref{sec:low-hard}, and the derived X-ray binary jet Lorentz factors are compared to those seen in AGN in \\S\\,\\ref{sec:AGN}. We discuss possible mechanisms for jet confinement and a method of using lightcurves to test confinement in \\S\\,\\ref{sec:confinement}.  A summary of the observed properties of the individual sources is given in Appendix \\ref{sec:details}. ",
        "conclusions": "If XRB jets are not confined, but are expanding freely, it is possible to constrain their Lorentz factors from measurements of the jet opening angles.  The small opening angles we observe are in this case a consequence of the transverse Doppler effect slowing the apparent expansion speed in the observer's frame. From the upper limits to the opening angles quoted in the literature, the Lorentz factors thus derived are significantly more relativistic (with a mean Lorentz factor $>10$) than is commonly assumed, and possibly no less so than AGN jets.  However, if the jets we observe do indeed have Lorentz factors in the commonly-assumed range of 2--5, then we can quantify the degree of confinement.  The lateral expansion speed perpendicular to the jet axis must then in some cases be $\\lesssim 0.15c$.  In most cases, we cannot exclude the possibility that the jets are unconfined from jet power constraints, nor from measurements of the proper motions of knots in XRB jets.  From the distances $d<<d_{\\rm max}$ and the observed deceleration in the jets of Cygnus X-3 and XTE\\,J1550--564, we know that the jets in these sources at least must be confined.  The observed opening angle of the jets in SS\\,433 suggests that its jets are also confined, possibly by the cold protons known to exist in the jets.  However, we cannot definitively rule out confinement in any of the other sources considered.  We do not find any evidence for a difference in the velocities of low/hard state jets and transient jets, although the observations only provided lower limits to the Lorentz factors in both cases.  In many cases, and as observed in Circinus X-1, XRB jets could well be significantly more relativistic than is commonly assumed, although this is unlikely in the case of high-mass XRBs.  While we cannot rule out many of the possible confinement mechanisms, in the absence of definitive lower limits to the jet Lorentz factors, the possibility that XRB jets are highly relativistic ($\\Gamma \\sim 10$) should not be ruled out."
    },
    "0601/astro-ph0601211_arXiv.txt": {
        "abstract": "{ The conventional picture of disk accretion onto magnetized neutron stars has been challenged by the spin changes observed in a few X-ray pulsars, and by theoretical results from numerical simulations of disk-magnetized star interactions. These indicate possible accretion during the propeller regime and the spin-down torque increasing with the accretion rate. Here we present a model for the accretion torque exerted by the disk on a magnetized neutron star, assuming accretion continues even for rapid rotators. The accretion torque is shown to have some different characteristics from that in the conventional model, but in accord with observations and numerical calculations of accretion-powered magnetized neutron stars. We  also discuss its possible applications to the spin evolution in X-ray pulsars. ",
        "introduction": "The interaction between a magnetized, rotating star and a surrounding accretion disk takes place in a variety of astrophysical systems, including T Tauri stars, cataclysmic variables, and neutron star X-ray binaries (Frank, King, \\& Raine \\cite{fra02}). Magnetic fields play an important role in transferring angular momentum between the accreting star and the disk. A detailed model for the disk-star interaction was developed by Ghosh \\& Lamb (\\cite{gho79a}, \\cite{gho79b}). In this model the stellar magnetic field lines are assumed to penetrate the accretion disk, and become twisted because of the differential rotation between the star and the disk. The resulting torque can spin up or down the central star, depending on the star's rotation, magnetic field strength, and mass accretion rate. The total torque $N$ exerted on a star with mass $M$ contains two components, \\begin{equation} N=N_0+N_{\\rm mag}. \\end{equation} The torque $N_0$ carried by the material falling onto the star is generally taken to be \\begin{equation} N_0=\\dot{M}R_0^2\\Omega_0 \\end{equation} where $\\Omega_0\\simeq (GM/R_0^3)^{1/2}$ is the (Keplerian) angular velocity at the inner radius $R_0$ of the disk, $G$ the gravitational constant, and $\\dot{M}$ the accretion rate. The magnetic torque $N_{\\rm mag}$ results from the azimuthal field component $B_{\\phi}$ generated by shear motion between the disk and the vertical field component $B_{\\rm z}$\\footnote{Here we have adopted cylindrical coordinate $(R, \\phi, z)$ centered on the star, and the disk is assumed to be located on the $z=0$ plane, perpendicular to the star's spin and magnetic axes.}, \\begin{equation} N_{\\rm mag}=-\\int_{R_0}^{\\infty}R^2B_{\\phi}B_{\\rm z}dR. \\end{equation} The torque $N_{\\rm mag}$ can be positive and negative, depending on the value of the ``fastness parameter\" $\\omega=\\Omega_{\\rm *}/\\Omega_0=(R_0/R_{\\rm c})^{3/2}$, where $\\Omega_{\\rm *}$ is the angular velocity of the star and $R_{\\rm c}\\equiv(GM/\\Omega_{\\rm *}^2)^{1/3}$ the corotation radius. A rapidly rotating star ($\\omega\\lsim 1$) can be spun down with steady accretion. The Ghosh \\& Lamb picture has been modified and extended in several directions by, for example, K$\\ddot{o}$nigl (\\cite{k91}), Campbell (\\cite{cam92}), Lovelace et al. (\\cite{lov95}), Wang (\\cite{wang95}), Yi (\\cite{yi95}), Li \\& Wang (\\cite{lw96}), Li, Wickramasinghe, \\& R$\\ddot{u}$diger (\\cite{lwr96}).  However, the above standard model has been challenged by recent observations on the spin changes in several X-ray pulsars (Bildsten et al. \\cite{bil97}) and accreting millisecond pulsars (Galloway et al. \\cite{gal02}), and by theoretical results from 2- and 3-dimensional numerical simulations of disk accretion to magnetized, rotating stars (Romanova et al. \\cite{rom02}, \\cite{rom03}, \\cite{rom04}). These observational and theoretical facts have raised a number of puzzling issues, including  (1) abrupt spin reversals in X-ray pulsars, (2) spin-down in millisecond X-ray pulsars throughout the outburst, (3) spin-down rates increasing with accretion rate, and (4) accretion in the ``propeller\" regime. These have motivated more recent investigations on the torques exerted by an accretion disk on a magnetized star (Nelson et al. \\cite{nel97}; Yi, Wheeler, \\& Vishniac \\cite{yi97}; Torkelsson \\cite{tor98}; Lovelace, Romanova, \\& Bisnovatyi-Kogan \\cite{lov99}; Locsei \\& Melatos \\cite{loc04}; Rappaport et al. \\cite{rap04}).  In this paper, attempting to solve some of the puzzles related to disk-accreting neutron stars, we have constructed a new model for the accretion torque based on some modifications of the Ghosh \\& Lamb picture. We introduce the model in section 2, and apply it to spin evolution in X-ray pulsars in section 3. In section 4 we discuss the similarities and differences between this work and the competent model proposed by Lovelace et al. (\\cite{lov99}). ",
        "conclusions": "Assuming accretion continues during the (soft) propeller stage, we have constructed a model for the magnetized torque on accreting neutron stars, which changes continuously when $\\omega$ varies from 0 to $\\gsim 1$ (without the singularity problem), and the spin-down torque increases with $\\dot{M}$. Following Yi et al. (\\cite{yi97}), we have applied the model to explain the torque reversal phenomena observed in a few X-ray pulsars.  The same problem has been investigated by Lovelace et al. (\\cite{lov99}). These authors have developed a model for magnetic ``propeller\"-driven outflows for rapidly rotating neutron stars accreting from a disk. They find that the inner radius $R_0$ of the disk depends not only on $\\mu$ and $\\dot{M}$, but also on the star's rotation rate $\\Omega_*$. The most interesting result is that for a given value of $\\Omega_*$, there may exist two solutions of $R_0$, with one $>R_{\\rm c}$ and the other $<R_{\\rm c}$, corresponding to two possible equilibrium configurations. In a transition of $R_0$, the propeller varies between being ``off\" and being ``on\", leading to jumps between spin-up and spin-down. The transitions are assumed to be stochastic or chaotic in nature, and could be triggered by small variations in the accretion flow and magnetic field configuration. The ratio of spin-down to spin-up torque is also shown to be of order unity.  This model shows some similar features as our work. For example, in both models (1) discontinuous transitions of the magnetic torque are invoked to account for the observed abrupt spin changes in X-ray pulsars, (2) the torques are of comparable magnitude, and (3) the spin-down torques increases with $\\dot{M}$. Especially, Eq.~(10) suggests loss of angular momentum at the inner edge of the disk, probably by the magnetic-driven outflows proposed by Lovelace et al. (\\cite{lov99}).  However, there exist significant differences between the two works. First, in the propeller regime, Lovelace et al. (\\cite{lov99}) let most of the accreting material be ejected from the star, while we still allow mass accretion to guarantees no significant variations in the X-ray luminosities during the transition of spin evolution, implying that $\\omega\\gg1$ in Lovelace et al. (\\cite{lov99}), and $\\omega\\gsim 1$ in this work. Second, the torque reversals in our model are caused by the change of the dynamical structure of the disk, as suggested by Yi et al. (1997), at a certain, critical mass accretion rate. In Lovelace et al. (\\cite{lov99}) the transition is stochastic, and could take place at any luminosities. Third, we have adopted the traditional Alf\\'ven radius as the inner radius $R_0$ of the disk in both accretion and propeller phases. In Lovelace et al. (1999) $R_0$ depends on $\\Omega_*$ also, and its value increases with $\\Omega_*$ in the propeller case. Obviously the hypothesis of $\\Omega_*$-depending $R_0$ is more realistic, but how $R_0$ changes with $\\Omega_*$ has not been well understood\\footnote{ Romanova et al. (\\cite{rom04}) indeed observed increasing $R_0$ with $\\Omega_*$ in their calculations. But this change is caused by the fact that the initial dipolar field becomes non-dipolar when the star rotates very rapidly.}. Future observations are expected to provide more stringent constraints on the mechanisms suggested in both models."
    },
    "0601/astro-ph0601027_arXiv.txt": {
        "abstract": "We have observed four red clump stars in the very old and metal-rich open cluster  \\amm \\ to derive its metallicity, using the high resolution spectrograph SARG mounted on the TNG. Using a spectrum synthesis technique we obtain an average value of [Fe/H] = +0.47 ($\\pm 0.04$, r.m.s. = 0.08) dex. Our method was tested on $\\mu$~Leo, a well studied metal-rich field giant. We also derive average oxygen and carbon abundances for \\amm \\ from synthesis of [O {\\sc i}] at 6300~\\AA \\ and C$_2$ at 5086~\\AA, finding [O/Fe] $\\simeq -0.3$ and [C/Fe] $\\simeq -0.2$. ",
        "introduction": "The determination of the abundances in stars of different age and Galactic location is one of the basic tools to interpret the chemical evolution of the Milky Way disk. Galactic open clusters are particularly well suited to this purpose (e.g., \\citealt{friel95}) since they reach ages as old as the disk itself, cover a large range in metallicities and ages, are observed in  regions of the Galactic disk likely characterized by different star formation histories, and their distances and ages can be determined with a precision not reachable for field stars, except for the nearest ones.  We are presently studying in an accurate and homogeneous way a significant sample of open clusters (\\citealt{bt06} and references therein): reliable distances, reddenings and ages are derived from photometry with the synthetic color-magnitude diagrams technique \\citep{tosi91}, and will be combined with metal abundances from high resolution spectroscopy to provide robust constraints on the current and past disk properties.  We have already presented the detailed chemical abundances of a few old open clusters  \\citep{bragaglia01,c04,c05}. To our sample we now add \\amm \\ which, with an age of about 9-10 Gyr,  is one of the oldest open clusters of our Galaxy and has super solar metallicity (e.g., \\citealt{pg98}; \\citealt{chaboyer99}; \\citealt{stetson03}; \\citealt{carney05}; \\citealt{king05}). This cluster, almost as old as the Galactic disk, is of paramount importance to study the time evolution of the disk properties. Apart from its age, NGC 6791 is interesting for his peculiar Horizontal Branch (HB), mostly composed by red stars, but with a (unusual) blue tail \\citep{ku92,liebert94}. Its study could be relevant for a number of issues, like e.g., the UV upturn in elliptical galaxies, the HB morphology and its connection with mass loss in globular clusters.  Despite its peculiarities and its importance for Galactic formation studies, only one detailed work dealing with the chemical pattern in \\amm, based on modern fine abundance analysis and high resolution spectroscopy, has been published so far: \\cite{pg98} analyzed  the coolest and brightest (at $V = 15.0, ~B-V = 0.48$)  blue  HB star. Besides this case, really high S/N,  high resolution spectra can be obtained in acceptable exposure times only for the brightest, hence cooler giants in this cluster. At the very large metallicity of \\amm \\ these pose a great challenge to the observers (see next Section) and great care has to be taken to ensure the reliability of the analysis of their extremely crowded spectra.  This paper is organized as follows: Sect. 2 presents our data; atmospheric parameters and iron abundances from spectrum synthesis are described in Sect. 3; Sect. 4 provides the abundances of carbon and oxygen;  a discussion and a summary are given in Sect. 5. ",
        "conclusions": "The first detailed study of \\amm \\ of which we are aware of is by \\cite{kinman65}, who published a photographic color-magnitude diagram and derived information on membership and intrinsic colors (i.e., reddening) from low resolution spectra.  There is a general agreement that this cluster is at about the same Galactocentric distance as the Sun, is about twice as old, and about twice as metal-rich. Nevertheless, there are significant differences between properties derived by different authors: $(m-M)_0 \\sim 12.6 - 13.6$, $E(B-V) \\sim 0.09 - 0.23$, age $\\sim 7 - 12$ Gyr, [Fe/H] $\\sim 0.1 - 0.4$ (for a recent review see \\citealt{stetson03}). \\amm \\ has also been observed with the ACS on board the Hubble Space Telescope (\\citealt{king05}; \\citealt{bedin05}), reaching almost the hydrogen-burning limit on the main sequence and defining the White Dwarfs cooling sequence.  Spectroscopic studies of \\amm \\ are fewer, but present interesting results. Since the pioneering work by \\cite{st71}  the cluster has been recognized to be more metal-rich than the Sun; they found  [M/H] = +0.75, but with a method that attributes [M/H] $\\simeq$ +0.6 to NGC 188 and M 67, both presently known to have nearly solar metallicity.  More recently, the red HB (that we call red clump) has been studied by low resolution spectroscopy by  \\cite{hufnagel95}; they wanted to investigate possible correlations between CH and CN band strengths and found none, at variance with similar studies on globular clusters.  \\cite{wj03} took low resolution spectra of 23 K giants in \\amm \\ and derived, using Lick/IDS indices, [Fe/H] = $+0.320 \\pm 0.023 \\pm \\sigma_{sys}$, where $\\sigma_{sys}$, the systematic error attached to their metallicity scale, is unknown. We have only one star in common (3019, their R25) for which the quoted abundance is [Fe/H] = +0.341.  All the four stars analyzed here were observed at low resolution by \\cite{friel89}, where we adopted identification and membership status from. The same group \\citep{friel93} later published their metallicities: [Fe/H] values are $+0.32 \\pm 0.28$, $+0.23 \\pm 0.33$, $+0.39 \\pm 0.15$, $+0.41 \\pm 0.21$ for stars 2014, 3009, 3019, and SE49, respectively. These values (that would give an average metallicity of about +0.34) do not differ much from our new ones, considering that the scale used by \\cite{friel93} gives on average metal-poorer values than others. The final result of their analysis depends  on the calibration of indices using standard stars with metallicity derived  by several different sources. Due to a recent revision of their abundance scale \\citep{friel02}, these abundances were lowered to  $+0.12 \\pm 0.13$, $+0.11 \\pm 0.13$, $+0.15 \\pm 0.11$ and $+0.16 \\pm 0.11$ for the same stars;  for their total sample of 39 stars they derived an average metallicity of +0.11,  $\\sigma = 0.10$  dex.  To date, the only published analysis based on high resolution spectroscopy remains that by \\cite{pg98}. They obtained 4 hours of integration at the 4~m Mayall telescope at Kitt Peak on a blue HB star that is a cluster member both by proper motion and RV. The summed spectrum has S/N and resolution lower than ours (S/N $\\sim$ 30 per pixel, R = 20000), but is less affected by blends because of the much higher temperature (about 7300 K). They employed both $EW$s (for Fe and atmospheric parameters determination) and spectrum synthesis (for all the other measured elements). They discussed the use of different scales, e.g. for $\\log gf$'s and solar reference abundances. On the scale they chose to adopt, their star has [Fe/H] = +0.37 ($\\pm 0.10$) dex; since their solar Fe is $\\log n$ = 7.67 and ours is 7.54, their value corresponds to [Fe/H] = +0.50 on our scale.  Taken at face value, the agreement with our results appears excellent, but it doesn't take into account many differences like e.g., $\\log gf$'s or model atmospheres and the like.  For instance, we have 4 Fe lines in common  with \\cite{pg98}, 1 of Fe {\\sc i} and 3 of Fe {\\sc ii} (see Table \\ref{t:abundances} and their table 1), and the average difference in $\\log gf$'s is about 0.29 (our values minus theirs). They also measured the abundances of several elements by means of extensive spectrum synthesis, finding strong overabundances with respect to the  solar ratio for N and Na (+0.5 or +0.4) and slightly lower ones for Mg and Si (+0.2) or Eu (+0.1). All the other elements, including C and O, display solar-scaled abundances. In our analysis we have measured so far only C and O abundances, finding for them sub-solar ratios. A complete comparison between the two analyses is beyond the purpose of the present paper, also because the chosen targets are quite different in properties, but it will be done once we have measured all other elements.  Summarizing, we have derived the metallicity of \\amm \\ analyzing  the high-resolution, high S/N spectra of four red clump stars, finding an average [Fe/H] = $+0.47\\pm 0.04\\pm 0.08$, where the two error values represent the random term and the systematic one, respectively.  This result has been obtained using spectrum synthesis of Fe {\\sc i} and  Fe {\\sc ii} lines, and adopting atmospheric parameters from the photometry. The soundness of our analysis has also been checked relatively to a well studied, metal-rich field giant star, $\\mu$ Leo.  We have measured O and C abundances using synthesis of the [O {\\sc i}] line at 6300.31 \\AA \\ and of the C$_2$ molecular feature at 5086 \\AA, finding [O/Fe] = $-0.32$ and [C/Fe] = $-0.23$ dex, on average, for \\amm.  \\amm \\ represents a challenge to any model for the formation of the Galactic disk, a possible missing link between globular and open clusters, a genuine puzzle that modern abundance analysis based on high resolution spectroscopy will help to unveil."
    },
    "0601/astro-ph0601577_arXiv.txt": {
        "abstract": "We present new relationships between halo masses ($M_h$) and several galaxy properties, including $r^*$-band luminosities ($L_r$), stellar ($M_{\\mathrm{star}}$) and baryonic masses, stellar velocity dispersions ($\\sigma$), and black hole masses ($M_{\\rm BH}$). Approximate analytic expressions are given. In the galaxy halo mass range $3\\times 10^{10}\\, M_{\\odot}\\leq M_h\\leq 3\\times 10^{13}\\, M_{\\odot}$ the $M_h$--$L_r$, $M_{\\mathrm{star}}$--$M_h$, and $M_{\\mathrm{BH}}$--$M_h$ are well represented by a double power law, with a break at $M_{h,{\\mathrm{break}}}\\approx 3 \\times 10^{11}\\, M_{\\odot}$, corresponding to a mass in stars $M_{\\mathrm{star}}\\sim 1.2\\times 10^{10}\\, M_{\\odot}$, to a $r^*$-band luminosity $L_r\\sim 5 \\times 10^9\\, L_{\\odot}$, to a stellar velocity dispersion $\\sigma \\simeq 88$ km s$^{-1}$, and to a black hole mass $M_{BH}\\sim 9\\times 10^{6}\\,M_{\\odot}$. The $\\sigma$--$M_h$ relation can be approximated by a single power law, though a double power law is a better representation. % Although there are significant systematic errors associated to our method, the derived relationships are in good agreement with the available observational data and have comparable uncertainties. We interpret these relations in terms of the effect of feedback from supernovae and from the active nucleus on the interstellar medium. We argue that the break of the power laws occurs at a mass which marks the transition between the dominance of the stellar and the AGN feedback. ",
        "introduction": "There is a well determined set of cosmological parameters $h\\equiv H_0/100 \\, \\mathrm{km\\, s}^{-1} \\, \\mathrm{Mpc}^{-1}=0.70 \\pm 0.04$, $\\Omega_M = 0.27 \\pm 0.04$, $\\Omega_{\\Lambda} = 0.73 \\pm 0.05$, $t_0=13.7 \\pm 0.2$ Gyr, $\\sigma_8=0.84\\pm 0.04$ emerging from a number of observations (the concordance cosmology, see Spergel et al. 2003). Also the cosmic density of baryons $\\Omega_{b} = 0.044 \\pm 0.004$ has been very precisely determined through both the power spectrum of Cosmic Microwave Background anisotropies and measurements of the primordial abundance of light elements (Cyburt et al. 2001; Olive 2002). An important complementary information is that the density of baryons residing in virialized structures and associated to detectable emissions is much smaller than $\\Omega_{b}$. In fact, traced-by-light baryons in stars and in cold gaseous disks of galaxies and in hot gas in clusters amount to a $\\Omega_{b,\\mathrm{lum}} \\approx (3-4) \\times 10^{-3}\\la 0.1 \\Omega_{b}$ (Persic \\& Salucci 1992, Fugukita et al. 1998, Fukugita \\& Peebles 2004). On the other hand, in rich galaxy clusters the ratio between the mass of the dark matter (DM) component and the mass of the baryon component, mainly in the hot intergalactic gas, practically matches the cosmic ratio $\\Omega_M/\\Omega_b$.  The circumstance that $\\Omega_{b}$ is a factor of about 10 larger than $\\Omega_{b,\\mathrm{lum}}$ puts forth both an observational and a theoretical problem. On one side, observations are needed to detect and locate these ``missing'' baryons (see for a review Stocke, Penton \\& Shull 2003). On the theoretical side, galaxy formation models have to cope with the small amount of baryons presently in gas and stars inside galaxies. No doubt that feedback from stars and AGNs played a relevant role in unbinding large amounts of gas and eventually removing them from the host DM halo (see, e.g., Dekel \\& Silk 1986; Silk \\& Rees 1998; Granato et al. 2001, 2004; Hopkins et al. 2005; Lapi, Cavaliere \\& Menci 2005), but we need to get a detailed quantitative understanding of these processes, which are crucial to comprehend galaxy formation. The stellar feedback is expected to depend on the star formation history and the total energy released to the gas ultimately depends on the total mass of formed stars and on the present-day galaxy luminosity. Correspondingly, the total energy injected by the AGN feedback is ultimately related to the final black hole (BH) mass. The fraction of the gas removed by feedbacks is expected to depend also on the binding energy of the gas itself, which is determined by the galaxy \\emph{virial} mass and by its density distribution. Therefore the relationships between the galaxy halo mass and the galaxy luminosity, the stellar and baryonic mass, the mass of the central BH, are expected to give extremely useful information on the process of galaxy formation and evolution. An additional relevant piece of information is the link between the galaxy halo mass and the velocity dispersion of the old stellar component. This paper is devoted to a statistical study of these relations.  The Halo Occupation Distribution (see, e.g., Kauffmann, Nusser \\& Steinmetz 1997; Peacock \\& Smith 2000; Berlind \\& Weinberg 2002; Magliocchetti \\& Porciani 2003), which specifies the probability $P(N|M)$ that a halo of mass $M$ is hosting $N$ galaxies, is a helpful statistics to establish the link between the host DM halo mass and the observed galaxy luminosity. An additional tool to explore this link is the formalism of the Conditional Luminosity Function (Yang, Mo \\& van den Bosch 2003), which describes how many galaxies of given luminosity reside in a halo of given mass. Following this approach, Yang et al. (2003) investigated the relation between halo mass and luminosity. However, particularly for high halo masses, $M_h\\ga 10^{13}\\, M_{\\odot}$, both methods give information on galaxy systems, more than on large galaxies.  Marinoni \\& Hudson (2002) and Vale \\& Ostriker (2004) suggest that a helpful starting point can be the simple hypothesis that there is a one-to-one, monotonic correspondence between halo mass and resident galaxy luminosity. Then, by equating the integral number density of galaxies as a function of their luminosity and stellar mass to the number density of galaxy halos, one gets a statistical estimate of the DM halo mass associated to galaxies of fixed luminosity or fixed baryon/stellar mass. However, a major problem of the method is the estimate of the mass function of halos hosting one single galaxy, the Galaxy Halo Mass Function (GHMF). To solve the problem, in this paper we use an empirical approach, which takes into account the results of numerical simulations (see, e.g., De Lucia et al. 2004 and references therein) on the halo occupation distribution by adding to the Halo Mass Function (HMF) the contribution of subhalos (Vale \\& Ostriker 2004; van den Bosch, Tormen \\& Giocoli 2005). At large masses we subtract from the HMF the mass function of galaxy groups (Girardi \\& Giuricin 2000; Martinez et al. 2002; Hein\\\"am\\\"aki et al. 2003; Pisani, Ramella \\& Geller 2003).  The mass around a galaxy up to a radial distance from its center much larger than the characteristic scale of light distribution can be inferred from detailed X-ray observations (see, e.g., O'Sullivan \\& Ponman 2004). Also  the statistical  measurements of the shear induced by weak gravitational lensing around galaxies (see Bartelmann, King, \\& Schneider 2001) yield important insights on the halo mass of galaxies (McKay et al. 2002, Sheldon et al. 2004). Though these mass estimates  have significant uncertainties and their extrapolations to the virial radii are not immediate, they nonetheless provide useful reference values to which we compare the outcomes of our method.  The role of stellar and AGN feedback has been discussed by several authors (see, e.g., Dekel \\& Silk 1986; Silk \\& Rees 1998). Granato et al. (2001, 2004) have implemented both feedbacks  in their model of joint formation of QSOs and spheroidal galaxies. More recently the feedbacks have also been introduced into hydrodynamical simulations (Springel, Di Matteo \\& Hernquist 2005). One of the purposes of this paper is  to show how the information coming from the relationships of the galaxy halo mass with measurable galaxy properties (such as the stellar, baryonic and central BH masses) can shed light on the role and on relative importance of the two feedbacks.  The plan of this work is the following. In \\S$\\,2$ we compute the galaxy stellar and baryonic mass functions, exploiting the luminosity function and the mass-to-light ratio of the stellar component inferred from kinematical mass modelling of galaxies. Then, in \\S$\\,3$, we derive the mass function of galactic halos, and, exploiting the relevant galaxy statistics (luminosity function, galaxy star/baryon mass function, velocity dispersion function) and the galaxy halo mass function, we investigate the relationships of the corresponding galaxy properties with the halo mass. The relation of the halo mass with the mass of the central BH in galactic spheroidal components is deduced in \\S$\\,4$, by comparing the central black hole mass function to the galaxy halo mass function. In \\S$\\,5$ we discuss the role of the stellar and AGN feedbacks in shaping the relationships between stellar and baryonic mass and halo mass. \\S$\\,6$ is devoted to the conclusions. ",
        "conclusions": "We computed the stellar and baryon mass function in galaxies exploiting $M/L$ ratios for stars and gas derived from galaxy kinematics. The results turn out to be in agreement with previous analyses based on stellar population models. The total baryonic mass density in galactic structures amounts to $\\Omega_b^G\\approx (3.6\\pm 0.7)\\times 10^{-3}$, of which $\\sim 40\\%$ resides in late-type galaxies. This result confirms the well-known conclusion that only a fraction $\\la 10\\%$ of the cosmic baryons are presently in stars and cold gas within galaxies.  The present-day galaxy halo mass function, i.e., the number density of halos of mass $M_h$ containing a single baryonic core, has been estimated by adding the subhalos to the halo mass function and by subtracting the contributions from galaxy groups and clusters. Such subtraction, which  is required to single out galactic halos, has however a minor effect over the mass range of interest here.  Approximated analytic relationships between the halo mass, $M_h$, the mass in stars, $M_{\\mathrm{star}}$,  and the $r^*$-band luminosity, $L_r$, have been obtained from the functional equations $\\mathrm{\\Phi}[>M_h(q)]= \\mathrm{\\Psi}(>q)$, where $\\mathrm{\\Phi}(>M_h)$ is the number density of galactic halos larger that $M_h$ and $\\mathrm{\\Psi}(>q)$ is the number density of galaxies with either stellar mass greater than $M_{\\mathrm{star}}$ or luminosity greater than $L_r$. The results are in good agreement with $M_h/L_r$ ratios inferred through X-ray mapping of the gravitational potential and through gravitational lensing. Both relations exhibit a double-power law shape with a break around $M_{h,{\\mathrm{break}}}\\approx 3\\times 10^{11}\\, M_{\\odot}$, corresponding to $M_{\\mathrm{star},\\mathrm{break}} \\approx 1.2\\times 10^{10}\\, M_{\\odot}$ and to an absolute magnitude $M_{r,{\\mathrm{break}}}\\approx -19.6$. A transition at about the same magnitude in the galaxy properties has been evidenced by Kauffmann et al. (2003) and Baldry et al. (2004).  An additional interesting outcome of our analysis is that the $M_{\\mathrm{star}}$--$M_h$ relation is already established at redshift $z\\approx 1.7$, in line with the theoretical expectation of the anti-hierarchical baryon collapse scenario (Granato et al. 2004).  Applying the same technique to the local velocity dispersion function of galaxies and to the black hole mass function, we have also computed the $\\sigma$--$M_h$ and $M_{\\rm BH}$--$M_h$ relationships. The former is quite close to a single power law $\\sigma \\propto M_h^{1/3}$. The latter is again a double power law breaking approximately at $M_{h,{\\mathrm{break}}}$, corresponding to $M_{\\rm BH}\\sim 9\\times 10^{6}\\, M_{\\odot}$. The associated velocity dispersion, $\\sigma \\simeq 88$ km s$^{-1}$, is very close to the first estimate of the critical velocity dispersion for the gas removal by SN explosions given by Dekel \\& Silk (1986), who found a critical halo velocity $V_{\\mathrm{crit}}\\sim 120$ km s$^{-1}$, corresponding to a critical velocity dispersion $\\sigma_{\\mathrm{crit}}\\sim 80\\,\\hbox{km}\\,\\hbox{s}^{-1}$.  As a test of our results, we combined the $M_{\\rm BH}$--$M_h$ relation [eq.~(\\ref{eq|MbhMh})] with the $\\sigma$--$M_h$ relation [eq.~(\\ref{eq|sigmaMh})]; as shown by the lower panel of Fig.~\\ref{fig|MbhMh}, the resulting $M_{\\rm BH}$--$\\sigma$ relation is consistent with the observational data.  The relationships we have obtained are model-independent and can be interpreted in terms of feedback effects by supernovae and AGNs in galactic structures. We also presented a simple feedback model, which nicely reproduces the approximate proportionalities $M_{\\rm BH}\\propto M_{\\mathrm{star}}\\propto M_h$ observed in the high mass range, and the break of these relationships at $M_{h,{\\mathrm{break}}}\\approx 3\\times 10^{11}\\, M_{\\odot}$.  At low masses, the $M_{\\mathrm{star}}$--$M_h$ relation derived here ($M_{\\mathrm{star}}\\propto M_h^{3.1}$) is steeper than that yielded by the model ($M_{\\mathrm{star}}\\propto M_h^{5/3}$), and would imply that only a tiny fraction of the baryons initially associated with the halo remains within it in the form of stars (and we know that the gas does not add much to the baryon content of galaxies). On the other hand, if the amount of stars formed is so low, for a standard Salpeter IMF the energy injected by SN explosions is insufficient to expel the residual gas if the baryon fraction is close to the cosmic value. Thus, if the slope of the $M_{\\mathrm{star}}$--$M_h$ relation really is as steep as its face value suggests, we must conclude that either the initial baryon fraction in low-mass galaxies was substantially lower than the cosmic value (due, e.g., to a pre-heating of the intergalactic medium hampering the infall of baryons into shallow potential wells) or that the SN feedback in these objects was substantially stronger than in more massive galaxies. As discussed in \\S$\\,$\\ref{sect:ghmf}, however, the uncertainties on the low luminosity portion of the LF are large enough to allow for a flatter slope, closer to the model prediction and almost sufficient to grant the gas removal by SN feedback.  The errors shown in the figures mostly reflect uncertainties in the $M/L$ ratio. We must not forget however, other error sources. For example, Fig.~\\ref{fig|MsLrMh}$b$ shows that different choices for the GHMF yield a systematic difference in the results, that, at the high-mass end, become comparable to the scatter considered in the same figure. Further uncertainties come from estimates of the GSMF; these are illustrated, in Figs.~\\ref{fig|MsLrMh}, ~\\ref{fig|Ratio}, ~\\ref{fig|Efficiency}, and ~\\ref{fig|feedback}, by comparisons with results obtained using the GSMF by Bell et al. (2003). Nevertheless, our approach provides results consistent with observations, and have comparable uncertainties. Moreover, since our approach bypasses any assumption on the DM profile, it could provide a valuable tool to discriminate among the different models of DM mass distribution in galaxies.  Our analysis has shown that the relationships presented above bear the imprint of the processes ruling the galaxy formation and evolution. Models should eventually comply with them."
    },
    "0601/astro-ph0601094_arXiv.txt": {
        "abstract": "In order to examine the ``giant impact hypothesis'' for the Moon formation, we run the first grid-based, high-resolution hydrodynamic simulations for an impact between proto-Earth and a proto-planet. The spatial resolution for the impact-generated disk is greatly improved from previous particle-based simulations. This allows us to explore fine structures of a circumterrestrial debris disk and its long-term evolution. We find that in order to form a debris disk from which a lunar-sized satellite can be accumulated, the impact must result in a disk of mostly liquid or solid debris, where pressure is not effective, well before the accumulation process starts. If the debris is dominated by vapor gas, strong spiral shocks are generated, and therefore the circumterrestrial disk cannot survive more than several days. This suggests that there could be an appropriate mass range for terrestrial planets to harbor a large moon as a result of giant impacts, since vaporization during an impact depends on the impact energy. ",
        "introduction": "In the current standard scenario of planet formation, the final stage of assemblage of terrestrial planets is collisions among proto-planets of about Mars mass \\citep{koku98,kort00}. It is widely accepted that at this final assemblage stage, the Moon is formed from the circumterrestrial debris disk generated by an off-set impact of the proto-Earth with a Mars-sized proto-planet, which is known as the ``Giant Impact (GI) hypothesis'' \\citep{hart75,came76}. The initial phase of the GI scenario, a giant impact and the formation of the circumterrestrial debris disk, has been studied by a series of hydrodynamic simulations \\citep{benz86,came00,canu01,canu04}. N-body simulations of the accumulation of the Moon from the debris disk, which is the last phase of GI scenario, revealed that a single large moon is formed just outside the Roche limit, at a distance of about three to four times the Earth's radius, within several months \\citep{ida97,koku00}. It is  believed that this GI scenario explains a number of mysteries concerning the origin of the Moon \\citep{came00}: Why is the Moon so large compared to satellites of other planets? Why is the Moon deficient in iron and volatiles compared to the Earth? Why does the Earth-Moon system have large angular momentum?  Almost all hydrodynamic simulations of GI in past decades used the Smoothed Particle Hydrodynamic (SPH) method, which is a particle-based Lagrangian scheme \\citep{lucy77,ging77}. In the SPH method, the numerical accuracy, in other words, the resolution, is determined by the number of SPH particles. The latest simulations \\citep{canu01,canu04} used $10^4-10^5$ SPH particles, which is an improvement of 1-2 orders of magnitude compared to simulations in the previous two decades \\citep{benz86,came00}. It is apparent, however, that the SPH method is not the best scheme to simulate an impact between two proto-planets. Firstly, SPH is not suitable to deal with the strong shocks and shear motion produced by the off-set impact. Moreover, since it is basically the Lagrangian scheme, the current SPH does not have resolutions fine enough for diffuse regions. This is critical problem particularly for GI, because the debris disk consists of only a few \\% of the total mass. As a result, even for a simulation with $10^5$ SPH particles, only a few $10^3$ SPH particles are used to represent the debris disk, and thus the fine structure of the disk is not resolved accurately. In the SPH formalism, the spatial resolution is determined by the `kernel size', which is variable with the local density. For a diffuse circumterrestrial disk, this kernel size can be as large as the radius of the disk itself \\citep[see Figure 1 of ][]{canu01}. This is the main reason that evolution of the debris is followed for very short period ($<1$ day) after the impact \\citep{canu01}. Therefore the post-impact evolution of the debris, which is a key for the GI hypothesis, is not well understood.  Alternatively, one can use grid-based Eulerian methods to overcome the above problems. \\citet{melo89} have written the only preliminary study on the GI by the grid-based method to date. However, they neglected self-gravity of the planets and debris, and the evolution is followed only for less than 1 hour around the impact. Therefore, it cannot be compared to the recent SPH simulations.   In this paper, we present the first three-dimensional hydrodynamic simulations of the giant impact followed by formation of a debris disk, taking into account the self-gravity, using a high-accuracy Eulerian-grid scheme. Our aim is to clarify the long-term (more than 100 hours after the impact) evolution of the debris with the highest numerical accuracy used in the previous simulations for GI. We pay attention especially to effects of pressure in the debris on formation of the large moon, rather than exploring the large parameter space for the mass ratio or orbits of the proto-planets. The previous SPH simulations with $10^4-10^5$ particles would be fine for studying the behavior of GI for the initial several hours, and they suggest that the fraction of mass orbiting the proto-Earth just after the impact is mainly determined by the orbital parameters. On the other hand, the long-term behavior of the debris, and therefore the final mass of the moon depends on the evolution of the debris disk. It is expected that the pressure in the low-density debris should affect the hydrodynamic nature of the disk, e.g., by generating shocks. Pressure in the debris is determined by the thermodynamics and the phase-change of the material after the impact, which is not fully understood for GI \\citep{stev87}. For example, if a substantial fraction of the material is in a hot vapor form after the impact, the thermal pressure dominates dynamics of the disk, but if liquid or solid is a major component of the debris, its behavior could be very different. Since there is no practical numerical codes to treat directly two-phase flow and phase-change, we run hydrodynamic simulations assuming two extreme equations of state, which approximately represent vapor or liquid-dominated materials.  The paper is organized as follows. In \\S 2, we briefly describe our numerical method, equations of state, and initial conditions. We show the numerical results in \\S 3. We give conclusions and discuss an implication on a necessary condition of formation of large satellites for the Earth-type planets in \\S 4. ",
        "conclusions": "We run for the first time three-dimensional grid-base hydrodynamic simulations of the giant impact between proto-Earth and a proto-planet taking into account self-gravity. We assume two types of the equation of state that phenomenologically represent hot gaseous material or liquid/solid material. Our numerical experiments suggest that in order to form a lunar-sized satellite from the circumterrestrial debris disk produced by the giant impact, the most fraction of the debris should not be in a pressure-dominated phase (e.g., hot vapor gas). Otherwise the subsequent disk evolution results in forming only a small satellite because of the fast angular momentum transfer associated with spiral shocks prior to the satellite accretion stage. The time-scale of the angular momentum transfer is an order of 10 days, therefore the accumulation to form a large satellite should be much faster than this. Yet our results may not rule out a possibility forming a satellite as large as the Moon from a pressure-dominated disk, if the disk mass is much larger than the current lunar mass. This might be the case for collisions with a massive impactor, but this causes another problem on a fraction of vapor in the debris (see discussion below).  In agreement with the previous SPH simulations \\citep[e.g.,][]{canu01, canu04}, we find that the predicted lunar mass at $t\\simeq 10$ hours after the collision does not strongly depend on choice of EOS and numerical methods for the same orbital parameters. This means that the early phase of GI is dominated by a gravitational process, not by thermodynamical processes. On the other hand, the late phase of GI, i.e., evolution of the debris disk, is sensitive to EOS, especially the pressure in the disk. In order to clarify this difference, the spatial resolutions of the previous SPH simulations were apparently not fine enough.  Our results give an important implication on a necessary condition for formation of large satellites for the Earth-type planets. The present results suggest that one of the important keys for GI scenario is the fraction of the vaporized debris after the giant impact. If the kinetic energy of the collision is much larger than the latent heat of the major component of the proto-planet, most of the proto-planet could be vaporized by the impact. If this is the case, the disk evolution would not lead to formation of a large moon as explained above. On the other hand, if the impact velocity is slower than a critical value, a large fraction of the material can be in a liquid or solid phase, and shocks do not dominate the angular momentum transfer in the debris disk. In this case, the post-impact evolution of the proto-planet could be similar to that in model B in our experiments, and therefore a single large moon could be formed. By assuming the impact is head-on and comparing the latent heat of SiO$_2$ \\citep{stev87} and the kinetic energy of the impact, we can roughly estimate the critical impact velocity as $\\simeq $15 km s$^{-1}$, which is slightly larger than the surface escape velocity of the Earth. In the grazing impact between proto-Earth and a Mars-sized proto-planet, shocks propagating in the planets involve vaporization, and its velocity is probably comparable (or smaller) to this critical value, and therefore it would be natural to postulate that most of the debris mass is in a liquid phase. This argument leads to an interesting suggestion: if the impact velocity is larger than the critical velocity, in other words, the mass of the planet is larger than, say, a few Earth mass, the giant impact never results in forming a large satellite.  Finally, one should recall that a correct phase-change even for the major component, such as SiO$_2$, during GI followed by formation of a debris disk is still unknown \\citep[e.g.][]{stev87}. For example, the post-impact expansion could produce more vapor in the debris material, and this may change the hydrodynamic property of the circumterrestrial disk. The final fate, namely whether the planet has satellites, and how large they are, depends on the nature of the phase-change. The thermo-dynamical process during catastrophic impact between the proto-planets is still too complicated to be explored using current numerical techniques including the SPH and grid-based methods. This ultimately requires a self-consistent numerical scheme to simulate two-phase flow taking into account the realistic thermodynamical processes."
    },
    "0601/astro-ph0601607_arXiv.txt": {
        "abstract": "{Search for specific chemical signatures of intermittent dissipation of turbulence in diffuse molecular clouds. We observed \\HCOp(1-0) lines and the two lowest rotational transitions of \\twCO\\ and \\thCO\\ with an exceptional signal-to-noise ratio in the translucent environment of low mass dense cores, where turbulence dissipation is expected to take place.  Some of the observed positions belong to a new kind of small scale structures identified in CO(1-0) maps of these environments as the locus of non-Gaussian velocity shears in the statistics of their turbulent velocity field \\ie\\ singular regions generated by the intermittent dissipation of turbulence. We report the detection of broad \\HCOp(1-0) lines ($10 {\\rm mK}< \\Tas < 0.5$ K). The interpretation of 10 of the \\HCOp\\ velocity components is conducted in conjunction with that of the associated optically thin \\thCO\\ emission.  The derived \\HCOp\\ column densities span a broad range, $10^{11}< N(\\HCOp)/\\Delta v <4 \\times 10^{12}$ \\cq/\\kms, and the inferred \\HCOp\\ abundances, $2 \\times 10^{-10}<X(\\HCOp) < 10^{-8}$, are more than one order of magnitude above those produced by steady-state chemistry in gas weakly shielded from UV photons, even at large densities.  We compare our results with the predictions of non-equilibrium chemistry, swiftly triggered in bursts of turbulence dissipation and followed by a slow thermal and chemical relaxation phase, assumed isobaric.  The set of values derived from the observations, \\ie\\ large \\HCOp\\ abundances, temperatures in the range of 100--200 K and densities in the range 100--10$^3$ \\cc, unambiguously belongs to the relaxation phase. The kinematic properties of the gas suggest in turn that the observed \\HCOp\\ line emission results from a space-time average in the beam of the whole cycle followed by the gas and that the chemical enrichment is made at the expense of the non-thermal energy.  Last, we show that the \"warm chemistry\" signature (\\ie\\ large abundances of \\HCOp, \\CHp, \\wat\\ and OH) acquired by the gas within a few hundred years, the duration of the impulsive chemical enrichment, is kept over more than thousand years. During the relaxation phase, the \\wat/OH abundance ratio stays close to the value measured in diffuse gas by the \\textit{ SWAS} satellite, while the OH/\\HCOp\\ ratio increases by more than one order of magnitude.}  \\titlerunning{Broad \\HCOp\\Jone\\ emission in dense core environments} \\authorrunning{Falgarone et al.} ",
        "introduction": "Parsec scale maps of molecular clouds with only moderate star formation provide a comprehensive view of the environment of dense cores before their disruption by young star activity. Whether the tracer is extinction, molecular lines or dust thermal emission, large scale maps reveal long and massive filaments of gas and dust. Dense cores are most often nested in these filaments, which are likely their site of formation.  The open questions raised by star formation thus include those on dense core and filament formation. Filaments being by essence structures at least one order of magnitude longer than thick, their observational study requires maps with large dynamic range.  These observations, now available, open new perspectives: filaments are ubiquitous and span a broad range of mass per unit length (e.g. \\cite{aa94}, \\cite{mizu95} and \\cite{haik04} for massive filaments and \\cite{phb04}, \\cite{fpw91}, \\cite{fpp01} for most tenuous ones).  Recent observations of the environment of low mass dense cores in \\twCO\\ and \\thCO\\ rotational transitions, have also brought to light previously hidden velocity structures.  The turbulent velocity field of core environments, alike turbulent gas, exhibits more large velocity shears at small scale than anticipated from a Gaussian distribution and the spatial distribution of the non-Gaussian occurrences of velocity shears delineates a new kind of small scale structure in these core environments: their locus forms a network of narrow filaments, not coinciding with those of dense gas (\\cite{pf03}, \\cite{phb04}, \\cite{hfp06}, hereafter Paper IV). It has been shown by \\cite{lis96} that they trace the extrema of the vorticity projected in the plane of the sky.  As discussed in \\cite{pf03}, these velocity structures of large vorticity likely trace either fossil shocks which have generated long-lived vorticity, or genuine coherent\\footnote{so-called because their lifetime is significantly longer than their period} vortices, both manifestations of the intermittency of turbulence dissipation, (\\cite{pf00}).  Interestingly, these structures might be the long-searched sites of turbulence dissipation in molecular clouds, as proposed by \\cite{J98} (hereafter JFPF), and thus the sites of the elusive warm chemistry put forward to explain the large abundances of \\CHp\\ (e.g. \\cite{crane95,gredel02,gredel97}), \\HCOp\\ (\\cite{ll96,luli00}) and now \\wat\\ (\\cite{neufeld02,plume04}) observed in diffuse molecular clouds.  Since all the above observations are absorption measurements, in the submillimeter, radio and visible domains, they suffer from line of sight averaging restricted to specific directions and from a lack of spatial information regarding the actual drivers of this \"warm chemistry\".  The present work is the first attempt at detecting the chemical signatures of locations identified in nearby clouds as the possible sites of the \"warm chemistry\" in diffuse molecular clouds, \\ie\\ the locus of enhanced dissipation of supersonic turbulence. It is part of a series of papers dedicated to the properties of this new kind of filaments in dense core environments, in the broad perspective of understanding how non-thermal turbulent energy is eventually converted and lost as dense filaments and cores form, and the timescales and spatial scales involved in the non-linear thermal and chemical evolution of the dissipative structures of turbulence.  Unfortunately, dissipative structures amount to very little mass \\ie\\ the impulsive releases of energy affect only a small fraction of the gas at any time and they are difficult to characterize because their emission is weak in most tracers.  A multi-transition analysis of CO observations, up to the $J=4-3$ transition, provides for this gas a full range of temperatures from 20K up to 200K and densities ranging from 10$^4$ \\cc\\ down to 100 \\cc, (\\cite{fp96}). Recent observations in the two lowest rotational lines of \\twCO\\ and \\thCO\\ show that the gas there is associated spatially, although not coincident, with gas optically thin in the \\twCO(1-0) line (\\cite{phb04}, \\cite{hf06}, hereafter Paper II).  The present  paper reports observations and  analysis of the \\HCOp(1-0)  emission  of a  few  such structures  previously identified  in  nearby molecular  clouds.  Each  of them  is characterized  by a  different value  of the  local velocity shear  measured  in the  \\twCO\\  lines  in  order to  sample different strengths (or  stages) of the dissipation process. In a companion paper, we  present the substructure in one of these filaments as detected in  a mosaic of 13 fields mapped in   the  \\twCO(1-0)   line   with  the   Plateau  de   Bure Interferometer (PdBI) (\\cite{fph}, herafter Paper III).  The target fields and the  observations are described in Section 2.  We  detail the  line analysis  done to  derive molecular abundances in  Section 3.  In Section 4,  we show  that they cannot  be reconciled  with current  models  of steady-state chemistry   and   that   an  alternative   model   involving non-equilibrium heating  and chemistry reasonably reproduces the main features of our  data, although this model is by no means final.  The discussion in  Section 5 is devoted to the limitations  of  our  study  and  the  implications  of  the interpretation we propose. ",
        "conclusions": "\\subsection{The complex space-time average hidden in the observed \\HCOp\\ emission lines.}  The results of the last section seem to be in contradiction with the relaxation scenario described above, in which the \\HCOp\\ abundance decreases with time. Another, probably related question, is: Why, in regard of the six orders of magnitude spanned by the \\HCOp\\ abundance in the proposed non-equilibrium evolution, the observations in emission point to a much narrower range of abundances, $2 \\,10^{-10} <X(\\HCOp)<10^{-8}$?  This may be understood in a framework in which, several phases of the evolution are present simultaneously in the beam, for individual vortices much thinner than the beamsize, here 0.015 pc (or $\\approx 3000$~AU) while vortices, in the JFPF model, have a radius $\\approx 15$AU. Observations sample a mixture of gas: the most diluted and \\HCOp-rich, still in the warm active layers of the vortices, as long as the vortex is alive, gas already out of the warm active layers and relaxing, up to the densest, \\HCOp-poor gas.  The older the burst, the larger the amount of relaxed gas.  The signal detected in emission is sensitive to a combination of the \\HCOp\\ column density, or simply the abundance $X(t)$ assuming the total gas column density being processed is constant with time, the density $n_{\\rm H}(t)$ that enters the line excitation and a term involving the duration of the phase, $\\approx t$, given the broad range of timescales involved.  We have computed the time integral of $X(t)\\,n_{\\rm H}(t)$, taken as a template for the signal in emission, assuming that we observe all the phases in the gas accumulated during at least $\\approx 10^4$~yr, the lifetime of the vortex. For all the cases, this quantity increases linearly with time up to a maximum at about $10^3$~yr and does not increase further.  The observations in emission are thus likely to be most sensitive to the the relaxation phase because it lasts longer and involves denser gas than the brief phase of chemical enrichment in gas too diluted to significantly excite the \\HCOp(1-0) line.  At the opposite, density does not weight the observations in absorption. The warmest and richest phase of the evolution, may have been detected by \\cite{lilu00} as the broad and weak absorption seen in the direction of B0355+508 over the whole velocity extent of the H~I absorption line.   More detailed comparison between abundances inferred from observations in absorption and in emission involves time and space averages which cannot be achieved at this stage of the study. The present observations have been obtained on one single ensemble of such dissipative structures, possibly generated simultaneously. Thus, space/time average is likely more simple than that required to interpret absorption measurements that sample depths of gas at parsec scales.   \\subsection{Further support for non-equilibrium chemistry: \"warm\" molecules abundance ratios} The present work stresses the need for a non-equilibrium mechanism to explain the large abundances of \\HCOp\\ detected in gas of moderate density and low shielding from the ambient ISRF. In this scenario, the gas is driven locally past a high temperature threshold by an impulsive release of turbulent energy. Once the gas has escaped the region of large dissipation, it cools down and condenses and the chemistry evolves accordingly.  \\begin{figure} \\includegraphics[height=\\hsize,angle=-90]{hotmol-nt.eps} \\caption{Time-dependent evolution of {\\it (top)} the OH (thin solid), \\wat\\ (thick solid), \\HCOp\\ (dashed), and \\CHp\\ (thin dashed) abundances, {\\it (middle)} the kinetic temperature and {\\it (bottom)} the total hydrogen density, along the isobaric cooling of a parcel of gas formerly enriched chemically in a burst of dissipation of turbulent energy. The initial conditions here are $n_{{\\rm H},i}$ =100~\\cc, $\\Gamma_d=10^{-21}$~\\eccs\\ and the shielding from ISRF $A_V$=1.} \\label{fig:hotmol} \\end{figure}  \\begin{figure} \\includegraphics[height=\\hsize,angle=-90]{hotmol-ratio.eps} \\caption{Time-dependent evolution of abundance ratios, OH/\\HCOp\\ (thick), \\wat/OH (thin) and \\CHp/\\HCOp\\ (thick dashed), along the same isobaric cooling as in Fig.~\\ref{fig:hotmol}.} \\label{fig:ratio} \\end{figure}  Our results are independent of the details of the processes at the origin of the warm chemistry. \\cite{fp98} and \\cite{gredel02} have shown that slow MHD shocks of velocities $V_s \\sim 10$ \\kms\\ reproduce the OH/\\HCOp\\ abundance ratios observed in the diffuse medium (\\cite{ll96}).  In MHD shocks, as in magnetized coherent vortices, gas is temporarily heated to temperatures as large as $10^3$ K. The triggered warm chemistry is characterized by coexisting large abundances of OH, H$_2$O, \\HCOp\\ and \\CHp, water and \\HCOp\\ being daughters of OH and \\CHp, respectively (JFPF).  Their time-dependent evolution in the relaxation phase, is displayed in Fig.~\\ref{fig:hotmol} for one of the cases discussed in Section 4 (initial density $n_{{\\rm H},i}=100$~\\cc, $\\Gamma_d=10^{-21}$~\\eccs\\ and $A_V=1$).  The abundance ratios along this evolution are interesting in their marked differences (Fig.~\\ref{fig:ratio}).  The \\wat/OH ratio varies by less than a factor 3 along the time-dependent evolution and up to $10^3$~yr after the beginning of the relaxation, stays close to the values found in diffuse molecular clouds by the {\\it SWAS} satellite (\\cite{neufeld02}, \\cite{plume04}).  The OH/\\HCOp\\ ratio varies by more than a factor 10 along the same evolution. We note however that its variations bracket the average value OH/\\HCOp$\\sim$ 30 found by \\cite{ll96} in their random sampling of the diffuse ISM provided by their absorption lines survey towards extragalactic sources.  In contrast, the \\CHp/\\HCOp\\ ratio varies very little, before $10^3$ yr only.  The value of the ratio, though, critically depends on the density (JFPF).  This follows from the fact that in our model of warm chemistry, \\HCOp\\ is not only a daughter molecule of \\CHp. It also forms via other routes such as CO$^+$ + \\HH\\ $\\rightarrow$ \\HCOp\\ + H, triggered by the endothermic charge exchange between \\Hp\\ and O ($\\Delta E/k=227$ K), and \\Hthp\\ + CO $\\rightarrow$ \\HCOp\\ + \\HH, also slightly enhanced in warm chemistry because both CO and \\Hthp\\ have increased abundances in the warm layer of the vortex.  This route may even be a dominant formation scheme for \\HCOp, if the large abundance of \\Hthp\\ detected on the line of sight towards $\\zeta$ Per by \\cite{mccall} happens to be the rule in the diffuse ISM.  There is also an efficient channel to form \\CtH\\ in the warm chemistry. It proceeds via the endothermic reaction ($\\Delta E/k = 1760$ K) CH + \\HH $\\rightarrow$ \\CHt\\ + {\\rm H}, followed by a series of fast reactions: \\Cp\\ + \\CHt\\ $\\rightarrow$ \\CtHp\\ + {\\rm H}, the hydrogenation of \\CtHp\\ into \\CtHtp\\, the recombination of \\CtHtp\\ and terminating with photodissociation of \\CtHt\\ which provides \\CtH.  These results are encouraging and suggest that impulsive chemical enrichment of the CNM by turbulence dissipation bears observable signatures up to several 10$^3$~yr after the end of the enrichment.  \\subsection{Large scale signatures of small scale activity. }  Other signatures, independent of chemistry, have been found for the existence in the CNM of gas at a few hundred Kelvin. They comprise collisional excitation of \\HH\\ far from UV sources susceptible to generate fluorescence (\\cite{gry02,lacour05,f05}). In the former cases, \\HH\\ is detected in UV absorption against nearby late B stars, in the third case, the pure rotational lines of \\HH\\ have been observed along a long line of sight across the Galaxy avoiding star forming regions. The rotational temperature of the warm \\HH\\ in these data is 276 K.  In all cases the fraction of warm \\HH\\ in the total column of gas sampled is the same, of the order of a few percent.  It is the same fraction as that required in the Solar Neighbourhood to reproduce the large observed abundances of \\CHp\\ within the CNM (JFPF). A few percent of warm gas, for which UV photons cannot be the sole heating source, is ubiquitous in the cold diffuse medium.  Last, we would like to stress the spatial scales presumably involved in the framework of the existing model, where individual vortices have transverse hydrogen column densities as small as $N_{\\rm H} \\approx 2\\times 10^{16}$\\cq\\ corresponding to initial values $n_{\\rm H,i}=100$\\cc\\ and $l_i=15$AU (JFPF).  We illustrate the point with the solutions of case C: $N(\\HCOp) \\approx 8\\times 10^{11}$\\cq, $X(\\HCOp)\\approx 2\\times 10^{-9}$, $n_H \\approx 500$\\cc. The total hydrogen column density of observed material is thus $N_{\\rm H}=N(\\HCOp)/X(\\HCOp)=4\\times 10^{20}$\\cq, corresponding to 2$\\times 10^4$ vortices in the IRAM beam (here we assume that the isobaric evolution occurs at constant mass, so that the transverse column density of vortices remains the same). For the gas at 500\\cc, the transverse size has a thickness of about 3 AU. The ensemble of the 2$\\times 10^4$ vortices, probably braided together, as vortex tubes do e.g \\cite{ppw94}, would form an elongated structure about 450 AU thick, if the vortices are in contact with one another.  The IRAM beam at the distance of the sources is a few times larger and the size of the structure seen with the PdBI is of this order of magnitude (Paper III).  This rough estimate just shows that the existence of chemical inhomogeneities at AU scales is not ruled out in the diffuse ISM.  It would certainly help understanding the remarkable similarity of the OH and \\HCOp\\ absorption line profiles reported by \\cite{lilu00}.   Chemical and thermal relaxation timescales are so short compared to the dynamical lifetime of molecular clouds (of at least a million year) that all the stages of the above evolution should co-exist along a random line of sight.  At present, the receiver sensitivity and the telescope beam sizes provide, when observed in emission, a signal which is the space and time average of myriads of small unresolved structures, likely caught at different epoch in their evolution. Molecular line observations in absorption in the UV, visible and now submm domains are much more sensitive and have presumably already started to sample this apparently untractable complexity."
    },
    "0601/astro-ph0601431_arXiv.txt": {
        "abstract": "On an empirical level, the most successful alternative to dark matter in bound gravitational systems is the modified Newtonian dynamics, or MOND, proposed by Milgrom. Here I discuss the attempts to formulate MOND as a modification of General Relativity. I begin with a summary of the phenomenological successes of MOND and then discuss the various covariant theories that have been proposed as a basis for the idea.  I show why these proposals have led inevitably to a multi-field theory.  I describe in some detail TeVeS, the tensor-vector-scalar theory proposed by Bekenstein, and discuss its successes and shortcomings.  This lecture is primarily pedagogical and directed to those with some, but not a deep, background in General Relativity ",
        "introduction": "There is now compelling observational support for a standard cosmological model.  It is most impressive that this evidence is derived from very different observational techniques applied to very different phenomena: from precise measurements of anisotropies in the Cosmic Microwave Background (CMB) \\cite{speal03}; from systematic photometric observations of the light curves of distant supernovae \\cite{pereal99,garn,tonry}; from redshift surveys mapping the distribution of observable matter on large scale and interpreting that distribution in the context of structure formation by gravitational collapse \\cite{sanch,eisen}.  Using the standard parameterised Friedmann-Robertson-Walker models (FRW), all of these observations imply a convergence to a narrow range of parameters that characterise the Universe; this convergence is rightly heralded as a remarkable achievement of the past decade.  However, the Universe that we are presented with is strange in its composition:  only five percent is the ordinary baryonic matter that we are familiar with; twenty-five percent consists of pressureless dark matter presumed to be fundamental particles that are as yet undetected by other means; and about seventy percent is the even stranger negative pressure dark energy, possibly identified with a cosmological term in Einstein's field equation, and emerging relatively recently in cosmic history as the dominate contributer to the energy density budget of the Universe.  A general sense of unease, primarily with this dark energy, has led a number of people to consider the possibility that gravity may not be described by standard four-dimensional General Relativity (GR) on large scale (see Sami, this volume)-- that is to say, perhaps the left-hand-side rather than the right-hand-side of the Einstein equation should be reconsidered.  Various possibilities have been proposed-- possibilities ranging from the addition of a scalar field with a non-standard kinetic term, K-essence \\cite{apms}; to gravitational actions consisting of general functions of the usual gravitational invariant, $F(R)$ theories (Sotiriou, this volume and \\cite{cct03,ceal}); to braneworld scenarios with leakage of gravitons into a higher dimensional bulk (\\cite{dvali} and Maartens, this volume).  But, in fact, there is a longer history of modifying gravity in connection with the dark matter problem-- primarily that aspect of the problem broadly described as ``missing mass'' in bound gravitational systems such as galaxies or clusters of galaxies.  The observations of this phenomenology have an even longer history, going back to the discovery of a substantial discrepancy between the dynamical mass and the luminous mass in clusters of galaxies \\cite{zwk33}. The precise measurement of rotation curves of spiral galaxies in the 1970's and 1980's, primarily by 21 cm line observations which extend well beyond the visible disk of the galaxy \\cite{Bos78,bg89}, demonstrated dramatically that this discrepancy is also present in galaxy systems.  A fundamental, often implicit, aspect of the cosmological paradigm is that this observed discrepancy in bound systems is due to the cosmological dark matter-- that the cosmological dark matter clusters on small scale and promotes the formation of virialized systems via gravitational collapse in the expanding Universe.  The necessity of clustering on the scale of small galaxies implies that there are no phase space constraints on the density of the dark matter and, hence, that it is cold, or non-relativistic at the epoch of matter-radiation equality \\cite{bdsl83}. The exact nature of the hypothetical cold dark matter (CDM) is unknown but particle physics theory beyond the standard model provides a number of candidates.  There are observational problems connected with the absence of phase space constraints in this dark matter fluid, problems such as the formation of numerous but unseen satellites of larger galaxies \\cite{deluc} and the prediction of cusps in the central density distributions of galaxies-- cusps which are not evident in the rotation curves \\cite{debeal}.  But it is usually taken as a article of faith that ``complicated astrophysical processes'' such as star formation and resulting feed-back will solve these problems.  The motivation behind considering modifications of gravity as an alternative to CDM is basically the same as that underlying modified gravity as an alternative to dark energy:  when a theory, in this case GR, requires the existence of a medium which has not been, or cannot be, detected by means other than its global gravitational influence, i.e., an ether, then it is not unreasonable to question that theory. The primary driver for such proposals has been the direct observation of discrepancies in bound systems-- galaxies and clusters of galaxies-- rather than cosmological considerations, such as that of structure formation in an expanding Universe.  The most successful of the several suggestions, modified Newtonian dynamics or MOND, has an entirely phenomenological rather than theoretical basis \\cite{m83a, m83b,m83c}. In accounting for the detailed kinematics of galaxies and galaxy groups, while encompassing global scaling relations and empirical photometric rules, MOND has, with one simple formula and one new fixed parameter, subsumed a wide range of apparently disconnected phenomena.  In this respect it is similar to the early proposal of continental drift by Alfred Wegener in 1912.  This suggestion explained a number of apparently disconnected geological and palaeontological facts but had no basis in deeper theory; no one, including Wegener, could conceive of a mechanism by which giant land masses could drift through the oceans of the earth.  Hence the idea was met with considerable ridicule by the then contemporary community of geologists and relegated to derisive asides in introductory textbooks. It was decades later, after the development of the modern theory of plate tectonics and direct experimental support provided by the frozen-in magnetic field reversals near mid-oceanic rifts, that the theory underlying continental drift became the central paradigm of geology and recognised as the principal process that structures the surface of the earth \\cite{wil}. I do not wish to draw a close analogy between MOND and the historical theory of continental drift, but only to emphasise the precedent:  an idea can be basically correct but not generally accepted until there is an understandable underlying physical mechanism-- until the idea makes contact with more familiar physical concepts.  The search for a physical mechanism underlying modified Newtonian dynamics is the subject here.  I begin with a summary of the phenomenological successes of the idea, but, because this has been reviewed extensively before \\cite{sm02}, I will be brief.  I consider the proposals that have been made for modifications of GR as a basis of MOND.  These proposals have led to the current best candidate-- the tensor-vector-scalar (TeVeS) theory of Bekenstein \\cite{jdb04}, a theory that is complicated but free of obvious pathologies. I summarise the successes and shortcomings of the theory, and I present an alternative form of TeVeS which may provide a more natural basis to the theory.  I end by a discussion of more speculative possibilities. ",
        "conclusions": "Here I have outlined the attempts that have been made to define modifications of gravity that may underly the highly successful empirically-based MOND, proposed by Milgrom as an alternative to dark matter in bound self-gravitating systems. These attempts lead inevitably to a multi-field theory of gravity-- the Einstein metric to provide the phenomenology of GR in the strong field limit, the scalar field to provide the MOND phenomenology most apparent in the outskirts of galaxies and in low surface brightness systems, and the vector field to provide a disformal relation between the Einstein and physical metrics-- necessary for the observed degree of gravitational lensing.  I re-emphasise that this process has been entirely driven by phenomenology and the need to cure perceived pathologies; there remains no connection to more a priori theoretical considerations or grand unifying principles such as General Covariance or Gauge Invariance.  It would, of course, be a dramatic development if something like MOND were to emerge as a incidental consequence of string theory or a higher dimensional description of the Universe, but, in my opinion, this is unlikely.  It is more probable that an empirically based prescription, such as MOND, will point the way to the correct theory.  The coincidence between the critical acceleration and $cH_0$ (or possibly the cosmological constant) must be an essential clue.  MOND must be described by an effective theory; that is, the theory predicts this phenomenology only in a cosmological context.  The aspect, and apparent necessity, of a preferred frame invites further speculation:  Perhaps cosmology is described by a preferred frame theory (there certainly is an observed preferred frame) with a long range force mediated by a scalar field coupled to a dynamical vector field as well as the gravitational metric.  With the sort of bi-scalar Lagrangian implied by TeVeS, the scalar coupling to matter becomes very weak in regions of high field gradients (near mass concentrations).  This protects the Solar System from detectable preferred frame effects where the theory essentially reduces to General Relativity. Because we live a region of high field gradients, we are fooled into thinking that General Relativity is all there is.  Only the relatively recent observations of the outskirts of galaxies or objects of low surface brightness (or perhaps the Pioneer anomaly) reveal that there may be something more to gravity.  On the other hand, it may well be that we have been pursuing a mirage with tensor-vector-scalar theories.  Perhaps the basis of MOND lies, as Milgrom has argued, with modified particle action-- modified inertia-- rather than modified gravity \\cite{m94a,m99}.  For a classical relativist this distinction between modified gravity and modified inertia is meaningless-- in relativity, inertia and gravity are two sides of the same coin; one may be transformed into another by a change of frame. But perhaps in the limit of low accelerations, lower than the fundamental cosmological acceleration $cH_0$, that distinction is restored \\cite{m05}.  It is provocative that the Unruh radiation experienced by a uniformly accelerating observer, changes its character at accelerations below $c\\sqrt{\\Lambda}$ in a de Sitter universe \\cite{m99}. If the temperature difference between the accelerating observer and the static observer in the de Sitter Universe is proportional to inertia, then we derive an inertia-acceleration relation very similar to that required by MOND \\cite{m99}. At present this is all very speculative, but it presents the possibility that we may be going down a false path with attempted modifications of GR through the addition of extra fields.  In any case, the essential significance of TeVeS is not that it, at present, constitutes the final theory of MOND. Rather, the theory provides a counter-example to the often heard claim that MOND is not viable because it has no covariant basis.  It is a pleasure to thank Jacob Bekenstein and Moti Milgrom for helpful comments on this manuscript and for many enlightening conversations over the years.  I also thank Renzo Sancisi for helpful discussions on the ``dark matter-visible matter coupling'' in galaxies and Martin Zwaan for sending the data which allowed me to produce Fig.\\ 4. I am very grateful to the organisers of the Third Aegean Summer School on the Invisible Universe, and especially, Lefteris Papantonopoulos, for all their efforts in making this school a most enjoyable and stimulating event."
    },
    "0601/hep-th0601189_arXiv.txt": {
        "abstract": "Recent observations confirm that our universe  is flat and consists of a dark energy component with negative pressure. This dark energy is responsible for the recent cosmic acceleration as well as determines the feature of future evolution of the universe. In this paper, we discuss the dark energy of the universe in the framework of scalar-tensor cosmology. In the very early universe, the gravitational scalar field $\\phi$ plays the roll of the inflaton field and drives the universe to expand exponentially. In this period the field $\\phi$ acts as a cosmological constant and dominates the energy budget, the equation of state ( EoS ) is $w=-1$. The universe exits from inflation gracefully and with no reheating. Afterwards, the  field $\\phi$ appears as a cold dark matter and continues to dominate the energy budget, the universe expands according to $\\frac{2}{3}$ power law, the EoS is $w=0$. Eventually, by the epoch of $z\\sim O(1)$, the field $\\phi$ contributes a significant component of dark energy with negative pressure and accellerates the late universe. In the future the universe will expand acceleratedly according to  $a(t)\\sim t^{1.31}$.\\\\ \\\\ PACS number(s): 98.80.Cq, 04.50.+h ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/hep-th0601230_arXiv.txt": {
        "abstract": "\\noindent The evolution of multiple scalar fields in cosmology has been much studied, particularly when the potential is formed from a series of exponentials. For a certain subclass of such systems it is possible to get ``assisted`` behaviour, where the presence of multiple terms in the potential effectively makes it shallower than the individual terms indicate. It is also known that when compactifying on coset spaces one can achieve a consistent truncation to an effective theory which contains many exponential terms, however, if there are too many exponentials then exact scaling solutions do not exist. In this paper we study the potentials arising from such compactifications of eleven dimensional supergravity and analyse the regions of parameter space which could lead to scaling behaviour. ",
        "introduction": "\\label{sec:introduction}  Extra dimensions are ubiquitous in unified models of gravity and particles, dating from the early work of Kaluza and of Klein to the modern ideas of string and M-theory. This raises the question of what to do with the extra dimensions? In a fully dynamical model we expect that at some time these extra dimensions should evolve, possibly having some impact on the Universe we observe today. In the context of dimensional reduction, where the extra dimensions take the form of a small, compact manifold, the dynamics of the internal space typically manifests itself through the dynamics of scalar fields in an effective theory. Such scalar fields are often termed moduli and their values describe the shape and size of the compact space. The evolution of scalar fields in cosmology is a well established area of study, finding applications in inflationary model building, quintessence and many others. Indeed, cosmologists seem able to solve most problems with the introduction of a new scalar field. A particularly attractive system of scalars is the one where the potential is formed from a series of exponential terms, such a system can be phrased in the form of an autonomous dynamical system whose critical points allow for a simplified description of the evolution \\cite{Copeland:1997et,Collinucci:2004iw}. Fortunately such a nice description does not go to waste as exponential potentials are common in unified gravity models with exponentials coming from dimensional reduction, gaugino condensation and instanton corrections \\cite{Derendinger:1985kk,Dine:1985rz,Witten:1996bn,Moore:2000fs}.  The critical points of the autonomous system formed in cosmology using scalars with exponential potentials reveals that {\\it tracking solutions} are possible, and indicate that {\\it assisted behaviour} can occur. By tracking solution we mean that the scalars evolve in such a way that their energy density remains at a constant fraction of the total energy density of the Universe, with the rest of the energy density being composed typically of a barotropic fluid. One may also have {\\it scaling} behaviour, where the energy density of the field evolves in proportion to $H^2$ without reference to other matter. Assisted behaviour refers to the behaviour of multiple scalars, where their combined effect can be such that the Universe expands more rapidly than one may naively expect by looking at the individual terms in the potential.  In this paper we aim to study dimensional reduction on homogeneous manifolds, namely those formed as a coset of compact Lie groups G/H. In order to be concrete we shall take eleven dimensional supergravity to be our starting point, whose bosonic field content comprises a metric and a three-form potential. As we are aiming for a four dimensional effective theory we shall reduce on the seven dimensional cosets which have been classified in \\cite{Castellani:1983yg}. Such coset reductions are familiar in the supergravity literature, where the supersymmetric stationary points are well known. Here however we concern ourselves with the full effective potential and the dynamics that follow, paying particular attention to the regimes which lead to tracking or scaling behaviour.  We shall structure the paper by first introducing the bosonic action of eleven dimensional supergravity, along with the ansatz for dimensional reduction. Having done that we shall produce the effective potentials for all the 7D cosets classified in \\cite{Castellani:1983yg}. With these in place we are able, after a brief introduction to scaling solutions, to study the scaling properties of supergravity reduced on coset manifolds. Throughout the paper we shall be using technical results for dimensional reduction and for cosets, the details of which can be found in the appendices. We end the paper with our concluding remarks, and comments for future work. ",
        "conclusions": "\\label{sec:conclusions}  In this paper we have given an explicit construction of the consistent dimensional reduction of eleven dimension supergravity over the homogeneous spaces classified in \\cite{Castellani:1983yg}. By including a Freund-Rubin four-form flux we have seen that this leads to an effective theory in four dimensions that consists of Einstein gravity coupled to a series of moduli fields, and that these fields experience a potential which is composed of a series of exponentials. Motivated by the large body of material which studies gravity+scalars with exponential potentials we studied the possibility of scaling solutions in this setup. However, for exact scaling solutions to exist there can be at most the same number of terms in the potential as there are fields; if this is violated then one finds an over constrained algebraic system for which there is no solution. In all the cosets of the classification we find that there are more terms in the potential than there are moduli fields. While this rules out the existence of exact scaling solutions, one still has the possibility that there are regions in field space where only a subset of terms are relevant, thereby allowing approximate scaling solutions. Having given a set of criteria for such regions to exist we showed that it is possible to get approximate scaling in each of the coset examples. The cases which we analysed explicitly were only the ones where there is a single dominant term, and we explicitly gave the exponent of the effective potential. We also noted that there are examples where more than one term could dominate, giving the possibility of assisted behaviour in these directions; we shall leave the detailed study of such examples for future work.  Aside from the Freund-Rubin flux that we included, there is also the possibility that one may include internal flux. This internal flux may take the form of a combination of axions coming from three-form potentials as $C_3\\sim\\phi(x)c_3(y)$, and non-exact flux, i.e. a four form flux that cannot be written globally as d$c$. In such a scheme, the consistent truncation would require us to take fluxes that matched the symmetries of the coset involved and would lead to a finite, manageable set of scalar fields. This would then lead to a four dimensional effective theory containing axions and geometrical moduli. Both of these types fields are to be expected in such dimensional reduction schemes and it remains to be seen how the presence of such axions in coset models affects the dynamics.   \\vspace{1cm} \\noindent {\\large\\bf Acknowledgements} The authors would like to thank Beatriz de Carlos, Mark Hindmarsh, Thomas House and Eran Palti for useful discussions. Both JLPK and PMS are supported by PPARC.  \\vskip 1cm"
    },
    "0601/astro-ph0601140_arXiv.txt": {
        "abstract": "Recent models for the emission  clouds within the Broad Line Region of quasars suggest that they are due to transient overdensities within an overall  turbulent medium.   If this  were  the case,  the broad  line emission  would spatially  appear fractal,  possessing structure  on a range of  scales. This  paper examines the  influence of  such fractal structure when a quasar is  microlensed by a population of intervening masses. It is found that  while the highest fractal levels can undergo significant  microlensing magnification, when  these light  curves are superimposed  to  create  an  emission  line  profile,  the  resultant emission line profile remains  relatively constant for physical models of the  Broad Line Region. It  is concluded that the  detection of the possible  fractal structure  of Broad  Line Regions  via gravitational microlensing is not practical. ",
        "introduction": "While quasars are  amongst the most luminous objects  in the Universe, the  majority of  their emission  is  generated within  a region  only parsecs in extent. At cosmological scales, such regions are well below the resolving power  of even the most powerful  telescopes, and so the spatial  structure  within  these  inner regions  cannot  be  directly observed and must be inferred by other, more indirect means.   In  recent years,  a general  picture for  quasar central  regions has emerged, with  a hot accretion  disk surrounding a  supermassive black hole   being  responsible   for  the   extensive   continuum  emission characteristic  of quasars. Beyond  this, on  the scale  of $\\sim1pc$, this  accretion disk  is surrounded  by a  large population  of clouds which  are illuminated  by the  central engine  and produce  the broad emission   lines,   with  widths   of   thousands   of  $\\kms$,   also characteristic  of  quasar spectra\\footnote{It  should  be noted  that alternative models  for the source  of the broad line  emission exist, including the accretion disk wind model of \\citet{1995ApJ...451..498M} which occurs on a considerably smaller scale.}.  Initial studies  of the Broad Line  Region (BLR) were  based on simple ionisation models, with predicted BLR sizes of 0.1 to a few parsecs in extent \\citep[e.g.][]{1979RvMP...51..715D}.  Recent studies, utilising reverberation mapping to  measure the scale of the  BLR, however, have found  evidence for a  smaller BLR,  more than  an order  of magnitude smaller       than      suggested      by       ionisation      models \\citep{1985ApJ...298..283P}.  Furthermore, \\citet{1999ApJ...526..579W} and \\citet{2000ApJ...533..631K} have demonstrated that the size of the BLR scales with quasar  luminosity, with $R_{BLR}\\propto L^{0.7}$, and BLRs were also found to possess significant ionisation stratification, with high-ionisation lines  arising in a region an  order of magnitude smaller than low-ionisation lines.  The  origin of  the BLR  has  proved problematic  to understand,  with suggestions  that  the emitting  material  could  be  the debris  from star-disk    collisions     at    the    heart     of    the    quasar \\citep{1994ApJ...434...46Z,1996ApJ...470..237A},  supernova explosions within                         quasar                         outflows \\citep{2001A&A...375..827P,2003A&A...408...79P},  or even  the bloated atmospheres of irradiated stars \\citep{1997MNRAS.284..967A}.  However, the  BLR must  possess  significant substructure,  with line  emission arising  from  dense  clouds   embedded  in  a  lower  density  medium \\citep{1981ApJ...245..396C},  with the  smoothness  of emission  lines suggesting  that  there  must  be  of  order  $10^8$  emitting  clouds \\citep{1998MNRAS.297..990A}, effectively ruling out the picture of the BLR being  composed of  a relatively small  number of large  clouds or bloated   stars.   Nevertheless,  such   a  situation   is  physically problematic, as  these clouds  should rapidly dissipate,  with various mechanisms  suggested  for their  confinement  [e.g.  magnetic  fields \\citet{1987MNRAS.228P..47R}].  Recently,  \\citet{2001ApJ...549..118B}  suggested  that  the  `clouds' within the BLR are not isolated, individual entities embedded within a confining  medium, but  rather represented  transient knots  of higher density within an overall turbulent  BLR. One conclusion of this study was  that the  BLR would  not  appear as  a smoothly-varying  emitting region. Instead, the  turbulent nature of the BLR  would result in the spatial  emission  from  the  region  possessing  an  overall  fractal structure.  Gravitational  microlensing  provides  an  opportunity  to  `see'  the structure in  the central  regions of quasars,  utilising differential magnification    effects    \\citep{1991AJ....102..864W}.     Recently, \\citet{fractal}  demonstrated   that  small-scale  fractal  structures within  quasars, arising  in an  x-ray emitting  hot corona  above the accretion disk, imprint themselves on the light curve of a microlensed quasar.  Hence, this paper  considers the influence of microlensing on an extended, fractal BLR.  However, as this is more extensive than the X-ray emitting  region, it is important  to compare it  to the natural scale length for gravitational  microlensing; the Einstein radius.  In the source plane, this is given by \\begin{equation} \\eta = \\sqrt{ 4 \\frac{ G M }{ c^2 } \\frac{D_{ls} D_{os}}{D_{ol}}} \\label{eqn1} \\end{equation} where  $M$ is  the mass  of the  microlensing body,  and  $D_{ij}$ are angular diameter distances between  the observer $(o)$, lens $(l)$ and source  $(s)$. For typical  cosmological lensing  configurations, with microlensing  stars with  mass $M\\lta  1\\msun$, this  scale  length is $\\eta\\lta 0.1{\\rm pc}$, and  is often substantially smaller. The X-ray emitting  region considered  in \\citet{fractal}  is small  compared to this  scale  and   hence  significant  microlensing  magnification  is expected.  As  discussed in \\citet{2004MNRAS.348...24L},  however, the extensive nature of  the BLR ensures it must  lie across a substantial portion of  the complex caustic network  that is seen  at high optical depths, and the effects of  microlensing cannot be approximated by the influence  of individual  caustic structures.   The structure  of this paper  is  as follows;  Section~\\ref{method}  describes the  numerical approach adopted in  this study, while Section~\\ref{results} discusses the  resultant   microlensing  light  curves   and  their  statistical properties.   The   conclusions  to   this  study  are   presented  in Section~\\ref{conclusions}.  \\begin{figure} \\centerline{ \\psfig{figure=fig1.ps,angle=270,width=3.0in}} \\caption[]{A   realisation  of   a   fractal  BLR   as  described   in \\citet{2001ApJ...549..118B}. Clearly clumps and subclumps can be seen, but the smallest structures  in this picture ($\\sim10^{-6}$pc) are too small to be seen by eye. \\label{fig1}} \\end{figure} ",
        "conclusions": "The nature of the broad  line emission regions of quasars has remained a  subject of discussion  for a  number of  years, with  recent models suggesting that the emitting clouds are temporary density enhancements in an overall  turbulent medium.  Such a scenario  predicts that these density enhancements should have a fractal distribution of structures. This   paper   has  investigated   the   influence  of   gravitational microlensing  on  this   extensive  fractal  structure,  finding  that different   levels  of   the  fractal   structure   undergo  differing magnifications,  with   the  smaller  sections   of  the  substructure suffering stronger  magnification. It  was found, however,  that while individual clouds  in the  fractal hierarchy displayed  quite dramatic variability due to microlensing,  the combination of light curves from the entire population of clouds  within the BLR resulted in an overall constant light curve.  While the  overall magnification  of the BLR  appears to  be constant, more significant magnification of substructures could be apparent when examining the  form of the  emission line profile,  i.e. substructures within  the  BLR  could  possess  coherent velocities  and  hence  may contribute  to a particular  velocity within  the emission  line. This paper  considered  a  simple  Keplerian  velocity  structure  for  the emission line  clouds, assigning a particular velocity  to all fractal hierarchies below  a particular level.  The result  of this procedure, however, revealed that coherent velocity structure cannot apply to the lowest levels of the fractal  hierarchy as the resultant emission line profile is  clearly too  structured, although dramatic  variability is seen throughout the line. Assigning coherent velocity structure to the higher  fractal hierarchy  does smooth  out the  form of  the emission line, but it also smooths out the light curves for each velocity bin. From this it is possible  to conclude that while individual structures within  a BLR  are being  microlensed and  are  undergoing significant magnifications, the  extensive nature of  the BLR and  the requirement that the resulting emission  line appears relatively smooth means that microlensing is  unlikely to reveal the putative  fractal structure of quasar BLRs.     \\newcommand{\\aap}{A\\&A} \\newcommand{\\apj}{ApJ} \\newcommand{\\apjl}{ApJ} \\newcommand{\\aj}{AJ} \\newcommand{\\mnras}{MNRAS} \\newcommand{\\apss}{Ap\\&SS} \\newcommand{\\nat}{Nature}"
    },
    "0601/astro-ph0601189_arXiv.txt": {
        "abstract": "{We report on diffuse X-ray emission associated with the nearby solitary millisecond pulsar \\PSR\\ detected with XMM-Newton and Chandra. The emission extends from the pulsar to the northwest by $\\sim 0.5$ arcmin. The spectrum of the nebular emission can be modeled with a  power-law of photon index $2.2\\pm0.4$, in line with the  emission originating from accelerated particles in the post shock flow. For PSR J0437$-$4715, PSR J0030+0451 and PSR J1024$-$0719, which all have spin parameters comparable to that of \\PSR, no diffuse emission is detected down to a $3-\\sigma$ limiting flux of $\\sim 4-7\\times10^{-15}$ ergs s$^{-1}$ cm$^{-2}$. ",
        "introduction": "Diffuse plerionic emission is often a special marker of young and powerful pulsars if their spin-down energy is in excess of $\\sim 10^{36}$ erg/s. Extended diffuse emission associated with older and less powerful pulsars, however, is rare and so far only seen in Geminga (Caraveo et al.~2003), PSR B1929+10 (Becker et al.~2005), PSR B2224+65 (Romani et al.~1997; Chatterjee \\& Cordes 2002) and the millisecond pulsars (MSPs) PSR B1957+20 (Stappers et al.~2003), J0437-4715 (Bell et al.~1995), and J2124$-$3358 (Gaensler, Jones \\&  Stappers 2002). While all three MSPs have bright bow-shock nebulae detected in  $H_\\alpha$, diffuse X-ray emission associated with it could be detected only from the black widow pulsar PSR B1957+20 in a recent Chandra observation (Stappers et al.~2003).  Comparable deep observations of almost all X-ray bright MSPs, however,  have been performed by XMM-Newton in previous years. These XMM-Newton data thus provides us a valuable data base to search for diffuse and extended emission components from these MSPs as well.  In this letter we report on a search for diffuse X-ray emission to be associated with the MSPs \\PSR, J0437-4715, J0030+0451 and J1024-0719. All these pulsars have comparable spin parameters (cf.~Table 1) so that differences in their emission properties are most likely cause by differences in the pulsar-ISM interaction/local environment rather than by differences in their total energy output.  PSR J2124-3358 and J1024-0719 were both discovered during the Parkes 436 MHz survey of the southern sky (Bailes et al.~1997). PSR J0030$+$0451 was discovered at 430 MHz during the Arecibo Drift Scan Search (Somer 2000) and independently in the Bologna sub-millisecond pulsar survey (D'Amico 2000), whereas PSR J0437-4715 was discovered in the Parkes southern sky survey (Johnston et al.~1993). PSR J0437-4715 was the first MSP of which pulsed X-ray emission was detected (Becker \\& Tr\\\"{u}mper 1993). X-ray emission  from \\PSR\\ and J1024-0719 was reported by Becker \\& Tr\\\"{u}mper (1998, 1999) in ROSAT HRI data whereas the X-ray counterpart of PSR J0030$+$0451 was discovered in the final ROSAT PSPC observation (Becker et al.~2000). A $H_\\alpha$ bow shock is seen around PSR J0437-4715 (Bell et al.~1995). However, its X-ray counterpart was not detected by ROSAT and Chandra (Becker \\& Tr\\\"umper 1999, Zavlin et al.~2002). Gaensler, Jones \\&  Stappers (2002) discovered an $H_\\alpha$-emitting bow shock nebula around PSR J2124-3358. This bow shock is very broad  and highly asymmetric about the direction of the pulsar's proper motion. The asymmetric shape might be caused by a significant density gradient in the ISM, bulk flow of ambient gas and/or anisotropies in the pulsar's relativistic wind (Gaensler, Jones \\&  Stappers 2002). Observation of PSR J0030+0451, performed with the ESO NTT in La Sila, did not detect diffuse $H_\\alpha$ emission associated with it (A.~Pellizzoni priv.~com.).  \\begin{table*} \\centering \\caption{Pulsar Parameters of PSRs J0030+0451, J2124-3358, J1024-0719 and J0437-4715 (from Manchester et al.~2005) \\label{ephemeris}} \\begin{tabular}{lcccc} \\hline\\hline Pulsar                       & PSR J0030+0451 & PSR J2124-3358 & PSR J1024-0719 & PSR J0437-4715\\\\ \\hline Right Ascension (J2000)  & $00^{\\rm h} 30^{\\rm m} 27.432^{\\rm s}$ & $21^{\\rm h} 24^{\\rm m} 43.862^{\\rm s}$ & $10^{\\rm h} 24^{\\rm m} 38.700^{\\rm s}$ & $04^{\\rm h} 37^{\\rm m} 15.787^{\\rm s}$\\\\ Declination (J2000)          & $+04^\\circ\\; 51'\\; 39.7\"$ & $-33^\\circ\\; 58'\\; 44.257\"$ & $-07^\\circ\\; 19'\\; 18.915\"$ & $-47^\\circ\\; 15'\\; 08.462\"$\\\\ Pulsar Period, $P$ (ms)        & 4.865453207369  & 4.9311148591481 & 5.1622045539088 & 5.7574518310720  \\\\ Period derivative $\\dot{P}$ ($10^{-20}$ s s$^{-1}$) & 1.0 & 1.33 & 1.85 & 1.87 \\\\ Age ($10^{9}$ yrs)          & 7.71 & 5.86 & 4.41 & 4.89\\\\ Surface dipole magnetic field ($10^{8}$ G) & 2.23 & 2.60 & 3.13 & 3.32 \\\\ Epoch of Period (MJD)      &  50984.4 & 50288.0 & 51018.0 & 51194.0 \\\\ Dispersion Measure (pc cm$^{-3}$)  & 4.3328  & 4.6152 & 6.491 & 2.6469 \\\\ Dispersion based distance (pc)   &  230 & 250 & 350 & 140            \\\\ Spin-down Luminosity ($10^{33}$) ergs s$^{-2}$ & 3.43 & 4.38 & 5.3 & 3.87\\\\ \\hline \\end{tabular} \\end{table*} ",
        "conclusions": "Basic characteristics of the X-ray bow shocks detected around rotation-powered pulsars are summarized in Table 3. \\PSR\\ is a new candidate to this group. Its tail is found to have an extent of $\\sim 0.5$ arcmin. Adopting the distance of $\\sim$ 250 pc, the tail has a length of $l\\sim1.1 \\times10^{17}$ cm. For the pulsar's proper motion velocity of 58 km s$^{-1}$ (Manchester et al.~2005), the timescale, $t_{\\rm flow}$, for the passage of the pulsar over the length of its X-ray trail is estimated to be $\\sim600$ yrs.  According to the discussion in Becker et al.~(2005) on the trail emission of PSR B1929+10 we estimate the magnetic field in the shocked region by assuming $t_{\\rm flow}$ to be comparable to the electron lifetime of the synchrotron emission. This yields $\\sim30\\mu$G for the inferred magnetic field strength in the emitting region. The magnetic field strength in the ISM is estimated to be $\\sim2-6\\mu$G (cf.~Beck et al.~2003 and references therein). Taking into account that the magnetic field in the termination shock might be compressed (e.g.~Kennel \\& Coroniti 1984), our estimation is approximately consistent if the compression factor is $\\sim 7$.   \\begin{table*} \\begin{center} \\caption{Properties of rotation-powered pulsars with X-ray/$H_\\alpha$ bow shocks.} \\begin{tabular}{ccccccccc} \\hline\\hline PSR & Other Name/ & \\.{E} & $L_{X}^{diff}$ (0.5$-$10 keV)$^{a}$ & $\\tau_{spin-down}$ & $H_\\alpha$ & X-ray & Radio & Reference\\\\ & associated SNR &erg s$^{-1}$ &erg s$^{-1}$ & years & bow shock & trail & nebula & \\\\ \\hline \\multicolumn{9}{c}{Young pulsars }\\\\ \\hline J0537-6910 & LMC: N157B &$4.8\\times10^{38}$ & $3.5\\times10^{36}$  & $5.0\\times10^{3}$  & - & d & - & 1,2  \\\\ B1757-24   & Duck &$2.6\\times10^{36}$ & $4.6\\times10^{32}$ &  $1.6\\times10^{4}$  & - & d & d & 3        \\\\ B1853+01   & W44 &$4.3\\times10^{35}$ & $5.7\\times10^{32}$ & $2.0\\times10^{4}$  & - & d & d & 4        \\\\ J1747-2958 & Mouse &$2.5\\times10^{36}$ & $5.3\\times10^{34}$ & $2.6\\times10^{4}$   & - & d & d & 5        \\\\ B1951+32   & CTB80 &$3.7\\times10^{36}$ & $2.5\\times10^{33}$ & $1.1\\times10^{5}$   & d   & d & d & 6,7,8      \\\\ \\hline \\multicolumn{9}{c}{Middle age and old pulsars}\\\\ \\hline B0740-28   & &$1.4\\times10^{35}$ & -  &  $1.6\\times10^{5}$  & d & - & - & 9 \\\\ B0633+17   & Geminga &$3.2\\times10^{34}$ & $8.9\\times10^{28}$ &  $3.4\\times10^{5}$ & - & d & - & 10       \\\\ B2224+65 & Guitar &$1.2\\times10^{33}$ & $3.7\\times10^{31}$  & $1.1\\times10^{6}$  & d & ? & - & 11,12            \\\\ B1929+10   &  &$3.9\\times10^{33}$ & $8.3\\times10^{29}$ &  $3.1\\times10^{6}$     & - & d & ? & 13        \\\\ \\hline \\multicolumn{9}{c}{Millisecond pulsars}\\\\ \\hline B1957+20   &  &$1.1\\times10^{35}$ & $1.9\\times10^{31}$ &  $2.2\\times10^{9}$   & d & d & - & 14       \\\\ J0437-4715 & &$3.9\\times10^{33}$ & -  & $4.9\\times10^{9}$  & d & - & - & 15,16 \\\\ J2124-3358 &  &$4.4\\times10^{33}$ & $10^{29}$  &  $5.9\\times10^{9}$  & d & d & - & 17,18     \\\\ \\hline \\end{tabular} \\\\ \\end{center} {\\bf a.} The X-ray luminosities of the diffuse emission are taken from the corresponding reference and recalculate into energy band of 0.5$-$10 keV for easy comparison. References: (1) Wang \\& Gotthelf 1998; (2) Wang et al. 2001; (3) Kaspi et al. 2001; (4) Petre et al. 2002; (5) Gaensler et al. 2004; (6) Migliazzo et al. 2002; (7) Moon et al. 2004; (8) Li et al. 2005; (9) Jones et al. 2002; (10) Caraveo et al. 2003; (11) Romani et al. 1997; (12) Chatterjee \\& Cordes 2002; (13) Becker et al. 2005; (14) Stappers et al. 2003; (15) Bell et al. 1995; (16) Zavlin et al. 2002; (17) Gaensler et al. 2002; (18) This letter; and Manchester et al. 2005 otherwise. \\end{table*}  Following Becker et al.~(2005), we apply a simple one zone model (Chevalier 2000) to estimate the spectral behavior and the X-ray luminosity of the nebular emission.  The X-ray luminosity and spectral index depend on the inequality between the characteristic observed frequency $\\nu^{\\rm obs}_{X}$ and the electron synchrotron cooling frequency $\\nu_{\\rm c}$ which is estimated to be $1.6\\times10^{17}$ Hz. Since in general $\\nu^{\\rm obs}_{X}>\\nu_{\\rm c}$, we concluded that the emission is in a fast cooling regime. Electrons with the energy distribution, $N(\\gamma)\\propto\\gamma^{-p}$, are able to radiate their energy in the trail with photon index $\\alpha=(p+2)/2$. The index $p$ due to shock acceleration typically lies between 2 and 3 (cf. Cheng, Taam, \\& Wang 2004 and references therein). Taking $p=2.35$ yields $\\alpha^{\\rm th}\\simeq2.2$ which is in accordance with the result from the observed value $\\alpha^{\\rm obs}=2.2\\pm0.4$. Assuming the energy equipartition between the electron and proton (Cheng, Taam, \\& Wang 2004), we take the fractional energy density of electron $\\epsilon_{\\rm e}$ to be $\\sim0.5$ and the fractional energy density of the magnetic field $\\epsilon_{B}$ to be $\\sim0.01$. Assuming a number density of ISM to be 1 cm$^{-3}$, the distance of the shock from the pulsar is estimated to be $\\sim3.6\\times10^{16}$ cm. With these estimates, the calculated luminosity, $\\nu L_{\\nu}$, is given as $\\sim10^{29}$ ergs s$^{-1}$ which is well consistent with the observed values of $1.3\\times10^{29}$ ergs s$^{-1}$ (0.1-2.4 keV) and $8.9\\times10^{28}$ ergs s$^{-1}$ (0.5-10 keV).  Although the general properties of the X-ray trail in \\PSR\\, are not in contradiction with properties observed in other pulsars there are still ambiguities which are not completely resolved. First, one should notice that the trail is misaligned with the direction of the pulsar's proper motion. As reported by Gaensler, Jones, \\& Stappers (2002), the head of the $H_\\alpha$ bow shock is found to be highly asymmetric about the pulsar's velocity vector with the apparent nebular symmetry axis deviated from the velocity vector by $\\sim30^{\\circ}$. Even though the misalignment of the X-ray trail seems to agree with the asymmetry of the $H_\\alpha$ nebula, deeper observations by XMM-Newton and Chandra are required in order to constrain the physical properties of this interesting nebula in higher detail."
    },
    "0601/astro-ph0601376_arXiv.txt": {
        "abstract": "This is the first part of an \\ha\\, kinematics follow-up survey of the SINGS sample. The data for \\NN\\, galaxies are presented. The observations were done on three different telescopes with \\FM, an integral field photon counting spectrometer, installed in the respective focal reducer of each telescope. The data reduction was done through a newly built pipeline with the aim of producing the most homogenous data set possible. Adaptive spatial binning was applied to the data cubes in order to get a constant signal-to-noise ratio across the field of view. Radial velocity and monochromatic maps were generated using a new algorithm and the kinematical parameters were derived using tilted-ring models. ",
        "introduction": "The Legacy survey SINGS (Spitzer Infrared Nearby Galaxies Survey) wants to characterise the infrared emission across the entire range of galaxy properties and star formation environments, including regions that until now have been inaccessible at infrared wavelengths (\\citeauthor{2003PASP..115..928K}, \\citeyear{2003PASP..115..928K}). SINGS will provide: \\begin{itemize} \\item new insights into the physical processes connecting star formation to the ISM properties of galaxies; \\item a vital foundation of data, diagnostic tools, and astrophysical inputs for understanding SPITZER observations of the distant universe and ultraluminous and active galaxies; \\item an archive that integrates visible/UV and IR/submillimeter studies into a coherent self-consistent whole, and enables many follow-up investigations of star formation and of the ISM. \\end{itemize}  The SPITZER observations will provide images in 7 different bands from 3.6\\um to 160\\um and spectroscopic data at medium and low resolution in the range 5--95\\um. These data will be used to trace the distribution and content of different dust components, from the PAHs and very small grains in the mid-IR, to the big grains in the far-IR (\\citeauthor{2003MNRAS.346..403D} \\citeyear{2003MNRAS.346..403D}, \\citeauthor{2003PASP..115..928K} \\citeyear{2003PASP..115..928K}). Ancillary multiwavelength observations will provide images in X-rays, UV (1300--2800 \\AA\\ imaging and spectrophotometry), BVRIJHK, \\ha, Pa$\\alpha$, FIR, submillimeter, CO and \\hI. A total of 20 ground- and space-based telescopes are providing supporting data.  These data will help understand the process of star formation and feedback mechanisms that are fundamental parameters regulating the formation and evolution of galaxies. History of star formation has been strongly different for galaxies of different morphological type and luminosity. While short events of star formation, probably triggered by violent merging, formed most stars in elliptical galaxies, late-type systems seem to have their star formation modulated by the angular momentum (\\citeauthor{1986ARA&A..24..421S}, \\citeyear{1986ARA&A..24..421S}) or by the mass of the initial system (\\citeauthor{2001ApSSS.277..401B}, \\citeyear{2001ApSSS.277..401B}). The process of star formation and feedback must thus be clearly understood in order to understand galaxies' evolution. However, these physical processes are still poorly known. The primordial atomic gas has to condense into molecular clouds to form stars. The newly formed stars inject metals into the interstellar medium via stellar winds, heat the dust and ionise the surrounding gas. It seems that the activity of star formation is regulated by the total gas surface density (\\citeauthor{1989ApJ...344..685K}, \\citeyear{1989ApJ...344..685K}), but it is still unclear what is the role of rotation in this process.  Even as important as it seems, no gathering of optical kinematical data was planned for the SINGS galaxies. This paper, by providing the \\ha\\, kinematics over the whole optical extent for \\NN\\, galaxies of the SINGS sample, wants to make up for this lack. A total of 58 SINGS galaxies are potentially observable in the \\ha\\, emission line (see section \\ref{sample_section}). The \\ha\\, kinematics of the 30 remaining galaxies of the observable part of the sample will be published in a forthcoming paper. These data were obtained with \\FM\\, on three different telescopes (see section \\ref{obs_runs_section}). \\FM\\, is an integral field spectrometer made of a photon-counting camera using a third generation photo-cathode, a scanning Fabry-Perot (FP) and a narrow-band interference filter. \\FM\\, was coupled to the focal reducer of the telescopes used. The photo-cathode used has a high quantum efficiency ($\\sim$ 30\\% at \\ha). This camera enables one to scan rapidly ($\\sim$ 5--10 minutes) the FP Free Spectral Range (FSR) and to cycle many times during an observation, thus averaging changing photometrical conditions, as compared to CCD observations where scanning must be done slowly to overcome the readout noise (details about the camera can be found in \\citeauthor{2003SPIE.4841.1472H}, \\citeyear{2003SPIE.4841.1472H} and \\citeauthor{2002PASP..114.1043G}, \\citeyear{2002PASP..114.1043G}). In this paper, section 2 gives an overview of the observational campaign and of the galaxies studied. Section 3 discusses how the data were reduced, processed and how the kinematical data and parameters were extracted. Section 4 provides all the maps extracted from the work done in section 3. Section 5 discusses the advantages of FP kinematical data as compared to other kinematical data. A short appendix is added to comment the observational characteristics of the galaxies presented. Once completed, the data set will be available in the SINGS database, as for the other SINGS ancillary surveys.   \\begin{figure} \\begin{center} \\includegraphics[width=0.5\\textwidth]{figures/galaxy_histo.eps} \\caption{The SINGS RC3 galaxy type distribution. The grey area shows the galaxies presented in this paper.} \\label{histo_figure} \\end{center} \\end{figure}  \\begin{figure*}% \\includegraphics[width=\\textwidth]{figures/great_map.eps} \\caption{The SINGS \\ha\\, kinematics sky coverage so far. The names in blue correspond to the observations done at the OMM, in green at the CFHT and in red at the ESO 3.6m telescope. Galaxies  relative sizes are to scale. The galactic equator is shown as the dotted line. Naturally, there is a strong concentration toward the Virgo cluster.} \\label{great_map_figure} \\end{figure*}  \\setcounter{table}{0} \\begin{table*} \\caption{Observational data for the SINGS \\ha\\, kinematics sample.} \\begin{tabular}{ccccccccccc} \\hline \\hline Galaxy  &  $\\alpha$(J2000)& $\\delta$(J2000) & Type & $\\Delta^{(1)}$ & $D_{25}^{b,i}$ $^{(2)}$ &  $B_T^{b,i}$ $^{(3)}$ &  $M_B^{b,i}$ $^{(4)}$ & Systemic Velocity$^{(5)}$\\\\ Name & (hh mm ss) & ($^o$ ' \") & RC3 & (Mpc) &  ($'$) &   &  & (\\kms)  \\\\ \\hline NGC 628   &   01 36 41.8   &   +15 47 00   &   SA(s)c   &   11.4   &  10.5 x 9.5  & 9.95            & $-$20.33      &     657\\\\ NGC 925   &   02 27 16.8   &   +33 34 41   &  SAB(s)d   &  9.3     &      11.2    & 10.6            & $-$19.24      &     554\\\\  % NGC 2403  &   07 36 54.5   &   +65 35 58   &  SAB(s)cd  &     4.2  & 21.4         & 8.5             & $-$19.61      &     132\\\\  % NGC 2798  &   09 17 22.9   &   +41 59 59   &  SB(s)a pec&   24.7   &   2.6 x 1    & 13.04           & $-$18.92      &    1726\\\\ NGC 2915  &   09 26 11.5   &   $-$76 37 35   &  I0        &   2.7    &   1.9 x 1    & 13.25           & $-$13.91      &     468\\\\ NGC 2976  &   09 47 15.4   &   +67 54 59   &   SAc pec  &    3.5   &   5.9 x 2.7  & 10.82           & $-$16.90      &       3\\\\ NGC 3049  &   09 54 49.6   &   +09 16 18   &  SB(rs)ab  &   19.6   &   2.2 x 1.4  & 13.04           & $-$18.42      &    1494\\\\ NGC 3031  &   09 55 33.2   &   +69 03 55   &   SA(s)ab  &    3.5   &  26.9 x 14.1 & 7.89            & $-$19.83      &     $-$34\\\\ UGC 5423  & 10 05 30.6   &   +70 21 52   &  Im        &    3.5   &   0.9 x 0.6    & 15.19           & $-$12.53      &     350\\\\ NGC 3184  &   10 18 17.0   &   +41 25 28   &  SAB(rs)cd &    8.6   &   7.4 x 6.9  & 10.36           & $-$19.31      &     592\\\\ NGC 3198  &   10 19 54.9   &   +45 33 09   &  SB(rs)c   &  14.5    &  7.8         & 11.1            & $-$19.70      &     660\\\\ % NGC 3521  &   11 05 48.6   &   $-$00 02 09   &  SAB(rs)bc &    9     &   11 x 5.1   & 9.83            & $-$19.94      &     805\\\\ NGC 3621  &   11 18 16.3   &   $-$32 48 45   &  SA(s)d    &   6.2    &   12.3 x 7.1 & 10.28           & $-$18.68      &     727\\\\ NGC 3938  &   11 52 49.4   &   +44 07 15   &  SA(s)c    &   12.2   &  5.4 x 4.9   & 10.90           & $-$19.53      &     809\\\\ NGC 4236  &   12 16 42.1   &   +69 27 45   &  SB(s)dm   &  2.2     & 16.4         & 9.5             & $-$17.21      &       2\\\\ % NGC 4321  &   12 22 55.2   &   +15 49 23   &  SAB(s)b   & 16.1     & 7.4          & 10.0            & $-$21.03      &    1590\\\\ % NGC 4536  &   12 34 27.1   &   +02 11 16   &  SAB(rs)bc &   25     &   7.6 x 3.2  & 11.16           & $-$20.82      &    1808\\\\ % NGC 4569  &   12 36 49.8   &   +13 09 46   &  SAB(rs)ab &   20     &   9.5 x 4.0  & 10.26           & $-$21.24      &    $-$235\\\\ % NGC 4579  &   12 37 43.6   &   +11 49 05   &  SAB(rs)b  &   20     &   5.9 x 4.7  & 10.48           & $-$21.18      &    1591\\\\ % NGC 4625  &   12 41 52.7   &   +41 16 25   &  SAB(rs)m pec&  9.5   &   2.2 x 1.9  & 12.92           & $-$16.97      &     609\\\\ NGC 4725  &   12 50 26.6   &   +25 30 06   &  SAB(r)ab pec& 17.1   &   10.7 x 7.6 & 10.11           & $-$21.05      &    1206\\\\ NGC 5055  &   13 15 49.3   &   +42 01 45   &  SA(rs)bc  &    8.2   &  12.6 x 7.2  & 9.31            & $-$20.26      &     504\\\\ NGC 5194  &   13 29 52.7   &   +47 11 43   &  SA(s)bc pec&   8.2   &  11.2 x 6.9  & 8.96            & $-$20.61      &     463\\\\ NGC 5398  &   14 01 21.5   &   $-$33 03 50   &  SB(rs)dm  &   15     &   2.8 x 1.7  & 12.78           & $-$18.10      &    1216\\\\ NGC 5713  &   14 40 11.5   &   $-$00 17 21   &  SAB(rs)bc pec & 26.6 &   2.8 x 2.5  & 11.84           & $-$20.28      &    1883\\\\ IC 4710   &   18 28 38.0   &   $-$66 58 56   &  SB(s)m    &   8.5    &   3.6 x 2.8  & 12.5            & $-$17.14      &     741\\\\ NGC 6946  &   20 34 52.0   &   +60 09 15   &  SAB(rs)cd & 5.5      &  14.9        & 7.92            & $-$20.78      &      46\\\\ % NGC 7331  &   22 37 04.1   &   +34 24 56   &   SA(s)b   &   15.7   &  10.5 x 3.7  & 10.35           & $-$20.63      &     816\\\\ \\hline \\cr \\end{tabular} \\begin{flushleft} $^{(1)}\\Delta$ : distance in Mpc, flow-corrected for H$_{0}$=70 \\kms$Mpc^{-1}$, as presented in \\citeauthor{2003PASP..115..928K} (\\citeyear{2003PASP..115..928K}).\\\\ $^{(2)}D_{25}^{b,i}$ : optical diameter at the 25 magnitude/arcsecond$^{2}$ in B, corrected for the effects of projection and extinction. Taken from the RC3.\\\\ $^{(3)}B_T^{b,i}$ : corrected total apparent magnitude in B. Taken from the RC3.\\\\ $^{(4)}M_B^{b,i}$ : corrected total absolute magnitude in B. Calculated from $\\Delta$ and $B_T^{b,i}$.\\\\ $^{(5)}$Systemic velocity : galaxy's systemic velocity. Taken from \\citeauthor{2003PASP..115..928K} (\\citeyear{2003PASP..115..928K}).\\\\ \\end{flushleft} \\label{sample_table} \\end{table*}  \\begin{center} \\begin{table*} \\caption{Journal of the Fabry-Perot Observations.} \\label{JOFP_table} \\begin{tabular}{c|c|ccc|cc|cccc|cc} \\hline \\hline Galaxy &  Date & \\multicolumn{3}{c}{Filter}  & \\multicolumn{2}{c}{Effective integration} & \\multicolumn{4}{c}{Fabry-Perot} & \\multicolumn{2}{c}{Sampling} \\\\ Name  & & $\\lambda_c^{(4)}$ & FWHM$^{(5)}$ & T$_{max}^{(6)}$ &  t$_{exp}^{(7)}$ & t$_{Channel}^{(8)}$ & p$^{(9)}$ & FSR$^{(10)}$ & \\textbf{\\textit{F}}$^{(11)}$ & \\textbf{\\textit{R}}$^{(12)}$ & nch$^{(13)}$ & stp$^{(14)}$ \\\\ % &  & (\\AA) & (\\AA) & (\\%)  & (min.) &  (min) & & (\\kms) & & & & (\\AA) \\\\% & (\\Arc) \\\\ \\hline NGC 628$^{(1)}$  & 18/11/2003 &   6598   &  18.2  &  73  & 149  & 2.33 & 899 & 333.47 & 23.6  & 21216 & 64  & 0.11  \\\\ NGC 925$^{(1)}$  & 11/02/2002 &   6584   &  15    &  75  & 132  & 2.75 & 765 & 391.88 & 16  & 12240 & 48  & 0.18  \\\\ NGC 2403$^{(1)}$  & 17/11/2002 &   6569   &  10    &  50  & 120  & 3.00 & 765 & 391.88 & 14 & 10710 & 40  & 0.21  \\\\ NGC 2798$^{(2)}$  & 04/04/2003 &   6608   &  16.2  &  69  & 96   & 2.00 & 899 & 333.47 & 16  & 14384 & 48  & 0.15  \\\\ NGC 2915$^{(3)}$  & 04/21/2004 &   6581   &  19.8  &  60  & 149  & 4.67 & 609 & 492.27 & 14.3  & 8580 & 32  & 0.34  \\\\ NGC 2976$^{(1)}$  & 30/01/2003 &   6581   &  19.8  &  60  & 184  & 3.83 & 899 & 333.47 & 16.5  & 14834 & 48  & 0.15  \\\\ NGC 3049$^{(2)}$  & 08/04/2003 &   6598   &  18.2  &  73  & 144  & 3.00 & 899 & 333.47 & 17.4  & 15643 & 48  & 0.15  \\\\ NGC 3031$^{(1)}$  & 06/02/2003 &   6581   &  19.8  &  60  & 272  & 5.55 & 899 & 333.47 & 14.3  & 12856 & 48  & 0.15  \\\\ UGC 5423$^{(2)}$ & 06/04/2003 &   6565   &  15    &  40  & 120  & 2.50 & 899 & 333.47 & 18.5  & 16632 & 48  & 0.15  \\\\ NGC 3184$^{(1)}$  & 18/02/2004 &   6584   &  15.5  &  74  & 162  & 3.37 & 765 & 391.88 & 17.6  & 13464 & 48  & 0.18  \\\\ NGC 3198$^{(1)}$  & 06/03/2003 &   6584   &  15.5  &  74  & 260  & 5.00 & 899 & 333.47 & 23  & 20976 & 52 & 0.14  \\\\ NGC 3521$^{(1)}$  & 19/02/2004 &   6584   &  15.5  &  74  & 120  & 2.50 & 765 & 391.88 & 16.0  & 12240 & 48  & 0.18  \\\\ NGC 3621$^{(3)}$  & 04/20/2004 &   6584   &  15.5  &  74  & 128  & 4.00 & 609 & 492.27 & 14.4  & 8640 & 32  & 0.34  \\\\ NGC 3938$^{(1)}$  & 11/03/2004 &   6584   &  15.5  &  74  & 128  & 2.66 & 765 & 391.88 & 16.8  & 12852 & 48  & 0.18  \\\\ NGC 4236$^{(1)}$  & 27/02/2004 &   6581   &  19.8  &  60  & 182  & 3.50 & 899 & 333.47 & 23  & 20977 & 52 & 0.14 \\\\ NGC 4321$^{(1)}$  & 25/02/2003 &   6608   &  16.2  &  69  & 260  & 5.00 & 899 & 333.47 & 23  & 20977 & 52 & 0.14  \\\\ NGC 4536$^{(1)}$  & 14/03/2004 &   6598   &  18.2  &  73  & 163  & 3.40 & 765 & 391.88 & 21.3& 16295 & 48 & 0.18  \\\\ NGC 4569$^{(1)}$  & 11/03/2002 &   6569   &  15.0  &  60  & 152  & 3.80 & 765 & 391.88 & 20.5& 15682 & 40 & 0.21 \\\\ NGC 4579$^{(1)}$  & 04/04/2002 &   6598   &  10.0  &  60  & 92   & 2.30 & 609 & 492.27 & 19.6& 14494 & 40 & 0.27 \\\\ NGC 4625$^{(2)}$  & 06/04/2003 &   6581   &  19.8  &  60  & 120  & 2.50 & 899 & 333.47 & 15.9  & 14294 & 48  & 0.15  \\\\ NGC 4725$^{(1)}$  & 19/02/2004 &   6584   &  15.5  &  74  & 120  & 2.50 & 765 & 391.88 & 18.7  & 14305 & 48  & 0.18  \\\\ NGC 5055$^{(1)}$  & 14/03/2004 &   6584   &  15.5  &  74  & 128  & 2.66 & 765 & 391.88 & 16.8  & 12852 & 48  & 0.18  \\\\ NGC 5194$^{(1)}$  & 18/05/2003 &   6581   &  19.8  &  60  & 246  & 5.14 & 899 & 333.47 & 20.7  & 18609 & 48  & 0.15  \\\\ NGC 5398$^{(3)}$  & 04/10/2004 &   6598   &  18.2  &  73  & 149  & 4.66 & 609 & 492.27 & 12.0  & 7308 & 32  & 0.34  \\\\ NGC 5713$^{(3)}$  & 04/13/2004 &   6608   &  16.2  &  69  & 150  & 6.25 & 609 & 492.27 & 9.5  & 5785 & 24  & 0.45  \\\\ IC 4710$^{(3)}$   & 04/15/2004 &   6598   &  18.2  &  73  & 48   & 2.00 & 609 & 492.27 & 8.8  & 5359 & 24  & 0.45  \\\\ NGC 6946$^{(1)}$  & 19/11/2002 &   6569   &  10    &  50   & 120  & 2.00 & 765 & 391.88 & 14    & 10710 & 40 & 0.21  \\\\ NGC 7331$^{(1)}$  & 03/11/2002 &   6584   &  15.5  &  74  & 174  & 3.62 & 765 & 391.88 & 15.9  & 12164 & 48  & 0.18  \\\\ \\hline \\cr \\end{tabular}\\\\ \\begin{flushleft} $^{(1)}$ OMM : Observatoire du mont M\\'egantic, Qu\\'ebec, Canada. 1.6m telescope. \\\\ $^{(2)}$ CFHT : Canada-France-Hawaii Telescope, Hawaii, USA. 3.6m telescope.\\\\ $^{(3)}$ ESO : European Southern Observatory, La Silla, Chile, 3.6m telescope.\\\\ $^{(4)}$ $\\lambda_c$ : non tilted filter central wavelength at 20$^{\\circ}\\mathrm{C}$.\\\\ $^{(5)}$ FWHM : non tilted filter Full Width Half Maximum at 20$^{\\circ}\\mathrm{C}$.\\\\ $^{(6)}$ T$_{max}$ : non tilted filter maximum transmission at $\\lambda_c$ and at 20$^{\\circ}\\mathrm{C}$.\\\\ $^{(7)}$ t$_{exp}$ : total effective exposure time in minutes (Total exposure time * mean counting efficiency).\\\\ $^{(8)}$ t$_{channel}$ : total effective exposure time per channel in minutes (total exposure time per channel * mean counting efficiency).\\\\ $^{(9)}$ p : FP interference order at \\ha.\\\\ $^{(10)}$ FSR : FP Free Spectral Range at \\ha\\, in \\kms.\\\\ $^{(11)}$ \\textbf{\\textit{F}} : mean Finesse through the field of view.\\\\ $^{(12)}$ \\textbf{\\textit{R}} : spectral resolution ($\\Delta\\lambda/\\lambda$) according to the computed \\textit{Finesse}.\\\\ $^{(13)}$ nch : number of channels done by cycle in multiplex observations.\\\\ $^{(14)}$ stp : wavelength step in \\AA.\\\\ \\end{flushleft} \\end{table*} \\end{center} ",
        "conclusions": "\\label{section_conclusion} The \\ha\\, kinematics of \\NN\\, galaxies of the SINGS survey were presented in this paper. The observations were made with \\FM, an integral field FP spectrometer and a photon counting camera. The raw data obtained at the telescope were processed through a new pipeline, an adaptive binning algorithm has been applied to achieve optimal signal-to-noise ratio and the radial velocity maps were finally extracted from the data cubes using a selective intensity weighted mean algorithm. Kinematical parameters were computed using a tilted ring model and most of them agreed within an acceptable error range with the photometrical parameters, except for a few problematic galaxies. It has been shown that high spatial resolution data is essential for mapping the velocity gradient at the centres of galaxies. The advantages of integral field spectroscopy over long slit spectroscopy were also presented.  The aim of this paper is to provide accurate optical kinematical data for the galaxies of the SINGS survey. These data will be used in a forthcoming paper to present rotation curves and mass models of the non-barred galaxies (Nicol et al. in preparation). The data of some of the barred galaxies of the SINGS sample that are also part of the \\BB\\, survey were used to derive the bar pattern speeds using the Tremaine--Weinberg method (\\citeauthor{1984ApJ...282L...5T} \\citeyear{1984ApJ...282L...5T}) in a paper presented by \\citeauthor{2005ApJ...632..253H} (\\citeyear{2005ApJ...632..253H}). These data might also be used to allow for some of the galactic star formation models to include kinematical data and thus try to determine what is the exact role of rotation in the star formation processes on a galactic scale.  The \\ha\\, kinematical data for all the observable SINGS galaxies will be made available to the community as soon as all the galaxies have been observed."
    },
    "0601/astro-ph0601006_arXiv.txt": {
        "abstract": "The solar seismic waves excited by solar flares (``sunquakes'') are observed as circular expanding waves on the Sun's surface. The first sunquake was observed for a flare of July 9, 1996, from the Solar and Heliospheric Observatory (SOHO) space mission. However, when the new solar cycle started in 1997, the observations of solar flares from SOHO did not show the seismic waves, similar to the 1996 event, even for large X-class flares during the solar maximum in 2000-2002. The first evidence of the seismic flare signal in this solar cycle was obtained for the 2003 ``Halloween'' events, through acoustic ``egression power'' by Donea and Lindsey. After these several other strong sunquakes have been observed. Here, I present a detailed analysis of the basic properties of the helioseismic waves generated by three solar flares in 2003-2005. For two of these flares, X17 flare of October 28, 2003, and X1.2 flare of January 15, 2005, the helioseismology observations are compared with simultaneous observations of flare X-ray fluxes measured from the RHESSI satellite. These observations show a close association between the flare seismic waves and the hard X-ray source, indicating that high-energy electrons accelerated during the flare impulsive phase produced strong compression waves in the photosphere, causing the sunquake. The results also reveal new physical properties such as strong anisotropy of the seismic waves, the amplitude of which varies significantly with the direction of propagation. The waves travel through surrounding sunspot regions to large distances, up to 120 Mm, without significant decay. These observations open new perspectives for helioseismic diagnostics of flaring active regions on the Sun and for understanding the mechanisms of the energy release and transport in solar flares. ",
        "introduction": "It was suggested long ago \\citep{Wolff1972} that solar flares, giant explosions on the Sun, may cause acoustic waves traveling through the Sun's interior, similar to the seismic waves on the Earth. Because the sound speed increases with depth the waves are reflected in the deep layers of the Sun and appear back on the surface, forming expanding rings of the surface displacement. Theoretical modeling \\citep{Kosovichev1995} predicted that the speed of the expanding seismic waves increases with distance because the distant waves propagate into the deeper interior where the sound speed is higher. First observations of the seismic waves caused by the X2.6 flare of July 9, 1996 \\citep{Kosovichev1998}, proved these predictions. These observations also showed that the source of the seismic response was a strong shock-like compression wave propagating downwards in the photosphere. This wave was observed immediately after the hard X-ray impulse which produced by high-energy electrons hitting the low atmosphere. This led to a suggestion that the seismic response can be explained in terms of so-called ``thick-target'' models. In these models, a beam of high-energy is related to heating of the solar chromosphere, resulting in evaporation of the upper chromosphere and a strong compression of the lower chromosphere \\citep[e.g.][]{Kostiuk1975, Livshits1981, Fisher1985, Kosovichev1986}. This high-pressure compression produces a downward propagating shock wave \\citep{Kosovichev1986} that hits the solar surface and causes sunquakes. This shock observed in SOHO/MDI Dopplergrams as a localized large-amplitude velocity impulse of about 1 km/s or stronger represents the initial hydrodynamic impact resulting in the seismic response. In addition, \\citet{Kosovichev1998a} found that the seismic wave was anisotropic, with a significant quadrupole component.  The following observations of solar flares made by the Michelson Doppler Imager (MDI) instrument on the NASA-ESA mission SOHO did not show noticeable sunquake signals even for strong X-class flares. This search was carried out by calculating an ``egression'' power for high-frequency acoustic waves during the flares \\citep{Donea1999}. It became clear that sunquakes are a rather rare phenomenon on the Sun, which occurs only under some special conditions. Surprisingly, seven years later several flares did show strong ``egression'' signals indicating new potential sunquakes \\citep{Donea2005} (for a list see http://www.maths.monash.edu.au/\\~\\,adonea). It is interesting to note that the flare of July 9, 1996, was the last strong of the previous solar activity cycle, and the new strong sunquake events are observed in the declining phase of the current activity cycle after the maximum of 2000-2001. It appears that during the rising phase of the solar cycle and during its peak the solar flares are rather a ``superficial'' coronal phenomenon not affecting much the solar surface and interior. This could happen if the topology of magnetic field of solar active regions which produce flares changes in such a way that the magnetic energy is released at lower altitudes in the declining phase of the solar cycle than in the rising and maximum phase. Here, I present analysis of new observations of the seismic response to solar flares from the SOHO and RHESSI space observatories, which show that the sunquakes were indeed caused by the hydrodynamic impacts of high-energy electrons accelerated in solar flares confirming the initial result of \\citep{Kosovichev1998}, and determine basic properties of the flare-generated seismic waves by investigating their time-distance characteristics. ",
        "conclusions": "The new observations from SOHO and RHESSI provide unique information about the interaction of the high-energy particles accelerated in solar flares with solar plasma and the dynamics of the solar atmosphere during solar flares. These data also provide unique information about the interaction of acoustic MHD waves with sunspots, showing explicitly propagation of wave fronts through sunspot regions. This opens opportunity for developing new methods of helioseismology analysis of flaring active regions, similar to the methods of Earth-quake seismology.  \\clearpage"
    },
    "0601/astro-ph0601230_arXiv.txt": {
        "abstract": "We calculate the physical structure of protoplanetary disks by evaluating the gas density and temperature self-consistently and solving separately for the dust temperature. The effect of grain growth is taken into account by assuming a power-law size distribution and varying the maximum radius of grains $a_{\\rm max}$. In our fiducial model with $a_{\\rm max}=10\\mu$m, the gas is warmer than the dust in the surface layer of the disk, while the gas and dust have the same temperature in deeper layers. In the models with larger $a_{\\rm max}$, the gas temperature in the surface layer is lower than in the fiducial model because of reduced photo-electric heating rates from small grains, while the deeper penetration of stellar radiation warms the gas at intermediate height. A detailed chemical reaction network is solved at outer radii ($r \\ge 50$ AU). Vertical distributions of some molecular species at different radii are similar, when plotted as a function of hydrogen column density $\\Sigma_{\\rm H}$ from the disk surface. Consequently, molecular column densities do not much depend on disk radius. In the models with larger $a_{\\rm max}$, the lower temperature in the surface layer makes the geometrical thickness of the disk smaller, and the gaseous molecules are confined to smaller heights. However, if we plot the vertical distributions of molecules as a function of $\\Sigma_{\\rm H}$, they do not significantly depend on $a_{\\rm max}$. The dependence of the molecular column densities on $a_{\\rm max}$ is not significant, either. Notable exceptions are HCO$^+$, H$_3^+$ and H$_2$D$^+$, which have smaller column densities in the models with larger $a_{\\rm max}$. ",
        "introduction": "Planetary systems are thought to be formed in circumstellar disks, called protoplanetary disks, around young low-mass stars. The physical structure and evolution of the disks have been investigated by observation and by theoretical modeling. The disks are in hydrostatic equilibrium in the vertical direction, while the temperature is determined mainly by accretion at inner radii ($\\lesssim$ several AU) and by irradiation from the central star at outer radii. Hence the calculation of temperature is important in determining the physical structure of disks. In recent years several groups have been constructing disk models by solving the equations of radiation transfer. \\citet{dal98} obtained the vertical structure by solving one-dimensional radiation transfer at each disk radius, while \\citet{nom02} and \\citet{dul02} solved the disk structure using a two-dimensional radiation transfer code. Although these models differ quantitatively, they all flare up significantly --- the scale height is larger at larger radii. It should be noted that the density and temperature distributions are coupled; warm temperature in the outer radii makes the flared-up structure, which enhances the local irradiation (i.e. heating) from the central star. \\citet{dal99} pointed out that these model disks are geometrically too thick. In order to reproduce the observed spectral energy distribution and images of edge-on disks, some fraction of grains should be larger than interstellar dust \\citep{dal01}. This grain growth decreases the vertical height of the disk photosphere where the optical depth to the stellar radiation is unity, and thus the local heating rate by irradiation and hence the scale height. The presence of mm-sized grain is also required to account for the relatively high intensity of the millimeter dust continuum emission.  The theoretical models mentioned above, however, assumed that the gas temperature is equal to the dust temperature. Since the vertical density distribution is determined by gas-pressure gradient, the validity of this assumption should be checked. Recently \\citet{kd04} tackled this problem; thermal balance is solved to obtain gas temperature, while the density distribution and the dust temperature are fixed to those of \\citet{dzn02}, in which astronomical silicate grains of 0.1 $\\mu$m were assumed. It is found that dust and gas temperatures are equal to within 10\\% for the visual extinction of $A_{\\rm v} \\gtrsim 0.1$ mag, and that the gas temperature exceeds the dust temperature at larger heights. A similar result is obtained by \\citet{jfz04}, who calculated the gas temperature by using the density distribution and the dust temperature of \\citet{dal99}. \\citet{jfz04} also investigated effect of dust settling; the gas temperature is calculated with the dust abundance reduced by two orders of magnitude above certain heights. They found that the dust settling increases the amount of hot gas if PAHs remain well-mixed with the gas.  In the present paper, we calculate gas and dust temperature separately together with the gas density distribution in the disk. The gas temperature is obtained by assuming a simple chemical equilibrium and detailed thermal balance, including gas-dust heat exchange, at each position in the disk. Then the vertical density distribution is determined from hydrostatic equilibrium. The dust temperature is obtained by solving the two-dimensional radiation transfer \\citep{nom02}. We calculate the temperature and density profiles self-consistently by iteratively solving the equations. The effect of grain growth is investigated by assuming a power-law size distribution of grains and varying the maximum grain radius $a_{\\rm max}$ \\citep{mn93}.  We also calculate molecular distributions in the disk models obtained above by solving a detailed chemical reaction network. Several molecular lines, such as CO, HCO$^+$, HCN and CN are detected by radio observations \\citep{dut97, qi03, thi04}, which are sensitive to the region outside $\\sim 100$ AU because of the beam size ($\\gtrsim 1''$) of the current instruments. Although emission bands of CO and H$_2$O in inner hot region ($r\\lesssim 1$ AU) are detected in the infrared wavelength \\citep{ncm03,ctn04}, little constraints are given yet on the molecular abundance and their spatial distribution. Therefore we constrain our discussion of the detailed chemistry to radii $\\ge 50$ AU.  \\citet{aik02} and \\citet{zah03} calculated the molecular distribution using the disk model of \\citet{dal99}, in which interstellar-sized grains are well mixed with gas. The disk can be schematically divided into three layers. In the surface layer, molecules are dissociated by ultraviolet (UV) radiation and X-rays from the central star and the interstellar field, while heavy-element species are frozen onto grains in the midplane with high density and low temperature. Gaseous molecules observed at radio wavelengths are mostly present in the intermediate layer where the gas temperature and density are relatively high and the UV radiation is not too harsh. The model explains the observational results; the freeze-out in the midplane lowers the averaged abundances of heavy-element species and the UV radiation causes the high abundances of radical species. However, there are some quantitative disagreements. The models of \\citet{aik02} overestimate CO column density in DM Tau, and underestimate column densities of organic species in LkCa15. \\citet{zah03} somewhat improved the agreement by including scattering of stellar UV, which enhances the UV intensity and hence the photo-dissociation rate of CO within the disk. The dissociation of CO produces carbon atoms, which are transformed to other organic species. Therefore it is interesting to see how the molecular abundances are affected by grain growth, which will enhance the UV penetration into the disk.  The remainder of the paper is organized as follows.  A description of the physical and chemical model is given in \\S 2. Numerical results of the physical structure and the molecular distribution are presented in \\S 3. In \\S 4, we check the self-consistency of our model and discuss implications for molecular line observations. A summary is given in \\S 5. ",
        "conclusions": "\\subsection{Self consistency of the model} Chemistry and physical structure in disks are coupled; molecular abundances depend on physical parameters such as density, and line emission from certain species such as C$^+$ and O is the main coolant in the surface layers. It is, however, very time-consuming if we couple the calculation of physical structure with the detailed time-dependent chemical model. Instead, we have constructed a physical disk model adopting a chemical equilibrium equation of \\citet{nl97}, and then performed a detailed chemical network calculation. Here we check the self consistency of our models; the temperature and density distributions are re-evaluated using the molecular abundances obtained from the detailed chemical reaction network.  Figure \\ref{re_evaluate} shows the vertical distribution of density and temperature at radius of $r=201$ AU and time $t=1\\times 10^6$ yr in the model with $a_{\\rm max}= 1$ cm. The solid lines represent our original values, while the dotted lines are obtained using the molecular abundances from the detailed reaction network. In the surface layer ($z/r\\gtrsim 0.4$) the temperature is slightly higher than the original value; the self shielding of CO, which is not included in the simple equilibrium equation, lowers the C$^+$ abundance, which is a major coolant in the surface layer. The gas density in the surface layer is thus higher than the original value, which then slightly lowers the gas temperature at $z/r\\sim 0.3$.  It should be noted that the difference between the original and re-evaluated structures is small enough that the detailed chemical reaction network would give almost the same abundance distribution for the re-evaluated structure. Therefore we can conclude that our models of physical structure and molecular abundances are essentially self-consistent.  \\subsection{Effects of hydrodynamical processes} We did not include hydrodynamic motions such as turbulence within the disk. Such motion will have two major effects on chemistry: (i) it tends to smear the stratification of molecular abundances and (ii) transport of chemically active species and/or mother molecules will modify the chemical reaction network.  The former effect, (i), can be estimated by comparing the mixing and chemical timescales. Following \\citet{aik96}, the mixing timescale over distance $D$ can be written as \\begin{equation} t_{\\rm mix} \\approx \\frac{D^2}{l_{\\rm turb} v_{\\rm turb}}. \\label{eq:mix} \\end {equation} The equation stands for either radial mixing and vertical mixing, as long as the transport process is turbulence with typical eddy size of $l_{\\rm turb}$ and eddy speed of $v_{\\rm turb}$. For example, at radius of 200 AU, the midplane temperature is 10 K, and hence the sound speed is $\\sim 2\\times 10^4$ cm s$^{-1}$, and the scale height is $\\sim 36$ AU. If we set $l_{\\rm turb}$ and $v_{\\rm turb}$ to be 10 \\% of the scale height and sound velocity, respectively, it will take $\\sim 9\\times 10^4$ yr to mix the matter over the scale height ($\\sim 36$ AU). The chemical timescale is short $\\lesssim 10^4$ yr in the surface and midplane layers of the disk (\\S 2.2), while the chemistry in the intermediate molecular layer is almost in a pseudo-steady state at $10^5-10^6$ yr (\\S 2.2). Therefore the overall three-layer model cannot be smeared out by the mixing in the vertical direction, although the distribution within the molecular layer can be somewhat averaged and the molecular layer can be extended in the vertical direction. For example, \\citet{ddg03} observed two rotational transitions ($J=2-1$ and $1-0$) of $^{13}$CO in DM Tau, and obtained the gas temperature of 13 K, which is lower than the freeze-out temperature of CO ($\\sim 20$ K). The existence of such cold CO gas can be explained by the vertical mixing. Equation (\\ref{eq:mix}) indicates that at the radius of $\\sim 200$ AU CO gas can migrate over $\\sim 10$ AU in the vertical direction within $10^4$ yr, which is comparable to the freeze-out timescale in the midplane regions of DM Tau.\\footnote{The midplane density in DM Tau is lower than that in our disk model \\citep{dut97}.}  Accretion and radial mixing would not significantly modify our results, because the vertical distributions of molecular abundance at different radii are similar at least in the outer radius (\\S 2.2, \\S 3.2).  The latter effect, (ii), will enhance the gas-phase abundance of grain-surface products in the intermediate molecular layer; species such as H$_2$CO are formed on the grain surface in the midplane and are sublimated to the gas phase in the warm molecular layer (Willacy et al. in prep; Semenov et al. in prep).  \\subsection{Implications for molecular line observations} The derivation of molecular abundances from line intensities is not straightforward because of spatial variation of molecular abundances and physical condition (i.e. temperature and density). Hence a prediction of line intensities from model disks is important. Calculation of non-LTE 2D line transfer, however, is beyond a scope of the present work. Here we make some qualitative comments on molecular line intensities by comparing the spatial distribution of molecular abundances and physical conditions.  Firstly, we compare the distributions of gas density (Figure \\ref{rhotem_z}) and molecular abundances (Figure \\ref{abun_z}). At $z/r\\lesssim 0.5$, where various molecular species reach their peak abundance, the gas density $n_{\\rm H}$ is $\\gtrsim 10^6$ cm$^{-3}$. Hence most of the major molecular lines in millimeter wavelength can be excited. Secondly, distributions of abundances and temperature are compared. Because of vertical distribution of abundances, each molecular line and the dust continuum trace different layers of the disk. Significant amounts of CN, C$_2$H and CO exist in the upper layer in which gas and dust temperature is not coupled. As the maximum grain size becomes larger, the gas temperature in the surface layer decreases, which could be observable through these molecular lines. Species such as HCN and H$_2$O exist in the layer in which the gas and dust temperatures are mostly coupled, but are higher than their midplane values. On the other hand, a major fraction of dust exists in the midplane. Our results suggest that owing to the vertical temperature gradient the molecular lines can be observed even in the disks which are optically thick in the dust continuum. Although the detailed chemical network is solved only in the outer region ($\\ge 50$ AU) of small-mass disks in the present work, the vertical distribution of temperature and molecular abundances in the surface layers would be qualitatively similar in the inner regions and/or more massive disks."
    },
    "0601/astro-ph0601556_arXiv.txt": {
        "abstract": "Various approaches to quantum gravity suggest the possibility of  violation of Lorentz symmetry at very high energies. In these cases we expect a modification at low energies of the dispersion relation of photons that contains extra powers of the momentum suppressed by a high energy scale. These terms break boost invariance and can be tested even at relatively low energies. We use the light curves of the very bright short Gamma-Ray Burst GRB 051221A and compare the arrival times of photons at different energies with the expected time delay due to a modified dispersion relation. As no time delay was observed, we set a lower bound of $0.0066 \\, E_{pl} \\sim 0.66 \\cdot 10^{17}$ GeV on the scale of Lorentz invariance violation. ",
        "introduction": "Various quantum gravity theories suggest that Lorentz symmetry is broken or modified at very high energies \\cite{Amelino-Camelia:2002dx}. In most of these models the low energy limit of the photon dispersion relation is deformed by an addition of extra powers of the particle momentum. These modifications are suppressed by a high energy scale, beyond which Lorentz symmetry is broken. Among other implications, such terms cause the photon speed to depend on the energy and to differ from the classical speed of light, $c$ \\cite{Coleman:1998ti}.  Amelino-Camelia et al. \\cite{Amelino-Camelia:1997gz} proposed to use the arrival times of photons from cosmological Gamma Ray Bursts (GRBs) to set experimental bounds on the energy scale of Lorentz invariance violation. The idea was pursued by several groups which obtained observational bounds using ensembles of bursts with known redshfits \\cite{Ellis:2002in,Ellis:2005wr}, or using a single very powerful and atypical burst GRB 021206 \\cite{boggs-2004}. In this paper we use the very bright short burst GRB 051221A to test Lorentz symmetry and set a new bound on the possible scale of Lorentz symmetry violation. ",
        "conclusions": "We set new bounds on the possible energy scale of Lorentz invariance violation. To do so, we used the simultaneity of peak emission times in the light curves in different energy bands observed by Swift-BAT and Konus-Wind for GRB 051221A. The order of magnitude of the bounds depends on how the photon dispersion relation is deformed. In the models where a cubic term in the momentum is added ($n=1$) the best bound we obtained is 0.0066 $E_{pl}$. When a quartic term is added ($n=2$) the effect is much smaller and the corresponding best bound is $5.1 \\cdot 10^{-13} E_{pl}$.  The present work complements previous studies which found bounds of similar order of magnitude. Ellis et al. \\cite{Ellis:2002in,Ellis:2005wr} performed a statistical analysis in an ensemble of bursts with known redshifts (to correct for intrinsic time delays), finding the bounds $\\xi_1 > 5.6 \\cdot 10^{-4} $ and $\\xi_2 > 2.4 \\cdot 10^{-13}$ at a 95\\% of confidence level. In general, intrinsic delays render difficult to constrain Lorentz violations using a single burst, but we can still extract lower limits if there is simultaneity in the peak emissions. In our study, since no time delays are observed in the light curves with 2 msec bins, we can assume that, at that resolution, all photons were emitted simultaneously.  This is, of course, an assumption. We cannot rule out the possibility in which intrinsic spectral temporal shifts are exactly compensated by the time delays produced by the modified dispersion relation, Eq.  \\ref{mod_disp_rel}. While, in principle, this is possible, we believe it to be extremely unlikely.  Boggs et al. \\cite{boggs-2004} considered an unique extremely bright burst GRB 021206 observed by RHESSI. This burst has an almost flat spectrum in the band 1 - 17 MeV, which allowed estimates of the time delay up to 17 MeV. The redshift of GRB 021206 is not known as no host galaxy was detected. Assuming a redshift of $z \\simeq 0.3$ and using the simultaneity of the peaks at different energies, Boggs et al., \\cite{boggs-2004} obtained the bounds $\\xi_1 > 0.015 $ and $\\xi_2 > 4.5 \\cdot 10^{-12}$ for $n=1,2$.  These lower bounds are higher than ours by factors of 2.5 and 9 respectively.  However, the redshift of this burst was not measured, but only estimated from the spectral and temporal properties of the burst.  From this point of view, the present estimate, which is based on a burst with a known redshift, is more robust."
    },
    "0601/astro-ph0601283_arXiv.txt": {
        "abstract": "{This paper - the first of a short series dedicated to the long-standing astronomical problem of de-projecting the bi-dimensional apparent morphology of a three-dimensional mass of gas - focuses on the density distribution in real Planetary Nebulae (and all types of expanding nebulae). We introduce some basic theoretical notions, discuss the observational methodology and develope the accurate procedure for the  determination of the matter radial profile within the sharp portion of nebula in the plane of the sky identified by the zero-velocity-pixel-column (zvpc) of high-resolution spectral images. Moreover, a series of evolutive snapshots is presented, combining illustrative examples of model- and true-Planetary Nebulae. Last, the general and specific applications of the method (and some caveats) are discussed. ",
        "introduction": "The spatial distribution of gas constitutes the most important observational parameter for Planetary Nebulae (PNe): when combined with gas dynamics, it supplies fundamental information on the mechanisms and physical processes driving nebular evolution (mass-loss history, wind interaction, ionization, magnetic fields, binarity of the central star etc.) and allows us a direct comparison of each real object with theoretical evolutionary models, hydro-dynamical simulations and photo-ionization codes.  In spite of this, the whole astronomical literature dedicated to PNe contains only a handful of papers dealing with the true spatial structure, which is generally obtained by deconvolution of volume emissivity from the apparent image using Abel's integration equation for spherically symmetric systems and ``ad-hoc'' algorithms for inclined axi-symmetrical nebulae (Wilson \\& Aller 1951, Lucy 1974, Soker et al. 1992, Volk \\& Leahy 1993, Bremer 1995). These rare - quite rough and controversial - results  stress  the huge difficulties so far encountered in de-projecting the bi-dimensional appearance of a three-dimensional structure (Aller 1994).  Recently, Sabbadin et al. (2005 and references therein) have suggested that the dynamical properties of the ionized gas represent the key for overcoming the stumbling block of de-projection in PNe and other types of expanding nebulae, like Nova and Supernova Remnants, shells around Population I Wolf-Rayet stars, nebulae ejected by Symbiotic Stars, bubbles surrounding early spectral-type Main Sequence stars etc..  Historically, the general kinematical rule for the bright main shell of PNe goes back to Wilson (1950), based on Coud\\'e spectra - no image-derotator - of 26  targets: the high-excitation zones expand more slowly than the low-excitation ones, and there is a direct correlation between the expansion velocity and the size of the monochromatic image.  Weedman (1968) secured Coud\\'e spectra (+ image-derotator) of a sub-set of Wilson's list (10 PNe), covered along the apparent major axis. Assuming ``a priori'' the prolate spheroid hypothesis with a/b=1.5, Weedman derived a typical expansion law of the type $V_{\\rm exp}$=s(R-R$_0$), where s is the (positive) slope of the expansion velocity gradient and R$_0$ the radius at which $V_{\\rm exp}$=0. In most cases R$_0\\simeq$0 or R$_0$$<<$R$_{\\rm neb}$ (i.e. the outflow is nearly ballistic, Hubble-type).  The classical papers by Wilson (1950) and Weedman (1968) belong to the heroic age of photographic plates. Later on, the introduction of linear 2-D detectors at high efficiency (CCDs) greatly enhanced flux accuracy in a wide spectral range, and the use of echelle spectrographs allowed the observers to extend the analysis to large, faint objects. Even so, a odd dichotomy marks out the whole astronomical literature: high-resolution spectra of PNe are used to obtain either the detailed kinematics in a few ions (generally,  [O III], H I and [N II]) or the ``average'' nebular fluxes.  In the former case (i. e. kinematics), long-slit observations are performed through an interference filter to isolate a single order containing one or two nebular emissions. Valuable examples concern the kinematics of: 15 multiple shell PNe (Guerrero et al. 1998; single position angle (PA), instrumental spectral resolution $\\Delta$V=6.5 km s$^{-1}$), NGC 2438 (Corradi et al. 2000; single PA, $\\Delta$V=4.3 km s$^{-1}$) and MZ 3 (a symbiotic-star nebula), covered at various PA by Santander-Garcia et al. (2004, $\\Delta$V=6.0 km s$^{-1}$) and Guerrero et al. (2004, $\\Delta$V=8.0 km s$^{-1}$). We must notice, however, that the standard observational procedure - strongly reducing the useful spectral range by means of an interference filter - represents an excessive precaution in most cases, since the PN is an emission line object. As shown by Benetti et al. (2003) and Sabbadin et al. (2004), a PN larger (even much larger) than the order separation of the echellograms can be covered with a single exposure in the whole spectral range, the only limit (variable from one instrument to an other)  being the effective superposition of spectral images belonging to different orders  (and not the mere order separation).  In the latter dichotomous case (average fluxes from high-dispersion spectroscopy), line intensities are integrated over the entire slit (whose length is smaller than order separation) and provide ``mean'' physical conditions and ionic and chemical abundances within the whole slice of nebula covered by the spectrograph, thus losing the detailed (i. e. pixel-to-pixel) information on kinematics and flux. Representative examples are given by Hyung \\& Aller (1998) and Hyung et al. (2001).  A courageous effort for overcoming the foregoing impasse has been performed by Gesicki \\& Zijlstra (2000) and Gesicki et al. (2003 and references therein): they secured high-resolution spectra of a number of compact PNe (including Galactic Bulge objects and nebulae in the Sagittarius Galaxy and in the Magellanic Clouds) and combined integrated emission profiles of forbidden and recombination lines with the photo-ionization code of a spherical nebula (Torum model), inferring that acceleration is a quite common property in the PNe innermost layers (i. e. ``U''-shaped expansion profile), due to the dynamical contribution by the shocked, hot wind from the central star.  Unfortunately, a detailed comparative analysis (Sabbadin et al. 2005) proves that the ``U''-shaped expansion profile is a spurious, incorrect result caused by a combination of unsuited assumptions, spatial resolution and diagnostic choice.  All this confirms that, due to the nebular complexity and large stratification of the radiation and kinematics, the recovery of the spatio-kinematical structure needs observations:  (a) at adequate ``relative''  spatial (SS) and spectral (RR) resolutions (SS=R/$\\Delta$r, R=apparent radius, $\\Delta$r=seeing+guiding; RR=$V_{\\rm exp}$/$\\Delta$V, $\\Delta$V=instrumental spectral resolution),  (b) in a wide spectral range containing emissions at low, medium and high ionization and  (c) at several slit positions over the nebula,  to be combined with a straightforward and versatile method of analysis.  To this end, in 2000 we started a multi-PA spectroscopic survey of bright PNe in both hemispheres, observed  with ESO NTT+EMMI (80 echelle orders covering the spectral range $\\lambda\\lambda$3900-8000 $\\rm\\AA\\/$ with $\\Delta$V=5.0 km s$^{-1}$) and Telescopio Nazionale Galileo (TNG)+SARG (54 orders, spectral range $\\lambda\\lambda$4600-8000 $\\rm\\AA\\/$, $\\Delta$V=2.63 km s$^{-1}$). The pixel-to-pixel analysis of flux and velocity in a large number of emissions gives the bi-dimensional ionic structure (tomography) of each nebular slice intercepted by the spectrograph slit and a 3-D rendering program, assembling all tomographic maps, provides the accurate spatial distribution of the kinematics, physical conditions ($T_{\\rm e}$ and $N_{\\rm e}$) and ionic and chemical abundances within the nebula.  Besides the - in fieri - detailed spatio-kinematical study of individual PNe (Ragazzoni et al. 2001, Turatto et al. 2002, Benetti et al. 2003, Sabbadin et al. 2004, 2005), we decided to deepen a general, long-lasting open problem: inferring the radial density profile in expanding nebulae.  This introductory, didactic paper contains the detailed procedure for the determination of the gas distribution in real PNe and is structured as follows: Sect. 2 gives some basic theoretical notions, Sect. 3 describes the practical application to  high-dispersion spectra, Sect. 4 presents a series of evolutionary snapshots for model- and true-PNe, Sect. 5 contains the general discussion and Sect. 6 draws the conclusions.  In a future, applicatory paper we will compare the density profiles observed in a representative sample of targets with the expectations coming from recent hydro-dynamical simulations and theoretical models, in order to disentangle the evolutive phenomenology and physical processes responsible for PNe shape and shaping. ",
        "conclusions": "The presence of an expanding nebular gas is the typical signature of the instability phases in stellar evolution, characterized by a large, prolungated  mass-loss rate (PNe, Symbiotic Star nebulae, shells around Wolf-Rayet and LBV (luminous blue variable) stars), or even an explosive event (Nova and Supernova remnants).  Mass-loss of evolved stars occupies a strategic ground between stellar and interstellar physics: it raises fundamental astrophysical problems (e. g. origin and structure of winds, formation and evolution of dust, synthesis of complex molecules), plays a decisive role in the final stages of stellar evolution and is crucial for galactic enrichment in light and heavy elements.  So far, image de-projection represented an unsourmontable obstacle for the accurate recovery of the spatial structure in real nebulae, thus impeding any detailed comparative analysis with theoretical models. Such a frustrating situation is now overcome by tomography and 3-D recovery developed at the Astronomical Observatory of Padua (Italy).  The radial density reconstruction  illustrated in this paper represents the umpteenth proof that properly secured, reduced and analysed spectra at high-resolution constitute a peerless tool for deeping the spatio-kinematics, physical conditions, ionic structure  and evolution of all classes of expanding nebulae, thus bridging the present, wide  and pernicious gap between  model- and true-nebulae."
    },
    "0601/hep-ph0601047.txt": {
        "abstract": " ",
        "introduction": "The hierarchy problem exists in both particle physics and cosmology. In the former case, this is the hierarchy problem between gravity and standard particle physics mass scales, while in the latter case, it is the presence of a very small but finite cosmological constant or the dark energy indicated by the analysis of type-Ia supernovae observations \\cite{Perlmutter}, whose magnitude is as small as 120 orders smaller than its theoretically natural value, i.e., the Planck scale (or a similar scale).\\cite{Weinberg} \\ The origin of this small cosmological constant is not well understood yet. There has yet been no definite convection established in theoretical studies.  Recently, one of the present authors attempted to construct a model \\cite{Yoshimura} which simultaneously solves the above two hierarchy problems by radically changing cosmology in the same spirit as the ideas of Dirac \\cite{Dirac} and Brans and Dicke.\\cite{B-D,F-M} \\ In this model, a theory with two dilaton fields coupled to the scalar curvature considered.  The existence of many dilatons is conceivable in light of higher-dimensional theories. Moreover, it has been argued that an effective gravity mass scale (the square root of the inverse of Newton's constant) increases in the inflationary stage \\cite{Guth} (for a review of inflation, see Refs.~\\citen{Linde1} and \\citen{Kolb}), and that the small cosmological constant or the dark energy density in the present universe could be dynamically realized in the case that two dilatons have approximately an $O(2)$ symmetry, taking the fundamental mass scale to be of TeV. In such a TeV scale model, where the potential of the dilatons is of order $ \\sim (\\mathrm{TeV})^4$, however, it is conjectured that a naive estimate of the curvature perturbation gives a magnitude which is much smaller than the recent observational results of the anisotropy of the cosmic microwave background (CMB) radiation obtained from Wilkinson microwave background probe (WMAP). \\cite{Bennet} \\ Moreover, as is shown in \\S 3, power-law inflation could be realized in this model. In the standard inflation models, however, inflation is driven by the potential energy of a scalar field, called an ``inflaton\", as it slowly rolls down the potential hill. This slow roll over in the quasi-de Sitter stage is necessary to account for the nearly scale-invariant spectrum of the primordial curvature perturbation out of the quantum fluctuation of the inflaton, which is also suggested by the observational results obtained from WMAP.\\cite{Spergel}  However, in recent years a new mechanism for generating the primordial curvature perturbation has been proposed, in which a late-decaying massive scalar field provides the dominant source of the curvature perturbation.\\cite{Moroi,Enqvist,Lyth,Lyth2} \\ In this scenario, the dominant part of the curvature perturbation originates from the quantum fluctuations of a new scalar field, called the ``curvaton\", which is different from the inflaton. In contrast to the situation with the usual mechanism, with this new mechanism, inflation need not be of the slow-roll variety. Hence, even in the case of power-law inflation, the nearly scale-invariant spectrum of the curvature perturbation demanded by the observations can be realized. Instead, it is required that the curvaton potential be sufficiently flat during inflation. %, that is, the curvaton mass should be %much smaller than the Hubble parameter in the inflationary stage.  The purpose of the present paper is to argue that the curvaton scenario can be realized in the above theory, in which two dilatons are introduced along with the coupling to the scalar curvature.\\cite{Yoshimura} \\ In particular, we show that there exists a scalar field corresponding to the curvaton in the framework of this theory without introducing any other scalar field that plays the role of the curvaton, and discuss the case in which the curvaton scenario is realized. Also, we are able to construct the curvature perturbation with sufficiently large amplitude and nearly scale-invariant spectrum suggested by the observational results obtained from WMAP.  This paper is organized as follows. In \\S 2 we describe our model and derive field equations from its action. In \\S 3 we show that power-law inflation can be realized in this model. Next, in the framework of this theory the existence of a scalar field corresponding to the curvaton is demonstrated in \\S 4. Then, in \\S 5, we study the case in which the curvaton scenario can be realized in this model. Finally, \\S 6 is devoted to a conclusion. Throughout the present paper we use units in which $k_\\mathrm{B} = c = \\hbar = 1$ and denote Newton's constant by $G = {m_{\\mathrm{Pl}}}^{-2}$, where $m_{\\mathrm{Pl}} = 1.2 \\times 10^{19}$GeV is the Planck mass. ",
        "conclusions": "In the present paper, we have shown that the curvaton scenario is realized in a theory with two dilatons coupled to the scalar curvature \\cite{Yoshimura}. This theory has been considered in order to simultaneously solve the hierarchy problem between gravity and particle physics mass scales and the small cosmological constant or the dark energy problem. We have found that when both coupling constants %$\\epsilon_1$ and $\\epsilon_2$ between two dilatons and the scalar curvature %$\\epsilon_1$ and $\\epsilon_2$ are much smaller than unity, power-law inflation in which the power-law exponent is much larger than unity is realized in this model, and that when the difference between the values of %$\\epsilon_1$ and $\\epsilon_2$ the two coupling constants is much smaller than these values themselves [in other words, the two dilatons have an approximate $O(2)$ symmetric coupling], there exists a scalar field corresponding to the curvaton whose mass is much smaller than the Hubble parameter in the inflationary stage in this theory. This is realized without introducing any other scalar field that may play the role of the curvaton. Furthermore, we have investigated the simple version of the curvaton scenario in which the curvaton has a quadratic potential, and we demonstrated that a curvature perturbation with a sufficiently large amplitude and a nearly scale-invariant spectrum suggested by observations obtained from WMAP can be generated.  \\if In the present paper we have discussed that the curvaton scenario could hold true in a theory with two dilatons coupled to the scalar curvature \\cite{Yoshimura}. As a result, we have found that in the case the two dilatons have approximately O(2) symmetry there could be a scalar field corresponding to the curvaton in the framework of this theory without introducing any other scalar fields playing a role of the curvaton and thus the curvaton scenario could be realized. Furthermore, we have discussed the simple version of the curvaton scenario in which the curvaton has the quadratic potential. As a result, the curvature perturbation with the enough amplitude and nearly scale-invariant spectrum suggested by the observational results obtained from WMAP could be generated. \\fi   %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% \\if The serious problem of this model is that it is exceedingly difficult for the present model to be compatible with the scenario proposed in Ref.~\\citen{Yoshimura}, %consistent with the scenario proposed in Ref.~\\citen{Yoshimura}. %it is exceedingly difficult to be successful at both the present model %and the scenario proposed in Ref.~\\citen{Yoshimura}, which simultaneously solves the above two hierarchy problems. This reason is as follows:  In the present model, \\fi  Finally, we note that in the present model, $\\varphi_{\\mathrm{r}} = \\sqrt{\\varphi_1^2 + \\varphi_2^2} \\approx \\gamma (t) M$ exists around the origin of the potential $V$ %$ %V[\\varphi_i] =V_{0} \\cos \\left( \\varphi_{\\mathrm{r}}/M \\right) + \\Lambda %$ in Eq.\\ (\\ref{eq:3}) in the inflationary stage; i.e., the case $\\gamma (t) \\ll 1$ during inflation is considered. After sufficient inflationary expansion, the amplitude of $\\varphi_{\\mathrm{r}}$ becomes large, and then $\\varphi_{\\mathrm{r}}$ approaches the first minimum of the potential. %i.e., $\\varphi_{\\mathrm{r}} = M\\pi$. After the oscillation epoch, $\\varphi_{\\mathrm{r}}$ remains around the first potential minimum. Finally, the dilatons decay into radiation through coupling to other lighter fields, and then the universe is reheated. By contrast, in the scenario proposed in Ref.~\\citen{Yoshimura}, the following case is considered: In the inflationary stage, $\\varphi_{\\mathrm{r}}$ goes over many local maxima of the potential and then changes for many periods; i.e., there exist the relations $\\gamma (t) \\gtrsim 1$ and $\\gamma (t_{\\mathrm{f}}) - \\gamma (t_{\\mathrm{i}}) \\gg 1$, where $t_{\\mathrm{f}}$ and $t_{\\mathrm{i}}$ are the time at the end and beginning of inflation, respectively. Thus, the situation for the evolution of $\\varphi_{\\mathrm{r}}$ during inflation considered in the present model is different from that in the scenario proposed in Ref.~\\citen{Yoshimura}. \\if This makes it difficult that both the present model and the scenario proposed in Ref.~\\citen{Yoshimura} stand together. \\fi %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%            \\if the situation of the evolution of $\\varphi_{\\mathrm{r}}$ during inflation in the present model is different from that in the scenario in Ref.~\\citen{Yoshimura}. The present model \\fi"
    },
    "0601/astro-ph0601626_arXiv.txt": {
        "abstract": "We estimate minimum energy magnetic fields ($B_{\\rm min}$) for a sample of galaxies with measured gas surface densities. The sample spans more than four orders of magnitude in surface density from normal spirals to luminous starbursts.  We show that the ratio of the minimum energy magnetic pressure to the total pressure in the ISM decreases substantially with increasing surface density.  For the ultra-luminous infrared galaxy Arp 220, this ratio is $\\sim10^{-4}$. Therefore, if the minimum energy estimate is applicable, magnetic fields in starbursts are dynamically weak compared to gravity, in contrast to our Galaxy and other normal star-forming spiral galaxies. We argue, however, that rapid cooling of relativistic electrons in starbursts invalidates the minimum energy estimate.  We critically assess a number of independent constraints on the magnetic field strength in starburst galaxies.  In particular, we argue that the existence of the FIR-radio correlation implies that the synchrotron cooling timescale for cosmic ray electrons is much shorter than their escape time from the galactic disk; this in turn implies that the true magnetic field in starbursts is significantly larger than $B_{\\rm min}$.  The strongest argument {\\it against} such large fields is that one might expect starbursts to have steep radio spectra indicative of strong synchrotron cooling, which is not observed. We show, however, that ionization and bremsstrahlung losses can flatten the nonthermal spectra of starburst galaxies even in the presence of rapid cooling, providing much better agreement with observed spectra.  We further demonstrate that ionization and bremsstrahlung losses are likely to be important in shaping the radio spectra of most starbursts at GHz frequencies, thereby preserving the linearity of the FIR-radio correlation.  We thus conclude that magnetic fields in starbursts are significantly larger than $B_{\\rm min}$.  We highlight several observations that can test this conclusion. ",
        "introduction": "\\label{section:intro}  The magnetic energy density of the Galaxy is observed to be in rough equipartition with the cosmic ray energy density and the turbulent pressure. The sum of these yields a total midplane pressure consistent with that required by hydrostatic equilibrium, given the mass distribution in the solar neighborhood (Boulares \\& Cox 1990). This equipartition magnetic field is roughly 6 $\\mu$G at the solar circle (e.g., Beck 2001). Similar magnetic field strengths are found in most normal star-forming spiral galaxies (Fitt \\& Alexander 1993; Niklas 1995).  Starbursts and ultra-luminous infrared galaxies (ULIRGs) have much higher surface densities and turbulent velocities than local spirals. As an example, Arp 220 has a gas surface density $\\sim 10^4$ times higher than the Galaxy (Downes \\& Solomon 1998).  If Arp 220 has a magnetic energy density in rough equipartition with its hydrostatic pressure, the implied magnetic field strength would be $\\sim0.03$\\,G on few-hundred parsec scales (Fig.~\\ref{fig:bmin}, Table \\ref{table:starburst}).  This is substantially larger than the value of $\\sim 1$ mG typically inferred in ULIRGs (e.g., Condon et al.~1991) using the observed radio emission and the classic ``minimum energy'' argument (Burbidge 1956; Longair 1994).  Throughout this paper we use the term ``minimum energy'' to refer to the magnetic field strength inferred using the observed radio flux and assuming comparable cosmic ray and magnetic energy densities ($B_{\\rm min}$; eq.~[\\ref{bmin}]).  We reserve the term ``equipartition'' for magnetic field strengths that are dynamically important with respect to gravity ($B_{\\rm eq}$; eq.~[\\ref{bequipartition}]).  We find that the large discrepancy between the minimum energy and equipartition magnetic field estimates in Arp 220 is generic to starburst galaxies; this discrepancy motivates the analysis of magnetic field strengths in luminous starbursts presented in this paper. In \\S\\ref{section:results} we present a sample of galaxies with measured gas surface densities and radio fluxes, for which both minimum energy and equipartition field strengths can be estimated.  In \\S\\ref{section:against} we summarize arguments in favor of $B \\sim B_{\\rm min}$ and against significantly larger fields.  Rebuttals to a subset of these arguments are given in \\S\\ref{section:for}, where we also provide independent arguments why magnetic fields in starbursts are likely to be $\\gg B_{\\rm min}$. These arguments draw heavily on the FIR-radio correlation for star-forming galaxies and the observed radio spectra of starbursts. Finally, in \\S\\ref{section:discussion} we summarize our results and discuss their implications.  \\  \\ \\begin{figure*} \\centerline{\\hbox{\\psfig{file=plot_bminp.ps,width=12.5cm}}} \\figcaption[x]{Minimum energy magnetic field ($B_{\\rm min}$; eq.~[\\ref{bmin}]) versus measured gas surface density ($\\Sigma_g$). Normal star-forming galaxies (open squares; Table \\ref{table:spiral}), starburst galaxies (filled squares; Table \\ref{table:starburst}), and the Galaxy ($B_{\\rm min}=6$ $\\mu$G at the solar circle, e.g., Beck 2001; $\\Sigma_g\\simeq2.5\\times10^{-3}$ g cm$^{-2}$, Boulares \\& Cox 1990) are shown.  To indicate the uncertainties at high $\\Sigma_g$, the open triangles show $B_{\\rm min}$ for IC 883 (Arp 193), Mrk 273, and the individual nuclei of Arp 220, as inferred using $\\Sigma_g$ and the radial size from Downes \\& Solomon (1998) (their ``extreme'' starbursts; see their Table 12) and with the radio flux from Condon et al.~(1991) at 8.44 GHz (see Table \\ref{table:extreme}).  The dotted line is the equipartition magnetic field (eq.~[\\ref{bequipartition}]) and the dashed line is the scaling $B_{\\rm min} \\propto \\Sigma_g^{2/5}$ (eq.~[\\ref{bminfir}]).  The latter scaling is expected if the cosmic ray electron cooling timescale is much shorter than the escape timescale from the galactic disk, in which case the true magnetic field is also significantly larger than the minimum energy estimate (see \\S\\ref{section:for}; eq.~[\\ref{bminfir}]).  The solid lines show the scalings $B \\propto \\Sigma_g^a$ (with a = 0.9, 0.8, and 0.7) used in \\S \\ref{section:for_breaks}. \\\\ \\label{fig:bmin}} \\end{figure*} ",
        "conclusions": "\\label{section:discussion}  Based on the analysis of the preceding sections, it is worthwhile contrasting the following two scenarios for the radio emission from starbursts:  \\noindent {\\it $B \\sim B_{\\rm min}$:} For the minimum energy magnetic field estimate to be applicable, the electrons must escape rapidly in a galactic wind with $\\tau_{\\rm cool} \\gtrsim \\tau_{\\rm esc}$.  The rapid escape implies that there is no cooling break at GHz frequencies, which is consistent with the observed nonthermal spectra of $\\alpha \\approx 0.75$.  Rapid escape is plausible given that hot outflowing galactic winds are observed from starbursts and that the cosmic rays are likely generated in the same supernova shocks that generate the wind.  This interpretation of magnetic fields in starbursts --- the standard one --- has two major drawbacks: (1) it is difficult to see how $\\tau_{\\rm cool} \\gtrsim \\tau_{\\rm esc}$ is applicable in ULIRGs where cooling times from IC alone are $\\sim 10^4$ yr and (2) it is very difficult to account for the FIR-radio correlation with $\\tau_{\\rm cool} \\gtrsim \\tau_{\\rm esc}$ because variations in $\\tau_{\\rm cool}$ or $\\tau_{\\rm esc}$ should lead to variations in the fraction of the electron energy radiated away via synchrotron radiation.  Either several coincidences or a complex feedback loop are then required to account for the linearity of the FIR-radio correlation.  The latter would be somewhat surprising given that $B_{\\rm min}$ itself implies that magnetic fields and cosmic rays are dynamically weak compared to gravity (Fig.~\\ref{fig:bmin}).  \\noindent {\\it $B \\gg B_{\\rm min}$:} An alternative possibility is that magnetic fields in starbursts are much stronger than is suggested by the minimum energy estimate.  Rapid cooling of relativistic electrons in starbursts will invalidate the minimum energy estimate if $\\tau_{\\rm cool} \\ll \\tau_{\\rm esc}$.  Magnetic fields in starbursts could then in principle be as large as $\\sim B_{\\rm eq}$, which is $\\sim 10$ times larger than $B_{\\rm min}$ for typical starbursts (Fig.~\\ref{fig:bmin}). This model naturally accounts for the linearity of the FIR-radio correlation because in the limit $\\tau_{\\rm cool} \\lesssim \\tau_{\\rm esc}$ the radio flux is nearly independent of the magnetic field strength (e.g., V\\\"olk 1989). In addition, the observed radio luminosities from star-forming galaxies are comparable to the estimated rate at which energy is supplied to relativistic electrons in supernova shocks (\\S\\ref{section:for_fir}).  This supports the hypothesis that $\\tau_{\\rm cool} \\lesssim \\tau_{\\rm esc}$.  Independent support for this hypothesis is provided by the fact that the observed scaling $B_{\\rm min} \\propto \\Sigma_g^{2/5}$ for the minimum energy field in starbursts (Fig. \\ref{fig:bmin}) follows directly from the Schmidt Law for star formation when $\\tau_{\\rm cool}\\lesssim\\tau_{\\rm esc}$ (eq. [\\ref{bminscale}]).  A standard objection to $\\tau_{\\rm cool} \\lesssim \\tau_{\\rm esc}$ is that the typical nonthermal spectral indices of starbursts ($\\alpha \\approx 0.75$) do not show the expected steepening due to strong cooling ($\\alpha \\approx 1$).  We have argued, however, that if $B \\gg B_{\\rm min}$ and if cosmic ray electrons interact with gas at approximately the mean density of the ISM, then ionization and bremsstrahlung losses flatten the radio spectra of starbursts at $\\sim {\\rm GHz}$ frequencies and reconcile $\\tau_{\\rm cool} \\lesssim \\tau_{\\rm esc}$ with the observed spectra (Fig.~\\ref{fig:slope}).  An important part of this argument is the realization that the ionization and bremsstrahlung loss times are similar to the synchrotron cooling time in all starburst galaxies for cosmic ray electrons emitting at GHz frequencies.  Thus, ionization and bremsstrahlung losses can modify the nonthermal spectra of starburst galaxies and yet maintain the linearity of the FIR-radio correlation (Fig. \\ref{fig:firradio}).  This interpretation of magnetic fields in starbursts --- that $B \\gg B_{\\rm min}$ --- has two potential drawbacks: (1) Given that galactic winds efficiently remove mass and metals from galaxies, it is unclear whether the cosmic rays actually interact with the bulk of the ISM in starbursts, which is required for ionization losses to be significant. (2) The nonthermal synchrotron halos observed in several systems are typically interpreted as cosmic-ray electrons advected out with a galactic wind, with synchrotron cooling only becoming important in the halo (e.g., Seaquist \\& Odegard 1991).  If this interpretation is correct, it requires $B \\sim B_{\\rm min}$ (\\S\\ref{section:against_halos}).  We have suggested, however, that extended synchrotron emission may be due to particles accelerated {\\it in situ} in galactic winds (\\S\\ref{section:for_halos}), in which case the observed halos cannot be readily used to constrain the magnetic field strength in galactic disks.  We believe that the arguments presented in this paper strongly favor $\\tau_{\\rm cool} \\lesssim \\tau_{\\rm esc}$ for the relativistic electrons in starbursts.  In this limit, the minimum energy argument underestimates the true magnetic field strength.  The observed radio flux is, however, nearly independent of the magnetic field and depends primarily on the rate at which energy is supplied to relativistic electrons.  Thus direct constraints on the true field strength are difficult to come by.  By themselves, these arguments do not imply that $B\\sim B_{\\rm eq}$, only that $\\tau_{\\rm cool} \\lesssim \\tau_{\\rm esc}$ and $B \\gtrsim B_{\\rm min}$.  Our strongest argument that magnetic fields are likely to be $\\gg B_{\\rm min}$ derives from the ratio of the ionization and bremsstrahlung loss times to the synchrotron cooling time: if cosmic ray electrons interact with gas having roughly the mean density of the ISM, magnetic fields {\\it must} be $\\gg B_{\\rm min}$ in starbursts in order for ionization and bremsstrahlung losses to not completely dominate over synchrotron losses (which would be inconsistent with the FIR-radio correlation and would require that an unreasonably large fraction of the supernova energy be supplied to cosmic ray electrons in order to account for the observed radio fluxes from starbursts; \\S\\ref{section:for_breaks}).  Our conjecture is thus that the magnetic field is significantly larger than $\\sim B_{\\rm min}$ (the canonical estimate) in many starbursts. Confirming this prediction would have significant implications for understanding the physics of star formation in starbursts, the effects of magnetic fields on the dynamics of galactic winds, and the role of magnetic stresses in driving gas to smaller radii to fuel a central active galactic nucleus.  Towards this end, we briefly summarize several observations that should help test the predictions of our model: (1) If ionization and bremsstrahlung losses are important, the radio spectra of starbursts should be flatter below $\\sim 1$ GHz than they are above $\\sim 1$ GHz (\\S\\ref{section:for_breaks} and Fig. \\ref{fig:slope}).  (2) If cosmic rays interact with gas at about the mean density of the ISM, as is required for ionization and bremsstrahlung losses to be important, there should be an appreciable gamma-ray flux from starbursts due to neutral pion production (\\S\\ref{section:for_gamma}).  We estimate $\\nu L_\\nu \\approx 2 \\times 10^{-5} L_{IR}$ above $\\approx 100$ MeV.  This prediction is testable with {\\it GLAST}. (3) Zeeman measurements of starbursts can directly probe the magnetic field strength in the dense phase of the ISM, and distinguish between $B \\sim B_{\\rm min}$ and $B \\gg B_{\\rm min}$. Existing upper limits in 4 ULIRGs from Killeen al.~(1996) are above $B_{\\rm min}$, but below the equipartition field (\\S\\ref{section:against_zeeman}).  (4) If $B \\gg B_{\\rm min}$, IC emission is unlikely to contribute significantly to the X-ray emission in starbursts (\\S\\ref{section:for_ic}).  It is interesting to compare our results on magnetic fields in starbursts with the inferred correlation between magnetic field strength, column density, and mass density in Galactic molecular clouds (e.g., Troland \\& Heiles 1986; Crutcher 1999).  Using Zeeman splitting to determine the magnetic field strength directly, Crutcher (1999) finds that over the range $0.01 \\, {\\rm g \\, cm^{-2}} \\lesssim \\Sigma_{\\rm H_2} \\lesssim 2 \\, {\\rm g \\, cm^{-2}} \\lesssim$ and $2.7\\lesssim \\log_{10}[n_{\\rm H_2}\\,({\\rm cm}^{-3})]\\lesssim6.8$, that $B\\propto n_{\\rm H_2}^{1/2}$ and $B\\propto \\Sigma_{\\rm H_2}$. These findings imply both that that the Alfv\\'en speed $v_A \\sim {\\rm constant} \\sim 10$ km s$^{-1}$ and that equipartition obtains. Observations of polarization in H$_2$O masers in dense star-forming regions provide evidence that these trends extend to yet higher densities ($10^8-10^{10}$ cm$^{-3}$; Sarma et al.~2002; Vlemmings et al.~2005).  These results suggest that the magnetic field is always comparable to the total pressure in Galactic molecular clouds, with $B\\approx B_{\\rm eq}$.  The sample of Galactic molecular clouds reviewed by Crutcher (1999) covers precisely the same surface densities represented by starbursts in our Figure \\ref{fig:bmin}.  It is thus particularly striking that the minimum energy estimate implies $B \\ll B_{\\rm eq}$ in starbursts, while direct Zeeman measurements of local molecular clouds with the same $\\Sigma_g$ imply $B \\sim B_{\\rm eq}$.  Our suggestion is that this difference arises in part from the inapplicability of the minimum energy estimate in starbursts."
    },
    "0601/hep-th0601205_arXiv.txt": {
        "abstract": "We investigate time dependent solutions (S-brane solutions) for product manifolds consisting of factor spaces where only one of them is non-Ricci-flat. Our model contains minimally coupled free scalar field as a matter source. We discuss a possibility of generating late time acceleration of the Universe. The analysis is performed in conformally related Brans-Dicke and Einstein frames. Dynamical behavior of our Universe is described by its scale factor. Since the scale factors of our Universe are described by different variables in both frames, they can have different dynamics. Indeed, we show that with our S-brane ansatz in the Brans-Dicke frame the stages of accelerating expansion exist for all types of the external space (flat, spherical and hyperbolic). However, applying the same ansatz for the metric in the Einstein frame, we find that a model with flat external space and hyperbolic compactification of the internal space is the only one with the stage of the accelerating expansion. Scalar field can prevent this acceleration. It is shown that the case of hyperbolic external space in Brans-Dicke frame is the only model which can satisfy experimental bounds for the fine structure constant variations. We obtain a class of models where a pare of dynamical internal spaces have fixed total volume. It results in fixed fine structure constant. However, these models are unstable and external space is non-accelerating. ",
        "introduction": "Introduction}  \\setcounter{equation}{0}  Recent astronomical observations abundantly evidence that our Universe underwent stages of accelerating expansion during its evolution. There are at least two of such stages: early inflation and late time acceleration. The latter began approximately at the redshift $z \\sim 1 $ and continues until now. Thus, the construction and investigation of models with stages of acceleration is one of the main challenge of the modern cosmology.   Among such models, the models originated from fundamental theories (e.g. string/M-theory) are of the most of interest. For example, it was shown that some of spacelike brane (S-brane) solutions have a stage of the accelerating expansion. We remind that in $D$-dimensional manifold S$p$-branes are time dependent solutions with ($p+1$)-dimensional Euclidean world-volume and apart from time they have $(D-p-2)$-dimensional hyperbolic, flat or spherical spaces as transverse/additional dimensions \\cite{GS}: \\ba{0.1} \\nonumber ds_D^2 = -e^{2\\gamma(\\tau )} d\\tau^2 &+& a_0^2(\\tau )(dx_1^2 + \\ldots + dx_{p+1}^2)\\\\ &+ &a^2_1(\\tau ) d\\Sigma^2_{(D-p-2),\\, \\sigma}\\, \\, , \\ea where $\\gamma (\\tau )$ fixes the gauge of time, $a_0 (\\tau )$ and $a_1 (\\tau )$ are time dependent scale factors and $\\sigma = -1,0,+1 $ for hyperbolic, flat or spherical spaces respectively\\footnote{Slightly generalized ansatz where the $(D-p-2)$-dimensional transverse space consists of the $k$-dimensional hyperspace $\\Sigma_{k,\\sigma}$ and $(q-k)$-dimensional Euclidean space was considered in \\cite{CGG}. Here, $D-p-2 = k+q$.}. Obviously, $p=2$ if brane describes our 3-dimensional space. These branes usually known as SM2-branes if original theory is 11-dimensional M-theory and SD2-branes in the case of 10-dimensional Dirichlet strings. For this choice of $p$, the evolution of our Universe is described by the scale factor $a_0$. In general, the scale factor $a_1$ can also determine the behavior of our 3-dimensional Universe. Hence, $D-p-2=3$ and we arrive to SM6-brane in the case of the M-theory and SD5-brane for the Dirichlet string. Usually, S$p$-brane models include form fields (fluxes) and massless scalar fields (dilatons) as a matter sources. If SD$p$-branes are obtained by dimensional reduction of 11-dimensional M-theory, then the dilaton is associated with the scale factor of a compactified 11-th dimension.  Starting from \\cite{GS}, the S-brane solutions were also found, e.g., in Refs. \\cite{CGG,Iv2,Oh1,Burgess}. It was quite naturel to test these models for the accelerating expansion of our Universe. Really, it was shown in \\cite{Oh2} that the SM2-brane as well as the SD2-brane have stages of the accelerating behavior. This result generalizes conclusions of \\cite{TW} for models with hyperbolic compact internal spaces. Here, the cosmic acceleration (in Einstein frame) is possible due to a negative curvature of the internal space that gives a positive contribution into an effective potential. This acceleration is not eternal but has a short period and the mechanism of such short acceleration was explained in \\cite{EG}. It was indicated in \\cite{Oh2} that the solution of \\cite{TW} is the vacuum case (the zero flux limit) of the S-branes. It was natural to suppose that if the acceleration takes place in the vacuum case, it may also happen in the presence of fluxes. Indeed, it was confirmed for the case of the compact hyperbolic internal space. Even more, it was found that periods of the acceleration occur in the cases of flat and spherical internal spaces due to the positive contributions of fluxes into the effective potential.  It is worth of noting that accelerating multidimensional cosmological models are widely investigated for last few years for different types of models. In general, such models can be divided into two main classes\\footnote{Apart of these models, interesting accelerating cosmologies following from non-linear models were proposed in \\cite{NO}}. First class consists of models where the internal spaces are stabilized and the acceleration is achieved due to a positive minimum of an effective potential which plays the role of a positive cosmological constant. General method for stabilization of the internal spaces was proposed in \\cite{GZ1} and numerous references can be found, e.g., in Refs. \\cite{GSZ,GMZ,Dark}. Models where both external (our) space and internal spaces undergo dynamical behavior constitute the second class of models. These models were considered, e.g., in \\cite{Gu,CV,Mohammedi} where a perfect fluid plays the role of a matter source and the cosmic acceleration happens in Brans-Dicke frame. Obviously, the S-branes accelerating solutions belong to the second class of models. Along with mentioned above Ref. \\cite{Oh2}, the accelerating S-brane cosmologies (in Einstein frame) were obtained and investigated, e.g., in Refs. \\cite{Roy,W2,GKL,CHNOW,Neupane}. Closely related to them accelerating solutions were also found in Refs. \\cite{CHNW,Iv3} (see also general discussion on inflationary cosmologies with the sum of exponential effective potentials in \\cite{W3}; the complete classification of solutions for such models according to their late-time behavior is given in \\cite{Jarv}). It should be noted that some of these solutions are not new ones but either rediscovered or written in different parametrization (see corresponding comments in Refs. \\cite{Iv2,Iv3}). For example, the first vacuum solution for a product manifold (consisting of $(n-1)$ Ricci-flat spaces and one Einstein space with non-zero constant curvature) was found in \\cite{Iv1}\\footnote{The first quantum solutions as well as the euclidian classical solution for this model in the presence of a massless minimally coupled scalar field were obtained in \\cite{Zhuk1}.}. This solution was generalized to the case of a massless scalar field in Refs. \\cite{BZpositive,BZnegative}. Obviously, solutions in Refs. \\cite{Iv1,BZpositive,BZnegative} are the zero flux limit of the Sp-branes and the result of \\cite{TW} is a particular case of \\cite{BZnegative}\\footnote{However, the period of accelerating expansion was not singled out in \\cite{BZnegative}. }. Some of solutions in \\cite{CHNOW,CHNW} coincides with corresponding solutions in Refs. \\cite{Iv1,BZpositive,BZnegative,BZ3}. An elegant minisuperspace approach for the investigation of the product space manifolds consisting of Einstein spaces was proposed in \\cite{IMZ}. Here, it was shown that the equations of motion have the most simple form in a harmonic time gauge\\footnote{For Eq. \\rf{0.1}, it  reads $\\gamma = (p+1)\\ln a_0 + (D-p-2)\\ln a_1 $. In the harmonic time gauge, time satisfies equation $ \\Delta [g] \\tau = 0$ \\cite{IMZ}.} because the minisuperspace metric is flat in this gauge. Even if the authors of the mentioned above papers did not aware of it, they intuitively used this gauge to get exact solutions. New solutions can be also generated (from the known solutions) with the help of a topological splitting when Einstein space with non-zero curvature is splitted into a number of Einstein spaces of the same sign of the curvature (see Refs. \\cite{GIM,split}). This kind of solutions was found, e.g., in Refs. \\cite{CHNOW,CHNW}.   Our paper is devoted to a model with the product of $n$ Einstein spaces where all of them are Ricci-flat but one with positive or negative curvature. We include massless scalar field as a matter source. As we mentioned above, the general solutions for this model was found in our papers \\cite{BZpositive,BZnegative}. Here, all factor spaces are time dependent. Obviously, these solutions are the zero flux limit of the S$p$-branes. The aim of the present investigations is twofold.  First, we give the detail analysis for the accelerating behavior of the external (our) space. At this stage, both the Ricci-flat space and non-zero curvature space may play the role of our Universe. The investigation is conducting in Einstein as well as Brans-Dicke frames. The transition between these two frames is performed with the help of the conformal transformation of the metric of the external spacetime. Such transformation does not destroy neither factorizable structure of $D-$dimensional metric ansatz nor the topology of factor spaces. However, scale factors of our Universe are described by different variables in the Brans-Dicke and Einstein frames. These variables are connected with each other via conformal transformation (see Appendix). Moreover, synchronous times are also different in both frames. Obviously, these different scale factors may behave differently with corresponding synchronous times. Precisely this interpretation we bear in mind when we write about different behavior of our Universe in different frames. For example we show that in Brans-Dicke frame, stages of the accelerating expansion exist for all types of the external space (flat, spherical and hyperbolic). However, in Einstein frame, the model with flat external space and hyperbolic compactification of the internal space is the only one with the stage of the accelerating expansion, in agreement with the results Refs. \\cite{TW,CHNOW}. A new result here is that scalar field can prevent the acceleration in the Einstein frame.  Second, we investigate the variation of the fine structure constant in our model. It is well known that dynamical internal spaces result in the variations of the fundamental constants (see, e.g., Refs. \\cite{GSZ,CV} and references therein). For example, the fine structure constant is inversely proportional to the volume of the internal space. However, there are strong experimental restrictions for the variations of the fundamental constants (see, e.g., \\cite{Uzan}). Thus, any multidimensional cosmological models with time dependent internal spaces should be tested from this point of view. In our paper, we show that considered model have a significant problem to satisfy these limitations for the variation of the fine structure constant. The case of the hyperbolic external space in the Brans-Dicke frame is the only possibility to avoid this problem, if there is no other way to explain the constancy of the effective four-dimensional fundamental constants in multidimensional models. For example, we propose models with the hyperbolic or spherical external space and two Ricci-flat internal spaces where the total volume of the internal spaces is the constant. Here, the dynamical factors of the internal spaces mutually cancel each other  in the total volume element. Thus, the effective fundamental constants remain really constant in spite of the dynamical behavior of the internal spaces. However, this model is unstable and the external space is non-accelerating. Anyway, such models are of special interest because indicate a possible way to avoid the fundamental constant variations in higher-dimensional theories.  The paper is structured as follows. In section \\ref{sec:2}, we explain the general setup of our model and present the exact solutions for a product manifold consisting of two factor spaces where only one of them is non-Ricci-flat. These solutions is carefully investigated in section \\ref{sec:3} (spherical factor space) and \\ref{sec:4} (hyperbolic factor space) for the purpose of the accelerating behavior of the external space. In section \\ref{sec:5}, we compare the rate of variations of the fine structure constant in our accelerating models with the experimental bounds. In section \\ref{sec:6}, we obtain and discuss a solution with three factor spaces where two dynamical internal spaces have the fixed total volume. The main results are summarized in the concluding section \\ref{sec:7}. ",
        "conclusions": ""
    },
    "0601/astro-ph0601318_arXiv.txt": {
        "abstract": "We analyzed the recently published   kHz QPO data in the neutron star low-mass X-ray binaries (LMXBs), in order to investigate the different correlations of the twin peak kilohertz quasi-periodic oscillations (kHz QPOs) in bright Z sources and in the less luminous Atoll sources. We find that a power-law relation $\\no\\sim\\nt^{b}$ between the upper  and the lower kHz QPOs with different indices: $b\\simeq$1.5 for the Atoll source 4U 1728-34 and $b\\simeq$1.9 for the Z source Sco X-1. The implications of our results for the theoretical models for kHz QPOs are discussed. ",
        "introduction": "With the advent of the Rossi X-ray Timing Explorer ({\\em RXTE}), our knowledge of the properties of the aperiodic variability of neutron star (NS) low mass  X-ray binaries (LMXBs) took a substantial step forward, especially initiated by the discovery of the  kilohertz quasi-periodic oscillations (kHz QPOs) in about twenty more LMXBs (see van der Klis 2000, 2004, for a review). The kHz QPOs in the power spectra of these systems cover the range of frequency from some hundred Hz to more than one  kHz, and they often occur in pairs in the persistent emission: the upper kHz QPO frequency ($\\nt$, hereafter the upper-frequency) and the lower kHz QPO frequency ($\\no$, hereafter the lower-frequency). The kHz QPOs  were soon found to behave in a rather regular way and the study of their phenomenology led to the discovery of tight correlations between. their frequencies and other observed characteristic frequencies (see, e.g., Psaltis et al. 1998, 1999ab;  Stella et al.  1999; Belloni \\et \\,2002). Without any doubt, {\\em RXTE} has provided a probe into the accretion flow in the non-Newtonian strong gravity regime where Einstein's  General Relativity might be tested (van der Klis 2000, 2004).  The correlation between the upper-frequency and the lower-frequency across different sources or for a particular  source, such as Sco X-1, can be roughly fitted  by a  power law function (see, e.g., Psaltis et al. 1998, 1999a), but also by a linear model (see Belloni et al. 2005).  The kHz QPO peak separation $\\dn\\equiv\\nt-\\no$ between the upper-frequency and lower-frequency   in a given source is generally inconsistent with a NS  spin frequency. In some sources, $\\dn$ is lower or higher than the NS spin frequency (when directly measured, see e.g. M\\'endez \\& van der Klis 1999; Jonker, M\\'endez, \\& van der Klis 2002b) or than the nearly coherent oscillation frequency ($\\nb$) observed during type~I X-ray burst that is identified to be the stellar spin frequency (Muno 2004; Strohmayer \\& Bildsten 2003; Wijnands et al 2003;  van der Klis 2004). The averaged peak separations  are found to be either close to (but not consistent with) the spin frequency $\\ns$ or to its half $\\ns$/2 (see, e.g., Wijnands \\et 2003; van der Klis 2004; Lamb \\& Miller 2001). The above observations offer strong evidence   against the simple beat-frequency model, in which the lower-frequency is the beat  between the upper-frequency and the NS spin frequency $\\ns$ (see e.g. Strohmayer et al.\\ 1996; Zhang et al. 1997; Miller et al. 1998), i.e. $\\no=\\nt-\\ns$. Nonetheless, with the discovery of  pairs of 30--450 Hz QPOs  from a few black-hole candidates with  frequencies ratios 3:2 (see, e.g., van der Klis 2004),   Abramowicz \\et  (2003ab) reported that the ratios of twin  kHz QPOs in  Sco X-1 tend to cumulate  around a value of 3:2, and they interpreted  it  as  an evidence of  a near $\\sim$  3:2 resonance. This was further argued by  Abramowicz \\et (2003ab) to be a promising link with the black-hole high-frequency QPOs (see e.g. van der Klis 2004). Moreover, the production mechanisms of kHz QPOs are  still  open issues: they have been identified with various characteristic frequencies in the inner accretion flow (see e.g. Stella \\& Vietri 1999;  Titarchuk et al. 1998;  Titarchuk \\& Osherovich 2000; Psaltis \\& Norman 2000; Lamb \\& Miller 2001; Zhang 2004).  In this paper,  in order to check the  predictions of the kHz QPO models, we analyze  the recently  published   kHz QPO data  by {\\em RXTE\\/}, which have been used by several authors (see, e.g., Belloni et al. 2005;  M\\'endez \\& van der Klis  1999, 2000; Psaltis et al. 1998, 1999ab;  van der Klis 2000, 2004, and original references therein). Most of the  data are   provided by T. Belloni, M. M\\'endez and D. Psaltis, and the others  are   extracted  from  the references listed in Table 1. Therefore, the  data we analyzed  here constitute a larger sample than  that presented  by  Belloni  \\et  (2005). In section 2, we  critically discuss the twin  kHz QPO correlation. The conclusions and discussions are given  in the last section. ",
        "conclusions": "We have analyzed an updated  sample of frequencies of the simultaneously detected twin kHz QPOs in LMXBs. Our main conclusions  are the following. (1) The power-law correlations were analyzed by means  of $\\chi^2$ tests: the  sources 4U 1728-34 and Sco X-1 are found to yield good power-law fits, with minimum $\\chi^2$-values  lower than 1. The   power-law indices are $b\\simeq1.5$ for the   Atoll  source 4U 1728-34 and  $b\\simeq1.9$  for the Z  source  Sco X-1. A similar power-law index was previously obtained with a slightly high minimum $\\chi^2$-value $\\sim 2$ by Psaltis et al. (1998). As it is known,  Atoll  and  Z sources show distinct properties in their spectra and luminosity (Hasinger \\& van der Klis 1989; van der Klis 2000), and we do not yet know what properties cause the  differences in their  power-law indices. Nevertheless, if the power-law relations with different indices for Atoll and Z sources are confirmed by future detections, then any kHz QPO models discarding the distinctions of the Atoll and Z sources will confront severe arguments. (2) Clearly, obeying such a power-law relation would contradict the constant peak separation and constant (3:2) peak ratio between kHz QPOs,  and in fact the plotted curves in Figure \\ref{fig-12} and Figure \\ref{fig-50} are incompatible with these constant relations. These  conclusions have been previously inferred with smaller samples of kHz QPO data (see, e.g. Psaltis et al. 1998; Psaltis et al. 1999ab; Belloni et al. 2005), but contrary to the suggestions by the simple beat model and   any   model that predict $\\dn$=constant  and   $\\nt =(3/2)\\no$, respectively. In addition, based on  the updated kHz QPO data  of LMXBs we find  that there is no extremely sharp concentration at  a 3:2 peak ratio as indicated by the $\\chi^2$ test in Figure 2c; the ratios are broadly distributed  from $\\sim 1.2$ to $\\sim 3.2$ over a  frequency range of some hundred Hz, as shown in Figure \\ref{fig-12}c. Therefore,   the  non-linear resonance model (see e.g. Abramowicz et al. 2003b; Rebusco 2004) can be consistent with this  distribution of peak ratios. Nevertheless, it is shown in Figure 1c that the value of $\\nt/\\no$ decreases systematically with increasing QPO frequency. (3) In a rough approximation, the kHz QPO frequency correlation seems to be consistent with the predictions by  the model based on the basic general relativistic frequencies around a compact object (Stella \\& Vietri 1999) and those by the model  based on the Alfv\\'en wave oscillation (Zhang 2004), with properly selected  parameters. Both models predicted a peak separation decreasing with increasing QPO frequency, but increasing when the upper frequency is low, for instance less than  $\\sim$700 Hz. Therefore, more kHz QPO detections are needed to confirm the predictions of the models. In conclusion,  if future data still support the conclusions obtained in the paper, they will pose new constraints on models for  explaining kHz QPOs."
    },
    "0601/astro-ph0601632_arXiv.txt": {
        "abstract": "We discuss the physical properties of four quasar jets imaged with the \\emph{Chandra} X-ray Observatory in the course of a survey for X-ray emission from radio jets \\citep{Marshall05}.  These objects have sufficient counts to study their spatially resolved properties, even in the 5 ks survey observations. We have acquired Australia Telescope Compact Array data with resolution matching \\emph{Chandra}. We have searched for optical emission with Magellan, with sub-arcsecond resolution.  The radio to X-ray spectral energy distribution for most of the individual regions indicates against synchrotron radiation from a single-component electron spectrum.  We therefore explore the consequences of assuming that the X-ray emission is the result of inverse Compton scattering on the cosmic microwave background.  If particles and magnetic fields are near minimum energy density in the jet rest frames, then the emitting regions must be relativistically beamed, even at distances of order 500 kpc from the quasar.  We estimate the magnetic field strengths, relativistic Doppler factors, and kinetic energy flux as a function of distance from the quasar core for two or three distinct regions along each jet. We develop, for the first time, estimates in the uncertainties in these parameters, recognizing that they are dominated by our assumptions in applying the standard synchrotron minimum energy conditions.  The kinetic power is comparable with, or exceeds, the quasar radiative luminosity, implying that the jets are a significant factor in the energetics of the accretion process powering the central black hole.  The measured radiative efficiencies of the jets are of order 10$^{-4}$. ",
        "introduction": "Following the remarkable discovery of an X-ray luminous, 100 kpc scale jet in \\pks\\ \\citep{Schwartz00, Chartas00}, we have embarked on a survey to investigate the occurrence and properties of such systems. Initial goals were to assess the frequency of detectable X-ray fluxes from radio-bright jets, to locate good targets for detailed imaging and spectral followup studies, and where possible to test models of the X-ray emission by measuring the broad-band, spatially resolved,  spectral energy distributions (SED) of jets from the radio through the optical to the X-ray band \\citep{Marshall05, Schwartz03a,Schwartz03b}. With an X-ray detection of 12 jets out of the first set of 20 observed, \\citep{Marshall05}, the survey has been successful in meeting these objectives. For four of these jets we have 40 to 130 total X-ray counts in our 5 ks observations.  This suffices to construct broad-band spectral energy distributions, from which we can estimate magnetic fields, particle densities, Doppler beaming factors and kinetic fluxes for independent, spatially distinct emitting regions, using the models of synchrotron radiation and inverse Compton (IC) scattering on the cosmic microwave background (CMB), which we adopt below \\citep{Schwartz03a,Schwartz03b}.  Deeper \\axaf\\  as well as HST observations have been approved for all these sources (\\citet{Perlman04}, and in preparation).  In parallel work, \\citet{Sambruna02, Sambruna04} have undertaken a survey of 17 jetted radio quasars with Chandra and HST, with 10 exhibiting at least one knot in the Chandra images. We will compare our results with theirs in Section~\\ref{param}.   Our survey is described in \\citet{Marshall05} (hereafter Paper I). We selected objects from two parent samples for which radio maps had been obtained with $\\sim$ 1\\arcsec\\ -- 2\\arcsec\\ resolution at 1-10 GHz: a VLA sample (Dec $\\ge$ 0\\arcdeg) of flat spectrum quasars with core flux densities $S_{5 \\rm {GHz}} \\geq 1$ Jy \\citep{Murphy93} and an ATCA survey of flat spectrum Parkes quasars (Dec $<$ -20\\arcdeg) \\citep{Lovell97}, with core flux densities $S_{2.7 \\rm {GHz}} \\geq 0.34$ Jy.  Selection was then made for objects with radio jets which extend beyond 2\\arcsec\\ from the core.  The survey is comprised of sources for which we either anticipated extended structures with significant X-ray fluxes, based on scaling the extended 5 GHz flux using the X-ray to radio flux ratio of \\pks, (subsample ``A''), or else selected by the morphological criteria of a one-sided linear jet (subsample ``B''). PKS~0920-397 and PKS~1202-262 meet the ``A'' criterion, and PKS~0208-512, PKS~1030-357, and PKS~1202-262 meet the ``B'' criterion (Paper I).  The definition of our parent samples by flat radio core spectra tends to select powerful sources with one-sided relativistic pc scale jets beamed toward our line of sight. The 5 or 2.7 GHz selection frequency emphasizes the jet rather than lobe emission.  Section~\\ref{obs} presents the X-ray data and defines distinct spatial regions for further analysis. Section~\\ref{spectral} gives the broad band spectral energy distribution, and Section~\\ref{param} the physical properties deduced by assuming that IC/CMB produces the X-ray emission. Section~\\ref{sec:discussion} discusses implications of the IC/CMB mechanism, and mentions alternate emission mechanisms. In Appendix~\\ref{app:errors} we estimate the systematic uncertainties in the magnetic fields and Doppler factors which we derive, and in Appendix~\\ref{app:kinetic} we present the basis for our calculation of the flux of kinetic energy carried by the jets. ",
        "conclusions": ""
    },
    "0601/astro-ph0601404_arXiv.txt": {
        "abstract": "The dark matter dominated Fornax dwarf spheroidal has five globular clusters orbiting at $\\sim 1$\\,kpc from its centre. In a cuspy CDM halo the globulars would sink to the centre from their current positions within a few Gyrs, presenting a puzzle as to why they survive undigested at the present epoch. We show that a solution to this timing problem is to adopt a cored dark matter halo. We use numerical simulations and analytic calculations to show that, under these conditions, the sinking time becomes many Hubble times; the globulars effectively stall at the dark matter core radius. We conclude that the Fornax dwarf spheroidal has a shallow inner density profile with a core radius constrained by the observed positions of its globular clusters. If the phase space density of the core is primordial then it implies a warm dark matter particle and gives an upper limit to its mass of $\\sim 0.5$\\,keV, consistent with that required to significantly alleviate the substructure problem. ",
        "introduction": "The Fornax dwarf spheroidal is a dark matter dominated satellite orbiting the Milky Way. It has five globular clusters that are at a projected distance from the centre of 1.60, 1.05, 0.43, 0.24 and 1.43\\,kpc \\cite{mackey} as well as further substructure at a projected distance of 0.67\\,kpc \\cite{coleman}. These star clusters move within a dense background of dark matter and should therefore be affected by dynamical friction, causing them to lose energy and spiral to the centre of the galaxy. We will show later that, if Fornax has a cosmologically consistent density distribution of dark matter, the orbital decay timescale of these objects from their current positions is $\\lsim$\\,5\\,Gyr. This is much shorter than the age of the host galaxy, presenting us with the puzzle of why these five globulars have not merged together at the centre forming a single nucleus \\cite{ostriker,tremaine}.  Several groups have studied the origin of nuclei in galaxies: e.g. Lotz et al.\\,(2001) carried out Monte Carlo simulations, which show that some, but not all, of the nuclei of dwarf elliptical galaxies could indeed have formed through coalescence of their globular clusters. Additionally they observed several dE galaxies and found out that within the inner few scale lengths, their sample appeared to be depleted of bright clusters. Oh and Lin\\,(2000) used numerical simulations to show that in dwarf galaxies with relatively weak external tidal perturbations, dynamical friction can lead to significant orbital decay of globular clusters and the formation of compact nuclei within a Hubble timescale.  Oh, Lin and Richer\\,(2000) gave two possible models for the observed spatial distribution of Fornax globulars. One possibility they proposed is that the dark matter consists of massive black holes which transfer energy to the globulars, preventing them from sinking to the centre of the galaxy. Another possibility they investigated was to postulate a strong tidal interaction between the Milky Way and Fornax which also could inject energy into their orbits and the central core of the dSph. This latter idea is probably ruled out due to the proper motion observations of Fornax \\cite{dinescu} which suggest it is already at closest approach on an extended orbit which never takes it close to the Milky Way.  Here we investigate another possibility for the lack of a nucleus in Fornax, namely that the central dark matter distribution has a very shallow cusp or core which dramatically increases the dynamical friction sinking timescale \\cite{hernandez}. This would be inconsistent with dark haloes that form within the CDM cosmology which have cusps steeper than -1 on all mass scales from $10^{-6}$M$_\\odot - 10^{15}$M$_\\odot$ \\cite{dubinski,diemand}.  Controversial evidence for cored mass distributions in dwarf spiral galaxies has been debated for over a decade \\cite{moore2}. The inner structure of spheroidal galaxies is harder to determine, however Kleyna et al.\\,(2003) claimed that the second peak in the stellar number density in the nearby Ursa Minor dwarf spheroidal (UMi dSph), is incompatible with cusped cold dark matter haloes. With their observations they show that this substructure has a cold kinematical signature and that its radial velocity with respect to its host galaxy is very small. Such a cold configuration could only survive intact if the stars orbited within a cored mass distribution where the orbital frequencies are all identical (harmonic potential) and phase mixing does not occur.  The stellar kinematical data for Fornax suggest that it is dark matter dominated with a mass to light ratio of order 20 within its optical extent. Due to the uncertainty on the orbital anisotropy, the mass distribution can only be weakly constrained - the data is consistent with either cusped or cored density distributions \\cite{lokas2}. However the normalisation (or mass within the central 1\\,kpc) is better constrained. In the inner $\\sim$1\\,kpc of a cored halo the mean density is approximately six times lower than in a cusped halo. Furthermore the velocity distribution function of the background particles is hotter than a cusped halo. These facts conspire to significantly increase the dynamical friction timescale in a cored mass distribution.  In this paper we construct cored and cuspy dark matter potentials and calculate orbital decay and sinking times using high resolution numerical simulations together with analytic calculations ~\\cite{chandrasekhar}. The haloes are consistent with the kinematical data for Fornax. We follow circular and eccentric orbits of single and multiple globular clusters. Although many dynamical friction studies have been carried out before \\cite{white,hernquist,capuzzo}, we are not aware of any studies within constant density cores at the resolution used in this paper, although the recent study explored the effects of sinking objects on various cusp structures \\cite{merritt}. In section 2 we present the numerical methods we used, the analytical computation of the sinking times which are compared to our high resolution numerical simulations. In section 3 we discuss our results and draw our conclusions. ",
        "conclusions": "The Fornax dwarf spheroidal has five globular clusters orbiting at a projected radius  of $\\sim$ 1\\,kpc from its centre. Using a cuspy CDM potential with central slope steeper than 0.5 and normalised to match the inferred properties from the kinematical data, we find that these globulars would all sink to the halo centre within 5 Gyrs.  By contrast, we showed that if Fornax has a constant density central core then the dynamical friction time becomes arbitrarily long - the globulars stop sinking at the edge of the core, thus the present position of the {\\it innermost} globular gives a lower limit of the core radius of the dark matter distribution. Since CDM uniquely predicts that all haloes are cusped, this suggests that the dark matter distribution within Fornax is not cold, but may be warm dark matter or some other candidate. Alternatively, the mass distribution has been modified through some exotic dynamical phenomenon such as rapid mass loss \\cite{read}. If the phase space density of the core, measured as $Q\\sim 10^{-5}$M$_\\odot$\\,pc$^{-3}$(kms$^{-1}$)$^{-3}$ in our model, is primordial then it implies a warm dark matter particle of mass $\\sim 0.5$\\,keV \\cite{hogan,dalcanton}. This is consistent with that required to largely solve the substructure problem \\cite{moore3}.  Although the dwarf spheroidals around the Milky Way do not contain prominent nuclei, about 30 percent of dwarf spheroidals (dE's) in clusters are nucleated. If these nuclei form by the merging of star clusters then we must conclude that these galaxies have cuspy mass distributions. This could be due to the fact that transformation to dE via galaxy harassment gives an exponential distribution of stars which usually dominate the central mass distribution. Therefore globulars could sink via friction against the stellar background. Most nucleated dE's are near the cluster centre where harassment is important which supports this idea.  Our numerical simulations show that the standard Chandreskhar estimate of sinking timescales works well for cuspy cores but fails completely for cored mass distributions. In addition we showed that over $10^6$ particles are required within the core region to suppress heating from particle noise \\cite{weinberg}. This is particularly important in a cored mass distribution where the potential minimum is quite shallow. The fact that cored mass distributions do not give rise to dynamical friction is an important result. We believe that this is due to orbit-resonant scattering. This will be the subject of a forthcoming paper \\cite{goerdt}.  We have shown that a natural solution to the timing problems for Fornax's globular clusters is a central dark matter core. A prediction of this model is that the clusters have a well-definied minimum radius. Fornax is not alone in showing indirect evidence for such a core. The UMi dSph galaxy also has substructure which appears to have survived longer than is possible within a cuspy halo. Understanding the origin of these constant density cores could be one of the most exciting challenges facing astronomy in the next few years."
    },
    "0601/astro-ph0601183.txt": {
        "abstract": "{We describe a highly flexible framework to solve 3D radiation transfer problems in scattering dominated environments based on a long characteristics piece-wise parabolic formal solution and an operator splitting method. We find that the linear systems are efficiently solved with iterative solvers such as Gauss-Seidel and Jordan techniques. We use a sphere-in-a-box test model to compare the 3D results to 1D solutions in order to assess the accuracy of the method. We have implemented the method for static media, however, it can be used to solve problems in the Eulerian-frame for media with low velocity fields. } ",
        "introduction": "With the increase in computer power in the last few years, 3D hydrodynamical calculations are becoming increasingly common in astrophysics. Most hydrodynamical calculations treat radiation in a simplified manner, since a full solution of the 3-D non-LTE radiative transfer problem is numerically much more expensive than the hydrodynamical calculation itself, which already stretch the limits of modern parallel computers. In many instances, such as the thermonuclear explosion of a white dwarf (thought to be the progenitor of Type Ia supernova), the goal of the hydrodynamical simulations is to understand the mode of combustion and to handle the effects of turbulence in as realistic manner as possible and radiative transfer effects are ignored. However, in general since the ultimate validation or falsification of the results of sophisticated hydrodynamical modeling will be via comparison with the observed radiation from the astrophysical object being studied, and the radiation strongly affects the physical state of the matter in the atmosphere of the object (where the observed radiation originates) the effect of detailed radiative transfer effects cannot be ignored.  In many multi-dimensional hydrodynamics codes \\cite[for example ZEUS3D,][]{ZEUS3D}, radiative transfer is treated in a simplified matter in order to determine the amount of energy transfered between the matter and radiation (the cooling function) although \\citet{DKC96} presented a full time-dependent 2-D NLTE radiative transfer code, based on a variable Eddington factor method and equivalent two level atom formulation.  Recently \\citet{RPDM05} presented a method of including 3-D radiative transfer into modern 3-D hydrodynamical codes (such as the ASCI FLASH code) but they only treated the solution of the radiative transfer equation in the absence of scattering, that is they their method only treats the formal solution of the radiative transfer equation and not the full self consistent scattering problem where the right hand side of the radiative transfer problem involves the radiation field itself. However, in astrophysical systems, the effect of scattering cannot be ignored and in fact it is due to scattering that the radiation field decouples from the local emission and absorption and the effects of the existence of the boundary are communicated globally over the atmosphere \\citep{mihalas78sa}. It is just this strong non-locality of the radiation field that makes the solution of the generalized radiative transfer problem so computationally demanding.  \\citet{steiner91} presented a 2-D multi-grid method based on the short-characteristics method \\citep{ok87,OAB} and showed that it worked in the case of a purely absorptive atmosphere, i.e., that the formal solution was tractable. \\citet{vath94} presented a 3-D short characteristics method and showed that it had adequate scaling on a SIMD parallel architecture. {In a series of papers 3-D, short characteristics methods for disk systems were presented by \\cite{adam90,hummel94a,hummel94b,hummeldachs92,papkalla95}. Our method is similar in spirit to these works, but we present a more detailed description of the construction of the approximate lambda operator (ALO), and the method of solution of the scattering problem.}  \\citet{BTB97} presented a multi-level, multi-grid, multi-dimensional radiative transfer scheme, using a lower triangular ALO and solving the scattering problem via a Gauss-Seidel method.  \\citet*{vannoort02} presented a method of solving the full NLTE radiative transfer problem using the short characteristics method in 2-D for Cartesian, spherical, and cylindrical geometry. They also used the technique of accelerated lambda iteration (ALI) \\citep{OAB,ok87}, however they restricted themselves to the case of a diagonal ALO. In addition they considered the case including velocity fields, but their method is feasible only in the case of a small velocity gradient across the atmosphere.  We describe below a simple framework to solve three dimensional (3D) radiation transport (RT) problems with a non-local operator splitting method. In subsequent papers we will extend this framework to solve 3D RT problems in relativistically moving configurations. Our method is similar to those described above, although we consider both short characteristics and long characteristics methods for the formal solution of the radiative transfer equation.  We show that long-characteristics produce a significantly better numerical solution in our test cases for a strongly scattering dominated atmosphere. Short-characteristics are known to be diffusive and it is apparent in our numerical results. We also implement a partial parallelization of the method, although we defer a full discussion of the parallelization to future work. ",
        "conclusions": "We have described a framework for solving three-dimensional radiative transfer problems in scattering dominated environments.  The method uses a non-local operator splitting technique to solve the scattering problem. The formal solution is based on a long characteristic piece-wise parabolic procedure. For strongly scattering dominated test cases (sphere in a box) we find good convergence with non-local $\\Lstar$ operators, as well as minimal numerical diffusion with the long characteristics method and adequate resolution. A simple MPI parallelization gives excellent speedups on parallel clusters. In subsequent work we will implement a domain decomposition method to allow much larger spatial grids. { Presently, we have implemented the method for static media, it can be used without significant changes to solve problems in the Eulerian-frame for media with low velocity fields.} {The distribution of matter over the voxels is, in the general 3D case, arbitrary. We chose a spherical test case to be able to compare the results our 1D code.}  In Figs.~\\ref{fig:eps=-4:contour:SC} and \\ref{fig:eps=-4:surface:SC} we compare the results for the $\\epsilon=10^{-4}$ test case using a simple implementation of the  short characteristics method and the long characteristics method used in this paper. The test grid contains $129^3$ voxel and $64^2$ angular points. The high diffusivity of the SC method is evident. Other authors \\cite[]{steiner91,BTB97,TBB95,vath94,ABTB94,vannoort02} have used SC methods in multi-dimensional radiative transport problems.  Short characteristics techniques are faster but require special considerations to reduce numerical diffusion \\cite[]{auer03}. { We may further look into the SC method in later papers.}   We have generalized the operator splitting to include larger bandwidth operators. They lead to faster convergence although they do require more memory and ultimately more computing time. Nevertheless, they will be useful for highly complex problems and we have developed a highly flexible approach to the construction of the $\\Lstar$ operator so that the bandwidth may be set for each spatial point individually as the problem and computational resources require.  We have designed an especially general and flexible framework for 3D radiative transfer problems with scattering. In future papers of this series we will describe its extension to line transfer problems, multi-level NLTE calculations, and differentially moving flows."
    },
    "0601/astro-ph0601297_arXiv.txt": {
        "abstract": "{The evolution of the line width - luminosity relation for spiral galaxies, the Tully-Fisher relation, strongly constrains galaxy formation and evolution models. At this moment, the kinematics of $z>1$ spiral galaxies can only be measured using rest frame optical emission lines associated with star formation, such as H$\\alpha$ and [\\ion{O}{iii}]5007/4959 and [\\ion{O}{ii}]3727. This method has intrinsic difficulties and uncertainties. Moreover, observations of these lines are challenging for present day telescopes and techniques. Here, we present an overview of the intrinsic and observational challenges and some ways way to circumvent them. We illustrate our results with the HST/NICMOS grism sample data of $z \\sim 1.5$ starburst galaxies.  The number of galaxies we can use in the final Tully-Fisher analysis is only three.  We find a $\\sim 2$ mag offset from the local rest frame B and R band Tully-Fisher relation for this sample. This offset is partially explained by sample selection effects and sample specifics. Uncertainties in inclination and extinction and the effects of star formation on the luminosity can be accounted for. The largest remaining uncertainty is the line width / rotation curve velocity measurement. We show that high resolution, excellent seeing integral field spectroscopy will improve the situation. However, we note that no flat rotation curves have been observed for galaxies with $z>1$.  This could be due to the described instrumental and observational limitations, but it might also mean that galaxies at $z>!$ have not reached the organised motions of the present day. ",
        "introduction": "The Tully-Fisher relation (hereafter TFR) is a tight empirical relation between the flat rotation curve (RC) velocity and the luminosity of spiral galaxies  (Tully \\& Fisher \\cite{tullyfisher}). The TFR has been used as a distance estimator and for measurements of the Hubble constant H$_0$ (e.g. Tully \\& Pierce \\cite{tullypierce}).  In addition to its empirical applications, the TFR is interesting in itself because it defines a tight relation between the total (dark matter dominated) mass of spiral galaxies and their luminosity produced by baryons. Furthermore, assuming a stellar mass-to-light ratio $M_*/L$, the stellar masses of galaxies can be calculated from their luminosities and a stellar mass TFR can be derived (Bell \\& De Jong \\cite{bell}). After addition of the gas mass, one obtains the baryonic mass TFR (Verheijen \\cite{verheijen}; Bell \\& De Jong \\cite{bell}; McGaugh et al. \\cite{mcgaugh}). The variations in the dark-to-baryonic mass ratio of galaxies are small and deviations from the baryonic TFR are absent down to very low mass galaxies (McGaugh et al. \\cite{mcgaugh}) although others claim a slight deviation for dwarf spirals (Stil \\& Israel \\cite{stil}). Semi-analytical models of galaxy formation struggle to explain simultaneously the slope, zero point and tightness of the TFR in all optical and near infrared bands (Van den Bosch \\cite{bosch}).  The tight fundamental relation between mass and luminosity is interesting to study in the context of galaxy evolution. The study of the evolving TFR with redshift can provide valuable information on the luminosity evolution of galaxies and the buildup of stellar mass as a function of galaxy mass. As we will show in this paper, the analysis of high redshift TFRs needs careful treatment of observational limits, selection effects, sample definitions and starburst influences, and high resolution high quality spectra.  In the local universe, HI is used to measure the velocity profiles of spiral galaxies. The gas disk in spiral galaxies extends 2-3 times further out than the stellar disk. HI measurements are currently limited to low redshift, beyond redshift $\\sim 0.2$ HI emission has not been observed (Zwaan et al. \\cite{zwaan}) and one has to rely on other kinematic tracers. Another gas tracer could be CO, but CO detections in the high redshift universe are currently limited to very CO bright galaxies, like submillimeter galaxies, quasi stellar objects (QSO) and radio galaxies (Hainline et al. \\cite{hainline} and references therein). One gravitationally lensed Lyman Break Galaxy (LBG) has been detected in CO (Baker et al. \\cite{baker04a}).  A second attempt to detect a LBG in CO failed, although the attempt was on the dustiest LBG known (Baker et al. \\cite{baker04b}).  Bright optical narrow emission lines like H$\\alpha$, H$\\beta$, [\\ion{O}{iii}]5007/4959 and [\\ion{O}{ii}]3727 can also be used to trace the rotation curve of galaxies. Their presence is limited to the stellar disk (or more precisely: the star forming disk). In the local universe, the agreement between HI and HII measurements of RCs is excellent (Courteau \\cite{courteau}). However, these lines shift out of the optical regime at redshifts 0.4 - 1.4. In recent years, high resolution near infrared spectrographs like the Infrared Spectrometer And Array Camera (ISAAC) at the Very Large Telescope (VLT) of the European Southern Observatory (ESO) have become available, opening the window out to redshift 2.4 (for H$\\alpha$) TFR studies.  Examples are Rigopoulou et al. (\\cite{rigopoulou}), who studied massive $z\\sim0.6$ galaxies, Barden et al. (\\cite{barden}) who found an offset from the local B band TFR of $\\sim 1 \\textrm{mag}$ at $z\\sim 0.9$, Lemoine-Busserolle et al. (\\cite{lemoine}) who used gravitational lenses to study two galaxies at $z\\sim 1.9$ and Pettini et al. (\\cite{pettini}), who studied Lyman Break Galaxies at $z\\sim3$.   We used ISAAC to study the kinematics of a sample of $z\\sim1.5$ H$\\alpha$ emitting galaxies and we present the results in this paper as a case study for $z > 1$ TFR studies.   Our sample is a subsample of the McCarthy et al. (\\cite{mccarthy}) HST/NICMOS grism survey sample. McCarthy et al. (\\cite{mccarthy}) surveyed 64 square arc minutes with the slitless NICMOS G141 grism and detected 33 emission line objects with varying 3$\\sigma$ limiting line fluxes down to $1\\times10^{-17} \\textrm{erg s}^{-1} \\textrm{cm}^{-2}$. They argue that the detected emission lines are H$\\alpha$ between redshift 0.75 and 1.9. The H$\\alpha$+[\\ion{N}{ii}]6548/6584 complex is not resolved due to the low spectral resolution ($R\\sim150$) of the grism and therefore contamination by other emission lines (particularly [\\ion{O}{iii}]5007/4959) cannot be excluded and no kinematic information is obtained. This sample is biased to galaxies with large H$\\alpha$ equivalent width, EW(H$\\alpha$), and H$\\alpha$ flux, F(H$\\alpha$), due to the low spectral resolution  of the grism. We chose this sample because it has clear selection criteria and all sources have known H$\\alpha$ fluxes.  Here, we present observations of 9 objects from the McCarthy et al. sample with the ISAAC at the VLT in medium resolution mode ($R\\sim3000-5000$). Our aim was to resolve the H$\\alpha$+[\\ion{N}{ii}] complex (or the [\\ion{O}{iii}]5007/4959 doublet) to confirm redshifts, measure accurate line fluxes and linewidths and if possible also rotation curves. We use this data to present our case study for $z>1$ TFRs.  Hicks et al. (\\cite{hicks}) also performed a follow-up of the HST/NICMOS grism sample. They observed 14 objects aiming to detect emission lines, particularly [\\ion{O}{ii}]3727 in the optical (R/I band) using LRIS at the Keck telescope. They observed in low resolution mode ($R\\sim350-700$ depending on the grating used) and therefore did not obtain any kinematic information. They confirmed the redshift from McCarthy et al. (1999) for 9 out of 14 objects. They explained the non-confirmations by twilight observations or the presence of a nearby bright star (emission lines may very well be not bright enough to detect in these two cases). In two cases, the [\\ion{O}{ii}]3727 line was outside the observed wavelength range and other emission lines like for example \\ion{C}{ii}]2326, \\ion{C}{iii}]1909 and \\ion{Mg}{ii}2800 might have been too faint to detect. The fifth non-detection was explained by reddening or the emission line detected by McCarthy et al. was not H$\\alpha$ but H$\\beta$ or [\\ion{O}{iii}]5007/4959. In the latter case, no bright emission lines are expected in the wavelength range observed.  Our follow-up is complementary in two ways: we observe the objects accessible from the southern hemisphere whereas Hicks et al. observed from the northern hemisphere. Only one object (J0931-0449) is in both samples. Second, we observe at higher resolution, resolving the emission lines.  This paper is organised as follows. The first part of the paper describes the case study data set: the observations of the NICMOS grism sample (Section 2), data reduction and analysis (Section 3), the sample properties (Section 4) and notes on individual galaxies (Section 5). The second part of the paper discusses high redshift TFRs using the earlier discussed dataset as an illustration with a strong focus on pitfalls (Section 6). We conclude with a summary and conclusion in the final section (Section 7).  Throughout this paper, we assume $\\Omega_M=0.3$, $\\Omega_\\Lambda=0.7$ and $H_0=70 \\textrm{ km s}^{-1}\\textrm{Mpc}^{-1}$. All magnitudes in this paper are Vega magnitudes. ",
        "conclusions": "We studied the challenges in measuring $z > 1$ TFR using emission line galaxies using a sample of H$\\alpha$ emitting galaxies. We conclude:  - We confirm only 5 out of 9 emission lines found by McCarthy et al. (\\cite{mccarthy}). Of the 5 remaining sources, one is a broad line AGN, one probably a narrow line AGN. Our conclusions are therefore based one a very small number of sources.  - Without reliable inclinations and extinctions, an ensemble averaged simultaneous inclination and extinction correction can be made assuming the relation between extinction and inclination follows that of local galaxies, which is in general a small correction because galaxies move approximately along the TFR.  - The most difficult challenge for high redshift TFR studies, are the velocity measurements. Besides the observational challenges and limitations, there is also the fundamental question if what we observe is comparable to what we observe in the local galaxies (for example the question if the emission lines extent to the flat part of the RC). For the current data set, this leaves considerably uncertainties in the measured line widths, which cannot be quantified.  - The high mass end of the local TFR is not sensitive to the sample used, therefore the choice of the local reference sample is not important for the high mass end of the high redshift TFR.  - Extending the study of the high redshift TFR to low mass galaxies will be very difficult due to the selection effects in velocity and magnitude: it would require very high spectral resolution observations of very faint galaxies.  - Star formation increases the luminosity of galaxies, but the effect on $\\log(L)$ is negligible for high mass galaxies with significant older population. Therefore star forming galaxies can be TFR-representative for all galaxies, and whether they are or not can be checked.   For the NICMOS galaxies, we measured a $\\sim 2$ mag offset of the $z \\sim 1.5$ rest frame B starburst TFR with respect to the local TFR. This offset is robust against the effects of inclination and extinction. The linewidths indicate sufficiently massive galaxies for TFR studies, the results is therefore also robust for the choice of local comparison sample. However, we cannot prove or disprove that the all linewidths are due to rotation. Moreover,  we cannot answer the question if the star burst luminosity dominates the total luminosity or not. This is due to the extreme nature of the galaxies and the sparseness of data for this sample."
    },
    "0601/astro-ph0601068_arXiv.txt": {
        "abstract": "Using 2D models of rotating stars, the interferometric measurements of $\\alpha$~Eri and its fundamental parameters corrected for gravitational darkening effects we infer that the star might have a core rotating 2.7 times faster than the surface. We explore the consequences on spectral lines produced by surface differential rotation combined with the effects due to a kind of internal differential rotation with rotational energies higher than allowed for rigid rotation which induce geometrical deformations that do not distinguish strongly from those carried by the rigid rotation. ",
        "introduction": "\\label{model}  An initial rigid rotation in the ZAMS switches in some $10^4$ yr into an internal differential rotation (IDR) (Meynet \\& Maeder 2000, MM2000). Nothing precludes then that IDR be present since the pre-main sequence phases implying a total angular momentum higher than allowed for rigid rotation. Inspired by the internal rotation laws obtained in MM2000, we adopt the step-like rotational law: $\\Omega(\\varpi)\\!=$ $\\Omega_{\\rm co}[1-p.e^{-a.\\varpi^b}]$, where $\\varpi\\!=$ distance to the rotation axis; $\\Omega_{\\rm co}\\!=$ core angular velocity; $p$ determines $\\Omega_{\\rm co}/\\Omega_{\\rm surf}$; $a\\!=\\!c(r_{\\rm co})$ where $r_{\\rm co}\\!=\\!R_{\\rm core}/R_{\\rm eq}\\!=$ distance at which $\\Omega\\!=$ $(1/2)\\Omega_{\\rm co}$; $b$ determines the steepness of the drop from $\\Omega_{\\rm co}$ to $\\Omega_{\\rm surf}$. From MM2000 we adopt $b\\!=\\!5$, so that the rotational law is stable against axi-symmetric perturbations [$\\partial{j}/\\partial\\varpi\\!>\\!0$; $j\\!=\\!\\Omega(\\varpi)\\varpi^2$] for $p\\!\\la\\!p_{\\rm max}(b\\!=\\!5)\\!=\\!0.73$. The geometrical deformation of stars is obtained by solving the 2D Poisson equation (Clement 1979): \\begin{equation} \\Delta\\Phi_G(\\theta,\\varpi) = -4\\pi\\rho(\\theta,\\varpi) \\label{poisson} \\end{equation} \\noindent using barotropic relations $P=P[\\rho(r)]$ of stellar interiors without rotation in different evolutionary stages (Zorec et al. 1989, Zorec 1992). Fig.~\\ref{f1}a shows stellar meridian cuts for $\\eta=$ $\\Omega^2_{\\rm surf}R^3_{\\rm eq}/GM$, several values of $p$ and $r_{\\rm co}=0.2$. Fig.~\\ref{f1}b shows $R_{\\rm eq}/R_o$ ($R_o=$ stellar radius at rest) against $\\tau_d=$ ${\\cal K}/|{\\cal W}|$ (${\\cal K}=$ rotational energy; ${\\cal W}=$ gravitational potential energy).\\par We considered that the apparent polar radius $R_{\\rm po}^{\\rm app}$ of $\\alpha$ Eri is well determined (Domiciano de Souza et al. 2003, Vinicius et al. 2005), while the equatorial radius $R_{\\rm eq}$ is obtained by asking the equivalent circular (determined spectrophotometrically and from visible interferometry) and the actual elliptical areas of the apparent stellar disc be equal. The relation between apparent and true stellar radii ratios: $R_{\\rm po}^{\\rm app}/R_{\\rm eq}=$ $\\{1-[1-(R_{\\rm po}^{\\rm true}/R_{\\rm eq})^2]^2\\sin^2i\\}^{1\\over2}$ leads to $F(p,\\tau)$ that can be evaluated with observed quantities and whose theoretical counterpart is (see Fig.~\\ref{f1}c): \\begin {equation} F(p,\\tau) = \\left[1-\\left(R_{\\rm po}/R_{\\rm eq}\\right)^2\\right]/\\left(V_{\\rm e}/V_{\\rm c,r}\\right)^2 \\label{funcf} \\end{equation} \\noindent where $V_{\\rm eq}=$ equatorial velocity; $V_{\\rm c}=$ critical velocity at rigid rotation. Thanks to interferometric data, we can look for models that reproduce the following observables: $R_{\\rm eq}/R_{\\odot}$, $R_{\\rm eq}/R_{\\rm po}$, $V\\!\\sin i$ corrected for gravity darkening effects (Fr\\'emat et al. 2005) and $F(p,\\tau_d)$. Thus, we infer: \\begin{equation} \\left.\\begin{array}{lclrclrcl} p&=& 0.624\\pm0.001 &      \\tau &=& 0.014\\pm0.001 & i&=& 52^o\\\\ \\eta&=& 0.69\\pm0.07   & V_{\\rm eq}&=& 308\\pm16 {\\rm km/s} & \\Omega_{\\rm co}/\\Omega_{\\rm surf}& = & 2.7 \\\\ \\end{array} \\right\\} \\end{equation} \\noindent For rigid rotation it is $p=0.0$ and $\\tau_d\\!\\la\\!0.015$. \\begin{figure}[] \\centerline{\\psfig{file=jzorec2_fig1.eps,width=13.5truecm,height=5truecm}} \\caption[]{a) Meridian cuts of models with $b=5$, $r_{\\rm co}=$ 0.2, $\\eta=0.8$ and several values of $p$. b) Equatorial radius $R_{\\rm eq}/R_o$ against $\\tau_d=$ $K/|W|$ and $p$. c) Function $F(p,\\tau_d)$ against $\\tau_d$ and $p$} \\label{f1} \\end{figure} ",
        "conclusions": ""
    },
    "0601/astro-ph0601542_arXiv.txt": {
        "abstract": "s{A hypercritical value for the magnetic field is determined, which provides the full compensation of the positronium rest mass by the binding energy in the maximum symmetry state and disappearance of the energy gap separating the electron-positron system from the vacuum. The compensation becomes possible owing to the falling to the center phenomenon. The structure of the vacuum is described in terms of strongly localized states of tightly mutually bound (or confined) pairs. Their delocalization for still higher magnetic field, capable of screening its further growth, is discussed.}   It is accepted that magnetic fields are stable in pure quantum electrodynamics (QED), and other interactions (weak or strong) or magnetic monopoles have to be involved to make the magnetized vacuum unstable \\cite{NO78}. In this Talk (see Ref. \\cite{SU05} for a detailed version) we establish, within the frame of QED, that there exists a hypercritical value of the magnetic field \\begin{eqnarray}\\label{final} B^{(1)}_{\\rm hpcr}\\simeq \\frac{m^2}{4e}\\exp\\left\\{\\frac{\\pi^{3/2}} {\\sqrt{\\alpha}}+2C_{\\rm E}\\right\\}\\simeq 10^{28}B_0\\simeq 10^{42}\\rm G\\,. \\end{eqnarray}  Here $\\alpha=e^2/4\\pi\\simeq 1/137,$ $B_0=m^2/e=4.4\\times 10^{13}$G, $m$ is the electron mass. The value (\\ref{final}) is less than the magnetic field $\\sim (10^{47}-10^{48})$ G expected to be present near superconductive cosmic strings \\cite{W95} and the one ($\\sim 10^{47}$G) produced at the beginning of the inflation \\cite{KK89}. The hypercritical magnetic field leads to the shrinking of the energy gap between an electron and positron that takes place due to the falling to the center. The latter is caused by the ultraviolet singularity of the photon propagator (or Coulomb potential) \\cite{goldstein} plus the dimensional reduction in the magnetic field \\cite{loudon} - \\cite{gusynin}. We discuss the vacuum structure around the hypercritical magnetic field and its possible decay under a further growth of the magnetic field - that may cause its screening.  We rely on  the theory of the falling to the center developed in \\cite{shabad} that implies deviations from the standard quantum theory manifesting themselves when extremely large electric fields near the  singularity become important. In that theory the singularity in the Schr\\\"{o}dinger-like equation yields a singular measure in the scalar product and hence the geometry of a black-hole-like object. (Stress, that the geometry is induced: no interaction of  gravitational origin is included.)  We proceed from the (3+1)-dimensional Bethe-Salpeter (BS) equation in an approximation, which is the ladder approximation once the photon propagator (in the coordinate space) is taken in the Feynman gauge: $D_{ij}(x)=g_{ij}D(x^2)$, $x^2\\equiv x_0^2-{\\bf x}^2$, $g_{ii}=(1,-1-1-1)$. In an asymptotically strong magnetic field this equation may be written \\cite{SU05} in the following (1+1)-form  covariant under the Lorentz transformations along the axis 3: \\begin{eqnarray}\\label{closed1}({{\\rm i}\\overrightarrow{\\hat{ \\partial_\\|}}- \\frac{\\hat{P_\\parallel}}2-m}) \\Theta(t,z)(-{\\rm i}\\overleftarrow{\\hat{ \\partial_\\parallel}}- \\frac{\\hat{P_\\parallel}}2-m)\\quad\\quad\\quad \\nonumber\\\\ = {\\rm i} 8\\pi\\alpha~\\sum_{i=0,3}D\\left(t^2-z^2-\\frac{P_\\perp^2}{(eB)^2}\\right)g_{ii}\\gamma_i \\Theta(t,z)\\gamma_i, \\end{eqnarray} where $\\Theta (t,z)$ is the 4$\\times 4$ (Ritus transform of) BS amplitude, $t=x^{\\rm e}_0- x^{\\rm p}_0$ and $z=x^{\\rm e}_3-x^{\\rm p}_3$ are the differences of the coordinates of the electron (e) and positron (p) along the time $x_0$ and along the magnetic field ${\\bf B}=(0,0,B_3=B)$. $P_\\|$ and $P_\\perp$ are projections of the total (generalized) momentum of the positronium onto the (0,3)- subspace and the (1,2)-subspace. Only two Dirac gamma-matrices, $\\gamma_{0,3},$ are involved, $\\hat{ \\partial_\\|}=\\partial_t\\gamma_0- \\partial_z\\gamma_3$, $\\hat{P_\\|}=P_0\\gamma_0-P_3\\gamma_3$.  Equation (\\ref{closed1}) is valid in the  domain, where the argument of  $D$ is greater than the electron  Larmour radius squared $(L_B)^2=(eB)^{-1}$. When $B=\\infty$, this domain covers the whole exterior of the light cone $z^2-t^2\\geq 0$.The argument of the original photon propagator $(x^{\\rm e}-x^{\\rm p})^2$ has proved to be replaced in (\\ref{closed1}) by $t^2-z^2-(\\widetilde{x}_\\perp^{\\rm e}-\\widetilde{x}_\\perp^{\\rm p})^2 =t^2-z^2-P_\\perp^2/(eB)^2$, where $\\widetilde{x}_\\perp^{\\rm e,p}$ are the center of orbit coordinates of the two particles in the transversal plane. Now that after the dimensional reduction this subspace no longer exists these substitute for the transversal particle coordinates themselves: $\\widetilde{x}_\\perp^{\\rm e,p}$are not coordinates, but quantum numbers of the transverse momenta. The mechanism of replacement of a coordinate by a quantum number is the same as % in \\cite{ShUs}.  In deriving equation (\\ref{closed1}) the expansion over the complete set of Ritus matrix eigenfunctions \\cite{ritus} was used in \\cite{SU05} that accumulate the dependence on the transversal spacial and spinorial degrees of freedom. This expansion yields an infinite set of equations, where different pairs of Landau quantum numbers $n^{\\rm e}$, $n^{\\rm p}$ are entangled, Eq. (\\ref{closed1}) being the equation for the ($n^{\\rm e}=n^{\\rm p}=0$)-component that decouples from this set in the limit $B=\\infty$.  In the ultra-relativistic limit $P_0=P_3=P_\\perp=0$ equation (\\ref{closed1}) is solved by the most symmetric Ansatz $\\Theta=I\\Phi$, where $I$ is the unit matrix.  The ultraviolet singularity  on the light cone ($x^2=0$) of the free photon propagator, $D(x^2)=-(\\rm i/4\\pi^2)(1$/$x^2)$, after this expression is used in eq. (\\ref{closed1}),  leads to falling to the center in the Schr$\\ddot{\\rm o}$dinger-like differential equation\\begin{eqnarray}\\label{last3} -\\frac{{\\rm d}^2\\Psi(s)}{{\\rm d} s^2}+\\left(m^2-\\frac 1{4s^2}\\right)\\Psi(s)=\\frac{4\\alpha}{\\pi}\\frac 1{s^2}\\Psi(s),\\quad (eB)^{-1/2}\\ll s\\leq \\infty, \\end{eqnarray} to which the radial part of equation (\\ref{closed1}) is reduced in the most symmetrical case, when the wave function $\\Phi (x)=s^{-1/2}\\Psi(s)$ does not depend on the hyperbolic angle $\\phi$ in the space-like region of the two-dimensional Minkowsky space, $t=s\\sinh\\phi,\\; z=s\\cosh\\phi,\\;s=\\sqrt{z^2-t^2}$.  The solution that decreases at $s\\rightarrow \\infty$ is given by the McDonald function with imaginary index: $\\Psi(s)=\\sqrt{s}\\;K_\\nu(ms),\\quad \\nu= {\\rm i}\\,2\\sqrt{\\alpha/\\pi}\\simeq 0.096\\,{\\rm i}\\,,$ oscillating when $s\\rightarrow 0.$ The falling to the center \\cite{QM} holds for any positive $\\alpha$.  According to \\cite{shabad}, the singular equation (\\ref{last3}) should be considered as the generalized eigenvalue problem with respect to $\\alpha$. The operator in the left-hand side is self-adjoint provided the standing wave condition is imposed, \\begin{eqnarray}\\label{stand} \\left.\\Psi (s)\\right |_{s=L_B}=0 \\,, \\end{eqnarray} that treats the Larmour radius as the lower edge of the normalization box. The discrete eigenvalues $\\alpha_n(L_B)$ condense in the limit $B=\\infty$ to become a continuum of states that form the (rigged) Hilbert space of vectors orthogonal with the singular measure $s^{-2}{\\rm d}s$. The latter fact allows to normalize them to $\\delta$-functions and interpret as free particles emitted and absorbed by the singular center. As long as the Larmour radius is much smaller than the only characteristic length in eq.(\\ref{last3}), electron Compton length, $L_B\\ll m^{-1}\\simeq 3.9\\times 10^{-11}$ cm., the small-distance asymptotic regime is reached, and nothing \"behind the horizon,\" $s<L_B,$ - where the two-dimensional equations (\\ref{closed1}) and (\\ref{last3}) are not valid - may affect the problem. In this way the existence of the limit of vanishing regularization length, impossible in the standard theory, is achieved.   For sufficiently large magnetic fields $L_B$ becomes so small that the region where $K_\\nu(ms)$ oscillates gets inside the domain of validity of equation (\\ref{last3}). The value of the magnetic field when this  happens for the first node is just (\\ref{final}). The corresponding Larmour radius is about fourteen orders of magnitude smaller than $m^{-1}$ and makes $\\sim 10^{-25}$ cm. Eq. (\\ref{final}) tells how large the magnetic field should be in order that the boundary problem (\\ref{last3}), (\\ref{stand}) might have a solution, in other words, that the point $P_0={\\bf P}=0$ might belong to the spectrum. Therefore, if the magnetic field exceeds the first hypercritical value  the positronium ground state exists \\footnote{A relation like (\\ref{final}) is present in \\cite{gusynin}. There, however, a different problem is studied and, correspondingly, a different meaning is attributed to that relation: it expresses the mass gained dynamically - in the course of spontaneous breakdown of the chiral invariance in massless QED - by a massless Fermion in terms of the magnetic field applied to it. Later, in \\cite{gusynin2} the authors revised that relation in favor of a different approximation. Supposedly, the revised relation may be of use in the problem of ultimate magnetic field dealt with here.} with its rest energy compensated for by the mass defect.  The ultra-relativistic state $P_\\mu=0$ has the internal structure of what was called a \"confined state\", belonging to kinematical domain called \"sector III\" in \\cite{shabad}, i.e. the one whose wave function behaves as a standing wave combination of free particles near the lower edge of the normalization box and decreases as $\\exp (-ms)$ at large distances. The effective \"Bohr radius\", i.e. the value of $s$ that provides the maximum to the wave function  makes $s_{\\rm max}=0.17 m^{-1}$. This is certainly much less than the standard Bohr radius $(e^2m)^{-1}$. Taken at the level of 1/2 of its maximum value, the wave function is concentrated within the limits  $0.006 ~m^{-1}<s<1.1~ m^{-1}$. But the effective region occupied by the confined state is still much closer to $s=0$, since - in accord with the aforesaid - the probability density of the confined state is the wave function squared $weighted ~with~ the~ measure$ $s^{-2}{\\rm d} s$ $singular~ in~ the~ origin$ \\cite{shabad} and is hence concentrated near the edge of the normalization box $s\\simeq 10^{-25}$cm, and not in the vicinity of the maximum of the wave function. The electric fields at such distances are about $10^{43}$ Volt/cm. Certainly, there is no evidence that the standard quantum theory should be valid under such conditions. This fact encourages the use of a theory that admits deviations from the standard quantum theory that close to the singularity point.  At $B=B^{(1)}_{\\rm hpcr}$ the total energy and momentum of a positronium in the ground state are zero. This state is not separated from the vacuum by an energy gap, and has maximum symmetry in the coordinate and spin space. Hence, it may be related to the vacuum and responsible for its structure.  At $B>B^{(1)}_{\\rm hpcr}$ the eigenvalues of the BS equation (\\ref{closed1}) for the total 2-momentum components $P_{0,3}$ of the e$^+$-e$^-$ system are expected to shift into the space-like region (we keep $P_\\perp=0$), whereas for $B<B_{\\rm hpcr}^{(1)}$ the c.m. 2-momentum of the real pair was, naturally, time-like. If $P_{0,3}\\neq 0,$ at least for far space-like $P_\\parallel$, the situation can be modelled by the same equation as (\\ref{last3}), but with the large negative quantity $~~m^2+P_\\|^2/4~~$ substituted for  $m^2$. Then the wave function would contain two oscillating exponentials for large space-like intervals, \\begin{eqnarray}\\label{hankel} \\exp\\left\\{\\pm{\\rm i} s\\sqrt{\\left|m^2+\\frac{P_\\|^2}4\\right|}\\right\\}\\exp\\left\\{{\\rm i} P_\\|\\frac{x_{\\rm e}+x_{\\rm p}}2\\right\\},\\quad  s\\gg (-P_\\|^2)^{-1/2},\\end{eqnarray} and two oscillating exponentials $\\exp (\\pm2{\\rm i}\\ln s\\,\\sqrt{\\alpha/\\pi})$ for small ones, $s\\sim L_{\\rm B}$. In the Lorentz frame, where $P_0=0, P_3\\neq 0$, and with the time arguments in the two-time BS amplitude equal to one another: $x_0^{\\rm e}=x_0^{\\rm p}$, the solution oscillates along the magnetic field with respect to the relative coordinate $x_3^{\\rm e}-x_3^{\\rm p}$ (mutually free particles) and with respect to the c.m. coordinate $x_3^{\\rm e}+x_3^{\\rm p}$ (vacuum lattice).  We are now in the kinematical domain called sector IV, or deconfinement sector in Refs.\\cite{shabad}. Here the constituents are free at large intervals and near the singular point $s=0$. The wave incoming from infinity is partially reflected, and partially penetrates to the singular point, the probability of creation of the delocalized (free) states being determined by the barrier transmission coefficient \\cite{shabad}. Such states may exist if one succeeds to satisfy e.g. periodic conditions, to be imposed on the lower and upper boundaries of the normalization volume, instead of  condition (\\ref{stand}), appropriate in sector III. The possibility to obey them is provided again by the falling to the center. Now, the delocalized states in two-dimensional Minkowsky space correspond to electron and positron that circle along Larmour orbits with vanishing radii in the plane orthogonal to the magnetic field and simultaneously perform, when the interval between them is large, a free motion along the magnetic field. They have magnetic moments and seem to be capable of screening the magnetic field. This provides the mechanism that may prevent the classical magnetic field from being larger than  a {\\it second hypercritical field}, for which the delocalization first appears. No sooner than the delocalized states are found in our present problem one may definitely claim the instability of the vacuum with the second hypercritical magnetic field or - which is the same - the instability of such field under the pair creation that might provide the mechanism for its diminishing. For the present, we state that the first hypervalue (\\ref{final}) is such a value of the magnetic field, the exceeding of which would already cause restructuring of the vacuum. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601418.txt": {
        "abstract": "{We studied the circumstellar environment of the  B{[}e{]} star HD 45677 through the analysis of the emission lines from ionized metals. We used the statistical approach of the self absorption curve method (SAC) to derive physical parameters of the line emitting region. The Fe II and Cr II double-peaked emission line structure is explained by the presence of a thin absorption component red shifted by $\\sim$ 3 $\\mathrm{km\\, s^{-1}}$. This absorption component can be interpreted geometricaly as being due to infalling material perpendicularly to the disk seen nearly pole-on, as indicated by the emission line structure. The Cr II and Fe II emission lines have a complex structure with two (narrow and broad) components, of 45 and 180 $\\mathrm{km\\, s^{-1}}$ FWHM for the permitted lines and 25 and 100 $\\mathrm{km\\, s^{-1}}$ FWHM for the forbidden ones, respectively. From our best data set of 1999, we obtained a Boltzmann-type population law whose excitation temperature is $3\\,900_{-600}^{+900}\\,\\mathrm{K}$ and $3\\,150_{-300}^{+350}\\,\\mathrm{K}$ for the narrow component lower and upper levels respectively. We obtained an excitation temperature of $3\\,400_{-300}^{+350}\\,\\mathrm{K}$ for the broad component upper levels. The forbidden lines are found to be formed in the outer regions with higher excitation temperatures of $10\\, 500 \\pm 1\\, 000 K$ and $8\\, 000 \\pm 1\\, 500 K$ for the narrow and broad components respectively in 1999. Our results are consistent with line formation in a rotating disk, around a young star. In the framework of a very simplified geometrical model, we argue that the narrow components are principaly emitted by an optically thin disk seen nearly pole-on, in a region whose minimum radius is estimated to be $4\\,10^{12}cm$, while the broad ones are formed in a disk-linked wind.} ",
        "introduction": "HD 45677 (FS CMa) is generally recognized as an early type star, usually classified as B{[}e{]}, or more precisely as showing the B{[}e{]} phenomenon, even if it is difficult to determine its exact evolutionary stage (de Winter et al. \\cite{dWvdA97}; Lamers et al.\\cite{LZWHZ98}).  The presence of a disk, and even more of a bipolar environment, was inferred by polarization measurements (Coyne \\& Vrba \\cite{CV76}; Schulte-Ladbeck et al. \\cite{SLSNCABBCMMTW93}) and by UV spectroscopic measurements of accreting gas (Grady et al. \\cite{GBS93}). These last authors agreed on the presence, around the star, of an actively accreting circumstellar, protoplanetary disk, and the presence of low-velocity absorption profiles in Fe {\\scriptsize II} lines, among other species, while optical spectroscopic measurements of strong absorption cores of a few lines are commented on by de Winter et al. (\\cite{dWvdA97}).  In the optical HD 45677 is characterised by a rich emission line spectrum, including the Balmer lines of hydrogen, and both permitted and forbidden lines of neutral and ionized metals (Mg {\\scriptsize I}, Mg {\\scriptsize II}, Na {\\scriptsize I}, Cr {\\scriptsize II}, Si {\\scriptsize II}, Mn {\\scriptsize II}, Ti{\\scriptsize ~II}, {[}Ni{\\scriptsize ~II}{]}, {[}S{\\scriptsize ~II{]}} and especially Fe {\\scriptsize II} and {[}Fe {\\scriptsize II}{]}).  Israelian et al. (1996) found a surface gravity of $\\log\\,\\mathrm{g}=3.85\\pm0.05$ (luminosity class V) from the photospheric wings of H$\\gamma$ and H$\\delta$, while examination of the He {\\scriptsize I} absorption lines suggested a spectral class of B2~($\\mathrm{\\mathrm{T_{eff}=}22\\,000\\pm1\\,500\\, K}$). The effective temperature was less certain than the surface gravity, due to a circumstellar contribution to the He {\\scriptsize I} lines. Israelian et al. (\\cite{Ia96}) also found that the photospheric Si {\\scriptsize II} 4128 and 4130 $\\textrm{\u00c5}$~lines were not blended, and could be only fitted by models with a projected rotational velocity $v\\,\\sin i$ of less than 70 $\\mathrm{km\\, s^{-1}}$.  Cidale et al. (\\cite{CZT01}) used the method of Barbier, Chalonge \\& Divan, based on the magnitude and position of the Balmer \\textit{photospheric} (not circumstellar) discontinuity. They obtained a spectral type of B2IV-V, corresponding to $\\mathrm{\\mathrm{T_{eff}=}21\\,500\\pm300\\, K}$ and $\\log\\, g=3.89$, in agreement with the previous determinations.  De Winter \\& van den Ancker (\\cite{dWvdA97}) studied the behaviour of HD 45677 during the last 25 years, and concluded that the large photometric variations of the star are best explained by a variable obscuration from circumstellar grains. They also derived a stellar instrinsic visual--ultraviolet energy distribution close to that of a B2 V star, with an infrared excess due to the dust emission. The recent spectral variations were discussed by Israelian et al. (\\cite{Ia96}) and Israelian \\& Musaev (\\cite{IM97}).  In the present study, we analyse the emission line spectrum in order to get an insight into the physical processes of line excitation and geometrical conditions of their forming regions. It is clearly impossible at the present stage to calculate a detailed synthetic spectrum for this badly known object, so in order to analyse the spectrum, we use the statistical approach offered by the Self-Absorption Curve method (Friedjung \\& Muratorio \\cite{FM87}; Muratorio \\& Friedjung \\cite{MF88}; Baratta et al. \\cite{Ba98}; Kotnik-Karuza et al. \\cite{KFS02}). ",
        "conclusions": "The SAC method has given consistent results. We have obtained the properties of what is probably a disk of HD 45677, including limits on a characteristic column density and a characteristic radius of the Fe {\\scriptsize II} emitting region. These estimates require reasonable assumptions, one of which is a true Boltzmann distribution of the populations of the lower even metastable levels and another that the Boltzmann--like law of the populations of the odd upper levels of strong optical lines are not more than what would be expected from a smooth fit between these odd and even population laws. We have demonstrated that the well known split emission line profile is in fact due to absorption components like in the ultraviolet. We have demonstrated that emission lines have narrow and broad components. We have shown that the double emission profile can be explained with a disk plus a flow (perhaps a wind) moving in or out of the disk. The Fe{\\scriptsize ~II} forbidden lines show a similar two component structure, but the smaller width of each component is consistent with their production in a more external part of the disk, where the Keplerian velocity is lower. The examination of the fluxes of the narrow {[}Fe {\\scriptsize II{]}} components has shown that the metastable $\\mathrm{Fe}{}^{\\textrm{+}}$ levels follow a Boltzmann--type law with a mean excitation temperature around 11~000~K. Such a value may actually represent the electron temperature of the {[}Fe {\\scriptsize II}{]} emitting region, if the metastable levels are in equilibrium with the ground term through electron collisions. This is probably the case of HD 45677, since we expect the electron density to be large enough (N$_e \\ge 10^6$ cm$^{-3}$) for bringing the metastable levels to LTE (Nussbaumer \\& Storey \\cite{NS88}). The low electron temperature with respect to the stellar surface temperature may be due to cooling of the {[}Fe {\\scriptsize II}{]} region by the large line emission of this ion. Indeed, low electron temperatures were predicted by Verner et al. (\\cite{VVKFHF99}) and Verner et al. (\\cite{VGBJID02}) for the $\\eta$ Car {\\it Weigelt blobs} that are strong $\\mathrm{Fe}{}^{\\textrm{+}}$ emitters. Fairly low {[}Fe {\\scriptsize II}{]} excitation temperatures were derived for $\\eta$ Car (12000 K, Viotti et al. (\\cite{VRB99}) and for the symbiotic star RR Tel (6600 K Kotnik-Karuza et al. \\cite{KFS02}).  The narrow components of the Fe {\\scriptsize II} permitted emission lines in HD 45677 are slightly, but significantly broader than the corresponding {[}Fe {\\scriptsize II}{]} components. From the SAC analysis of the Fe {\\scriptsize II} emission lines we have derived an $\\mathrm{Fe}{}^{\\textrm{+}}$ lower level excitation temperature much smaller than that derived from the {[}Fe {\\scriptsize II}{]} lines, namely around 3~900~K. We take these results as evidence of different (but possibly partially overlapping) emitting regions. The derived excitation temperatures of the lower and upper levels of the permitted Fe {\\scriptsize II} lines are much smaller than that expected if $\\mathrm{Fe}{}^{\\textrm{+}}$ were in thermodynamical equilibrium with the local plasma. Verner et al. (\\cite{VGBJID02}) calculated the formation of singly ionized iron in the Weigelt blobs of $\\eta$ Car. According to them Fe$^{\\textrm{+}}$ emission is the most important coolant in the blobs. Fe {\\scriptsize II} emission originates at electron temperatures $5\\,000\\, K\\leq T_{e}\\leq7\\,500\\, K$. This can also be seen by comparing their figs, 8 and 9. We argue that for HD 45677 the upper levels of the permitted low excitation transitions are mostly populated by the diluted stellar UV radiation followed by spontaneous decay, and that in the Fe {\\scriptsize II}  region the radiation density is high enough to make the radiation induced transitions more important than electron collisions. Similar considerations can be made for Fe{\\scriptsize ~II} and {[}Fe{\\scriptsize ~II}{]} broad components, as their emitting region parameters appear fairly close to the narrow components ones (see Table.~\\ref{TAB-fit_N} and Table.~\\ref{TAB-fit_W} for Fe{\\scriptsize ~II} lines and Table.~\\ref{TAB-fit_F} and Table.~\\ref{TAB-fit_G} for {[}Fe{\\scriptsize ~II}{]} lines).  The Fe {\\scriptsize II} metastable level, as well as the lower level (Table.~\\ref{TAB-fit_N}) excitation temperatures are suspected to increase with time. The high excitation anomalous multiplets may indicate the presence of inhomogeneities.  The shape of the iron broad component is insufficiently known now to decide about the structure of the non-disk component. Nevertheless the similarity of the physical parameters derived for narrow and broad components (T$_{\\textrm{exc}}$, $\\log\\frac{N_{O}}{g_{0}}$) suggest that both emitting regions, if not the same (different FWHM) are physically linked.  Our identification of the broad components, and the suspected variability of both emitting regions physical parameters should stimulate in future new higher quality observations (especially in the blue) in order to study more precisely the shape of the emission lines."
    },
    "0601/astro-ph0601224_arXiv.txt": {
        "abstract": "We investigate the statistical properties of cosmic baryon fluid in the nonlinear regime, which is crucial for understanding the large-scale structure formation of the universe. With the hydrodynamic simulation sample of the Universe in the cold dark matter model with a cosmological constant, we show that the intermittency of the velocity field of cosmic baryon fluid at redshift $z=0$ in the scale range from the Jeans length to about 16 h$^{-1}$ Mpc can be extremely well described by She-Leveque's universal scaling formula. The baryon fluid also possesses the following features: (1) for volume weight statistics, the dissipative structures are dominated by sheets, and (2) the relation between the intensities of fluctuations is hierarchical. These results imply that the evolution of highly evolved cosmic baryon fluid is similar to a fully developed turbulence. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601538_arXiv.txt": {
        "abstract": "In this manuscript I review the mathematics and physics that underpins recent work using the clustering of galaxies to derive cosmological model constraints. I start by describing the basic concepts, and gradually move on to some of the complexities involved in analysing galaxy redshift surveys, focusing on the 2dF Galaxy Redshift Survey (2dFGRS) and the Sloan Digital Sky survey (SDSS). Difficulties within such an analysis, particularly dealing with redshift space distortions and galaxy bias are highlighted. I then describe current observations of the CMB fluctuation power spectrum, and consider the importance of measurements of the clustering of galaxies in light of recent experiments. Finally, I provide an example joint analysis of the latest CMB and large-scale structure data, leading to a set of parameter constraints. ",
        "introduction": "The basic techniques required to analyse galaxy clustering were introduced in the 70s \\cite{peebles73}, and have been subsequently refined to match data sets of increasing quality and size. In this manuscript I have tried to summarise the current state of this field. Obviously, such an attempt can never be complete or unique in every detail, although it is still worthwhile as it is always useful to have more than one source of information. An excellent alternative viewpoint was recently provided by Hamilton \\cite{hamilton1,hamilton2}, which covers some of the same material, and provides a more detailed review of some of the statistical methods that are used. Additionally it is worth directing the interested reader to a number of good text books that cover this topic \\cite{coles_lucchin,dodelson,liddle_lyth,book_of_john}. In addition to a description of the basic mathematics and physics behind a clustering analysis I have attempted to provide a discussion of some of the fundamental and practical difficulties involved. The cosmological goal of such an analysis is consider in the final part of this manuscript, where the combination of cosmological constraints from galaxy clustering and the CMB is discussed, and an example multi-parameter fit to recent data is considered. ",
        "conclusions": ""
    },
    "0601/astro-ph0601681_arXiv.txt": {
        "abstract": "{136 stars which were known to be the members of open cluster NGC 752 were observed at $R$ band with ROTSE--IIId telescope located at the Turkish National Observatory (TUG) site. The data had been evaluated together with BV and 2MASS photometric data. A new practical method for separating dwarf and giant was described and applied. Evaluating the colour magnitude--diagrams with Padova isochrones revealed metallicity similar to the Sun and an age of 1.41 Gyr for the open cluster NGC 752. ",
        "introduction": "One of the problems of the Galactic astronomy is the estimation of Galactic model parameters of giant stars in our Galaxy. In many studies, the Galactic model parameters are estimated without any discrimination between dwarfs and giants, whereas some researchers estimated model parameters only for certain star categories (e.g. Pritchet 1983, Bahcall \\& Soneira 1984, Buser \\& Kaeser 1985 and Mendez \\& Altena 1996). A very recent work is an example for this where the Galactic model parameters were estimated using only giants (Cabrera-Lavers, Garzon \\& Hammersley 2005). The separation of field dwarfs and field giants plays an important role for such kinds of works. The most efficient classical methods of identifying dwarf and giant stars utilize spectroscopy. By inspecting spectral line profiles, one has to estimate surface gravity to discriminate between higher and lower pressure stellar atmospheres to be sure for the identification. This is, however time consuming and tiring. A rather easier procedure is to separate dwarfs and evolved stars (subgiants or giants) such as to obtain a luminosity function consistent with the local luminosity function of nearby stars due to Gliese \\& Jahreiss (1991) and Jahreiss \\& Wielen (1997). The procedure of this separation is based on the fact that the local luminosity functions obtained for many fields indicates a systematic excess of star counts relative to the luminosity function of nearby stars for the fainter segment, i.e. $M(V)\\geq5^{m}.5$, and a deficit for brighter segment, $M(V)<5^{m}.5$. (in $RGU$ system $M(G)\\geq6^{m}$ and $M(G)<6^{m}$, respectively). The works of Karaali (1992); Ak, Karaali \\& Buser (1998); Karata\\c{s}, Bilir \\& Karaali (2000); Karaali et al. (2000); Karata\\c{s}, Karaali \\& Buser (2001); Karaali, Bilir \\& Buser (2004); Bilir, Karaali \\& Buser (2004) and Karata\\c{s} et al. (2004) can be given as examples for application of this procedure.  Recently, a new method were suggested by Bilir et al. (2006) for separating the field dwarfs and field giants. This new method is based on the comparison of the Two Micron All Sky Survey (2MASS, hereafter) $J$, $H$, $K_{s}$ with the $V$ magnitudes down to the limiting magnitude of $V=16$. In this work we extend application of this method to the open cluster NGC 752 observed by a robotic telescope ROTSE-IIId (Akerlof et al. 2003).  This paper is organized as follows. In Section 2 the BV, 2MASS and ROTSE data are presented. In Section 3 the method is applied to ROTSE and 2MASS data, and the separation of dwarf and giant stars is tested. In Section 4 colour-magnitude diagrams (CMDs) of NGC 752 is compared to the Padova isochrones. Finally, the conclusion is given in Section 5.  \\begin{figure} \\center \\resizebox{5cm}{7cm}{\\includegraphics*{fig01.eps}} \\caption {$V \\times (B-V)$ CMD of 136 probable member stars in NGC 752.} \\end {figure}  \\begin{figure*} \\center \\resizebox{17cm}{5.51cm}{\\includegraphics*{fig02.eps}} \\caption {V and 2MASS magnitudes of 136 stars in NGC 752. (a) $J_{0} \\times V_{0}$, (b) $H_{0} \\times V_{0}$, and (c) $Ks_{0} \\times V_{0}$. The solid lines in diagrams were drawn according to eqs. (1), (2), and (3).} \\end {figure*} ",
        "conclusions": "In this study, we have reduced the relations between the 2MASS and $V$ magnitudes by which field dwarfs and giants can be separated, to the USNO A2.0 $R$-band magnitudes. The $R_{0}$ magnitudes of 136 stars in open cluster NGC 752 were transferred to the $V_{0}$ magnitudes, and the relations between $J_{0}$, $H_{0}$, $(K_{s})_{0}$ and $R_{0}$ were derived by means of the relations between $J_{0}$, $H_{0}$, $(K_{s})_{0}$ and $V_{0}$ given by Bilir et al. (2006). Dwarf and giant stars identified by the CMD of open cluster NGC 752 lie at different sides of the line representing the relation between the 2MASS and $R_{0}$ magnitudes, in the $J_{0} \\times R_{0}$, $H_{0} \\times R_{0}$ and $(K_{s})_{0} \\times R_{0}$ diagrams. The best one is the last diagram, i.e. $(K_{s})_{0} \\times R_{0}$. Thus, dwarf-giant separation could be carried out also in the ROTSE-IIId data. Proven to be successful, this practical method can provide good contributions to the studies of Galactic model parameters in which separation of dwarfs and giants were needed.  A set of Padova isochrones were fitted to the CMDs of the open cluster NGC 752 using BV and 2MASS photometric data. It turned out that the isochrone with chemical composition Z=0.019 and Y=0.273 which reveals an age of 1.41 Gyr for the open cluster NGC 752 could be fitted to all segments, i.e. main-sequence, turn-off and giant branch, of the $M_{J} \\times (J-K_{s})_{0}$ two-colour diagram. This result is very close to the age 1.24$\\pm$0.20 Gyr which Salaris, Weiss \\& Percival (2004) calculated from the morphology of 71 open clusters in our Galaxy. The metal-abundance of the cluster given by Daniel et al. (1994), i.e. $[Fe/H]=-0.15\\pm0.05$ dex, is a strong confirmation for our result."
    },
    "0601/astro-ph0601362_arXiv.txt": {
        "abstract": "{  The quasar B\\,2005+403 located behind the Cygnus region in our Galaxy, is a suitable object for studying the interplay between propagation effects, which are extrinsic to the source and source intrinsic variability. On the basis of VLBI experiments performed at 1.6, 5, 8, 15, 22, and 43\\,GHz during the years 1992 -- 2003 and parallel multi-frequency monitoring of the total flux density (between 5 -- 37\\,GHz), we investigated the variability of total flux density and source structure. Below 8\\,GHz, the point-like VLBI source is affected by scatter-broadening of the turbulent interstellar medium (ISM), which is located along the line of sight and likely associated with the Cygnus region. We present and discuss the measured frequency dependence of the source size, which shows a power-law with slope of $-1.91 \\pm 0.05$. From the measured scattering angle at 1\\,GHz of $\\theta = 77.1 \\pm 4.0$\\,mas a scattering measure of SM= $0.43 \\pm 0.04 {\\rm\\,m}^{-20/3}$\\,kpc is derived, consistent with the general properties of the ISM in this direction. The decreasing effect of angular broadening towards higher frequencies allows to study the internal structure of the source. Above 8\\,GHz new VLBI observations reveal a one-sided slightly south-bending core-jet structure, with stationary and apparent superluminally moving jet components. The observed velocities range from $6.3 - 16.8$\\,c. The jet components move on non-ballistic trajectories. In AGN, total flux density variations are often related to the emergence of new VLBI components. However, during almost eleven years no new component was ejected in B\\,2005+403. A striking feature in the flux density variability is a trough observed at 5 -- 37\\,GHz and between 1996 and 2001. This trough is more pronounced at higher frequencies, where it shows a prominent flux density decrease of $\\sim 50 - 60$\\,\\% on a time scale of $2-3$\\,yrs. The trough can be explained as a blending effect of decreasing and increasing jet component fluxes. Dense in time sampled flux density monitoring observations with the 100\\,m Effelsberg telescope reval intra-day variability (IDV) at 1.6\\,GHz with a modulation index of $m=1.0$\\,\\%. At 5\\,GHz less pronounced variations are seen ($m=0.5$\\,\\%). This and a relatively short variability time scale of $\\sim 0.1$\\,days imply a second, less dense or turbulent scattering screen at nearby (few to hundred parsec) distance. ",
        "introduction": "The propagation of radio waves through the ionized interstellar medium causes several effects, such as Faraday rotation and depolarization of polarized emission, dispersion of pulsar signals, scatter broadening of compact radio sources and intensity fluctuations caused by diffractive (DISS) and refractive (RISS) interstellar scintillation (ISS). The effect of scatter broadening plays an important role in the understanding of the compact object located in the center of our galaxy \\object{Sgr\\,A*} \\citep[e.g.][]{galcenter2,galcenter3,galcenter}. It is however also important for other galactic regions on the sky, e.g. for the Cygnus region and the quasar B\\,2005+403 located behind. The detailed study of scatter broadening in compact radio sources could lead to a better understanding of the interstellar medium. It also bears the potential to reveal fine details of the intrinsic source structure, particularly at longer wavelengths, where direct imaging often cannot provide sufficiently high angular resolution.  The other prominent propagation effect addressed in this paper is the so called Intra-Day Variability \\citep[IDV,][]{idv_discovery,idvh}, and the question if and how it is related to the interstellar scintillation. Recent IDV-surveys show that a large fraction (up to 30\\,\\%) of all compact flat spectrum radio sources show this effect, which is characterized by variability amplitudes of up to $20-30$\\,\\% and variability timescales ranging from less than one hour to several days \\citep[e.g.][]{reduc_idv,idvh,southern,lovell}. If interpreted via source intrinsic incoherent emission processes, such short variability timescales imply -- via the light travel time argument --  apparent brightness temperatures of $T_\\text{B} = 10^{16...19}$\\,K, far in excess of the inverse-Compton limit \\citep{compton}. With the assumption of relativistic Doppler-boosting the brightness temperatures can be reduced, which however, would imply uncomfortably large Doppler boosting factors of $\\delta \\simeq 20 - 200$.  Alternately, the IDV in the radio-bands is also explained extrinsically via scintillation of radio waves in the turbulent interstellar medium (ISM) of our Galaxy \\citep[e.g.][]{rick_uj,dt}. The main problem in this interpretation is that it lacks the explanation of observed radio-optical broad-band correlations in at least some sources (\\cite{corr_var}, see also \\cite{ww}). The recent detection of diffractive ISS in the source J\\,1819+3845 leads to micro-arcsecond source sizes and brightness temperatures of $\\sim 10^{14}$\\,K, which again requires Doppler boosting factors of $\\delta \\simeq 100$ \\citep{diff_1819}. It is therefore likely that the IDV phenomenon invokes both, a combination of source intrinsic {\\it and} extrinsic effects \\citep[e.g.][]{krich,int_extr}.   In this paper we present new measurements for the scatter broadened quasar B\\,2005+403, combining Very Long Baseline Interferometry  (VLBI) and flux density variability measurements obtained at various frequencies over the last decade. \\object{B\\,2005+403} is a flat spectrum quasar \\citep[$S_{\\rm 5\\,GHz} = 2.5 - 3.5$\\,Jy, $\\alpha_{\\rm 0.3/5\\,GHz} = 0.3$, with $S \\sim \\nu^{\\alpha}$;][]{spekt} at a redshift of $z=1.736$ \\citep{redshift}. It is located close to the Galactic plane at $l$\\,=\\,$76.82^\\circ$, $b$\\,=\\,$4.29^\\circ$ behind the \\object{Cygnus super-bubble} region. Earlier studies showed that interstellar scattering affects the VLBI image of the source, causing angular broadening at frequencies below 5\\,GHz \\citep{fey,mutel,desai}. The high value of the scattering measure ($ \\text{SM}$\\,=\\,$\\int_0^L C^{2}_{N}\\left( s \\right) ds $ \\footnote{The Scattering Measure is the path integral over the coefficient $C^{2}_{N}$ of the electron density fluctuation wavenumber spectrum.}), derived for the line of sight of B\\,2005+403 by \\cite{fey}, reflects the strong influence of interstellar medium.  In order to further investigate the scatter broadening and the amount of scintillation in B\\,2005+403, we combined all available flux density monitoring data with our VLBI observations. The wide frequency coverage of our data facilitates a study of the interplay of the frequency dependent scattering effects, which dominates below 8\\,GHz and the source intrinsic variability, which is best seen at and above 15\\,GHz. At these higher frequencies B\\,2005+403 has not been intensively studied with VLBI before. The VLBI data cover a time range of 11 yrs (from 1992 -- 2003) and allow us to study and monitor the structural evolution of VLBI-core and pc-scale jet. The scatter broadening observed in the VLBI images at the lower frequencies is measured and gives further constraints to the properties of the intervening ISM. Additional parameters for the ISM are obtained from the IDV monitoring of the source, performed with the Effelsberg 100 m telescope.  The paper is organized as follows. In Sect. 2, we present details of the VLBI observations and corresponding data reduction and introduce the long-term trends of the flux density variability. In Sect. 3 we discuss the source structure, the kinematics of the inner jet components and the possible location of the nucleus (VLBI core). In Sect. 3.4, we focus on the results of the low frequency VLBI observations and discuss the scatter broadening. In section 3.5 we discuss the long-term variability of B\\,2005+403 and relate it to changes of the VLBI structure. In Sect. 3.6 we discuss the implications from the observed intraday variability. Sect. 4 gives a concluding summary.  The following cosmological parameters are used throughout this paper: \\(H_{0}\\)\\,=\\,\\(71\\ \\text{km\\ s$^{-1}$\\ Mpc$^{-1}$,} \\ \\Omega_{\\text{m}}\\)\\,=\\,\\(0.27 \\ \\text{and} \\ \\Omega_{\\text{vac}}\\)\\,=\\,\\(0.73\\). ",
        "conclusions": "B\\,2005+403 is located in the Cygnus-region and previously was known as one of the most scattered broadened extragalactic sources in this region of the sky. In this paper we combined multi-frequency flux density monitoring and VLBI imaging observations. We also used high time resolution flux density monitoring observations, in order to search for possible Intra-Day Variability. The combination of these complementary data yields new and more accurate than previously published estimates of (i) the source structure and its kinematics and of (ii) the properties of the interstellar medium, which is responsible for the observed propagation effects.  High frequency VLBI imaging observations (at 15 - 43 GHz) spanning 11 years (1992 - 2003) reveal a one sided and bent core jet structure, with at least 5 embedded VLBI components. The inner jet components separate from the stationary assumed core C1 with apparent superluminal speeds of 6 - 17 c. The component nearest to the core (C2), however, appeared stationary ($\\beta_\\text{app} \\leq 0.3$). It is remarkable that the flux density variation of this stationary and steep spectrum component apparently violates the inverse Compton limit, indicating strong Doppler-boosting. This could be explained with a component path, which is oriented at a smaller angle to the line of sight than those of the other jet components. A systematic increase of the component speeds with increasing core separation is observed. The paths of the inner-most jet components is curved and their motion is not ballistical, suggesting a spatially curved (helical) jet. Beyond 2\\,mas core-separation, however, the components seem to move on linear (ballistical) trajectories. Continued future VLBI monitoring will be necessary to determine their paths more accurately.  The VLBI images at longer wavelengths (1.6 - 8 GHz) show a scatter broadened nearly point-like source. Combining our new size measurements with published data from earlier observations allowed to study the frequency dependence of the source size in greater detail. Below 8 GHz this dependence is best described by a non-Kolmogorov power law with slope of -1.9. For the size of the scattering disk at 1\\,GHz we obtain $\\left( 77.1 \\pm 4.0 \\right) $\\,mas. Above 8\\,GHz the measured sizes are increasingly larger than the prediction from the scattering law, and the internal source structure becomes visible. Based on the observed scatter broadening, several parameters of the scattering medium (scattering size, scattering measure, electron density) were determined and found to be in accordance with previously published estimates (NE2001 model, Cordes \\& Lazio, 2002). The NE2001 model places the scattering screen at a distance of 2.35\\,kpc. We however note that towards the direction of  B\\,2005+403 there is a degree-size extended supernova remnant \\object{G78.2+2.1}. It exhibits a patchy, inhomogeneous structure at optical, X-ray and radio wavelengths \\citep[][and references therein]{snr_xray}. Its distance is $ \\left( 1.5\\pm0.5 \\right)$ \\,kpc \\citep{snr_dist}, still compatible with the scattering model.  We discuss the observed long term flux density variability curves (5 - 37\\,GHz), which during 1996 - 2001 showed a remarkable flux density 'trough'. A tentative interpretation via a possible 'scattering event', similar to the 'extreme scattering events' observed in other sources, could be abandoned on the basis of the observed frequency dependence of the variability, which appeared more pronounced at higher frequencies. Instead, we show that the 'trough' is well explained by the summed flux density variability of the (independently) varying VLBI components (C1 - C3). In contrast to other sources, which often show correlations between jet component ejection and flux density variability \\citep[e.g.][and references therein]{var}, such a relation is not seen B\\,2005+403. In this source at least some (if not most) of the observed flux density variability results from the blending of the evolving jet components.  Dense in time sampled variability measurements performed with the 100\\,m Effelsberg telescope showed only weak intra-day variability in B\\,2005+403. This was in contrast to naive expectations, in which more pronounced variability was anticipated. The low variability indices ($m=1$\\,\\% at 18\\,cm, $m=0.5$\\,\\% at 6\\,cm) and the short variability time scale of $\\sim 0.1$\\,days cannot be explained by scattering effects from the kpc-screen, which is responsible for the scatter broadening. Instead, another and much closer scatterer is required. Using the thin screen approximation for refractive interstellar scintillation we determined the likely distance of the scattering screen to be $(6 \\leq L \\leq 128)$\\,pc. At 18\\,cm, its scattering size would lie in the range of  $0.5 \\leq \\theta_\\text{scat} \\leq 6.3$\\,mas. For B\\,2005+403 one therefore has to deal with the situation that a scattered broadened image (kpc-screen) is scattered a second time (pc-screen) and that the combination of both effects limits the accuracy of low frequency flux density variability measurements to about 0.5-1\\,\\% and the angular the resolution of VLBI observations to a few milliarcseconds. We note that weak refractive interstellar scintillation from a very nearby and ionized ISM (which most likely covers larger portions of the sky), may therefore also limit the calibration accuracy and sensitivity of planned large radio telescopes operating and low frequencies, like LOFAR or the square-kilometer array (SKA)."
    },
    "0601/astro-ph0601154_arXiv.txt": {
        "abstract": "We examine the stability of $\\ell=1$ and $\\ell=2$ g modes in a pair of nitrogen-rich Wolf-Rayet stellar models characterized by differing hydrogen abundances. We find that modes with intermediate radial orders are destabilized by a $\\kappa$ mechanism operating on an opacity bump at an envelope temperature $\\log T \\approx 6.25$. This `deep opacity bump' is due primarily to L-shell bound-free transitions of iron. Periods of the unstable modes span $\\sim 11-21\\,\\hr$ in the model containing some hydrogen, and $\\sim 3-12\\,\\hr$ in the hydrogen-depleted model. Based on the latter finding, we suggest that self-excited g modes may be the source of the 9.8\\,\\hr-periodic variation of WR~123 recently reported by \\citet{Lef2005}. ",
        "introduction": "\\label{sec:intro}  Of all of the differing excitation mechanisms responsible for free oscillations of stars \\citep[for a comprehensive review, see][and references therein]{Unn1989}, none is more ubiquitous than the opacity, or `$\\kappa$', mechanism. Put in simple terms, this mechanism works in layers of a star's envelope where the opacity $\\kappa$ increases in response to compression of the stellar material. Such an increase (initially produced for instance by stochastic perturbations to the equilibrium state) retards the escape of energy from the stellar core, with the trapped flux being deposited in the layer as heat. As \\citet{Edd1926} discusses, the supply of additional heat during compression creates a Carnot-cycle heat engine capable of converting part of the outflowing radiant energy into mechanical energy. In tandem with a restoring force that returns the layer back toward its equilibrium state (pressure for p modes, buoyancy for g modes), the heat engine creates a condition of pulsational instability (overstability).  The role played by the $\\kappa$ mechanism in the classical ($\\delta$) Cephei pulsators was first identified by \\citet{Zhe1953}. In these stars, the source of the instability is a peak in the opacity and its partial derivative $\\kappaT \\equiv (\\partial \\ln \\kappa/\\partial \\ln T)_{\\rho}$ at an envelope temperature $\\log T \\approx 4.65$ where helium undergoes its second ionization. At lower luminosities, the same mechanism is primarily responsible for the RR Lyrae and $\\delta$ Scuti classes of pulsating star \\citep[see][and references therein]{Cox1980}.  However, in the early-type $\\beta$ Cephei stars and slowly pulsating (SPB) stars, the helium opacity bump is situated too close to the stellar surface to play any significant role in destabilizing pulsation modes.  In fact, the excitation mechanism responsible for these two classes of star remained a mystery until the advent of new opacities from the OPAL and OP projects \\citep{RogIgl1992,Sea1994}. The treatment of spin-orbit splitting in same M-shell transitions of iron and nickel, absent in previous calculations, introduced a new opacity peak at $\\log T \\approx 5.3$. Stability analyses by \\citet{Cox1992}, \\citet{DziPam1993}, and \\cite{Dzi1993} revealed that this so-called `iron bump' can destabilize p modes in $\\beta$ Cephei stars and g modes in SPB stars. More-recent investigations by \\citet{Cha1997} and \\citet{Fon2003} have inferred that the same iron bump can also drive the pulsation of the short-period (EC14026) and long-period subdwarf B (sdB) stars.  In parallel with these developments, the new opacity data also impacted understanding of the pre-white dwarf GW Vir (pulsating PG1159) stars. \\citet{Sta1983,Sta1984} found that g modes in these objects could be $\\kappa$-mechanism excited by an opacity bump around $\\log T \\approx 6.2$, arising from K-shell photoionization of carbon and oxygen; however, an almost-pure C/O mixture was required for the instability to operate \\citep{Sta1991}. This abundance restriction is relaxed when updated OPAL/OP date are adopted \\citep[see][and references therein]{Gau2005}, due to the appearance of extra opacity around $\\log T \\approx 6.3$. The origin of this additional opacity appears to be spin-orbit effects in L-shell bound-free transitions of iron \\citep{RogIgl1992}.  From this brief survey of the literature, a natural question emerges. The iron bump in the OPAL/OP opacities leads to instability strips at both high masses ($\\beta$ Cephei and SPB stars) and low masses (sdB stars). Given that opacity bump at $\\log T \\sim 6-6.3$ drives pulsation in the low-mass GW Vir stars, is there by analogy a high-mass group of stars that also exhibit pulsation driven by this `deep opacity bump' (DOB)? In order for the pulsation to be remain non-adiabatic around the driving zone, and therefore allow the $\\kappa$-mechanism heat engine to operate, these putative stars must be exceedingly hot, with surface temperatures approaching $10^{5}\\,{\\rm K}$. In fact, these stars can be recognized as massive, population I, core helium burning Wolf-Rayet stars.  In this letter, we conduct a preliminary stability analysis of a pair of model Wolf-Rayet (W-R) stars, to evaluate the hypothesis that g modes can be excited in these stars by the DOB. We introduce the models in Sec.~\\ref{sec:models}, and describe the stability analysis in Sec.~\\ref{sec:stability}. Our results are presented in Sec.~\\ref{sec:results}, and then discussed and summarized in Sec.~\\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss}  In the preceding section, we demonstrate that low-degree, intermediate radial-order g modes are unstable in a pair of stellar models representative of WN stars, due to a $\\kappa$ mechanism operating on the deep opacity bump. This finding supports the central hypothesis of the paper advanced in Sec.~\\ref{sec:intro}, and likewise confirms WN stars as the high-mass cousins of GW Vir stars, exhibiting a relationship parallel to that between the high-mass $\\beta$ Cephei and SPB stars, and the low-mass sdB pulsators.  However, there are caveats to our results. Beyond the various (reasonable) simplifications we adopt in constructing our stellar models (Sec.~\\ref{sec:models}) and conducting the stability analysis (Sec.~\\ref{sec:stability}), we have neglected the effects of the strong, radiatively driven winds that are characteristic of all W-R stars. The stellar models incorporate only the evolutionary effects of wind-originated mass loss, and the stability analysis ignores the radial outflow at the outer boundary. As discussed by \\citet{Cra1996}, an outflow can allow ordinarily-trapped pulsation modes to propagate outward through the boundary. \\citet{Tow2000} demonstrated that this leakage of wave energy can result in an appreciable damping of pulsations. If the unstable g modes found herein are to remain unstable, the rate of energy input from the $\\kappa$ mechanism must exceed the damping due to leakage.  A thorough evaluation of this issue may have to await further theoretical advances in understanding how outflows influence pulsation. In the meantime, we discuss an observational development of particular relevance to the present work, one that poses some interesting questions regarding the role played by pulsation (whether leaking or not) in establishing structure and modulation in W-R winds. Although searches for \\emph{periodic} light variations in WN8 stars -- the subtype most characterized by variability -- have proven largely unsuccessful \\citep[see, e.g.,][and references therein]{Mar1998}, recent \\emph{MOST} observations of the WN8 star HD~177230 (WR~123) by \\citet{Lef2005} have detected an unambiguous 9.8\\,\\hr-periodic signal in the star's light curve.  Pulsation appears a promising candidate for the origin of the star's variation. W-R stars have been known for some time to be unstable toward strange modes \\citep[e.g.,][]{Gla1993}; however, there is debate over whether such modes are compatible with the observations of WR~123. On the one hand, \\citet{Lef2005} themselves argue that the periods typical to strange modes, being on the order of the dynamical timescale \\tdyn, are far too short to match the 9.8\\,\\hr\\ period detected by \\emph{MOST}. On the other hand, \\citet{Dor2006} reason that as a WN8 star WR~123 is expected to have a significantly larger radius, and hence longer strange-mode periods, than the helium-star models on which \\citet{Lef2005} based their conclusions. In support of their argument, \\citet{Dor2006} present a WN8 model having $\\Rstar=15.4\\Rsun$ that exhibits unstable radial strange modes with periods on the order of a half day, seemingly compatible with the observations.  \\begin{figure} \\leavevmode \\begin{centering} \\epsffile{figure3.eps} \\caption{Differential work functions, plotted as a function of envelope temperature, for the $\\ell=1$ modes of the two stellar models having the largest growth rates. The differential is with respect to the dimensionless radius $x\\equiv r/\\Rstar$, and the work $W$ is expressed in units of the total energy of pulsation \\Epul.} \\label{fig:work} \\end{centering} \\end{figure}  \\begin{table} \\leavevmode \\begin{centering} \\caption{Properties of the unstable $\\ell=1$ and $\\ell=2$ g modes of the ``WNL'' and ``WNE'' stellar models. The subscripts $_{\\rm min}$ and $_{\\rm max}$ indicate the minimum and maximum values of the relevant quantities, respectively.} \\label{tab:growth} \\begin{tabular}{@{}ccccccc} model & $\\ell$ & $P_{\\rm min}$/\\hr & $P_{\\rm max}/\\hr$ & $\\ncowl_{\\rm min}$ & $\\ncowl_{\\rm max}$ & $\\eta_{\\rm max}$ \\\\ \\hline ``WNL'' & 1 & 10.8 & 20.6 & -20 & -10 & $6.67 \\times 10^{-5}$ \\\\ ``WNL'' & 2 & \\multicolumn{5}{c}{\\emph{No unstable modes}} \\\\ ``WNE'' & 1 & 3.83 & 12.4 & -14 & -4 & $1.67 \\times 10^{-3}$ \\\\ ``WNE'' & 2 & 2.81 & 8.34 & -15 & -5 & $1.12 \\times 10^{-3}$ \\end{tabular} \\end{centering} \\end{table}  Here, there may be a problem. The large radii of the models studied by \\citet{Dor2006} are, we believe, linked to the authors' \\emph{ab initio} assumption of an envelope hydrogen abundance $X = 0.35$. Although representative of many WN8 stars, this value is inappropriate to WR~123, which is well-known to be largely depleted of hydrogen \\citep[$X \\lesssim 0.005$; see, e.g.,][]{Cro1995}. It seems probable that the absence of hydrogen in the star is a consequence of it having already shed any residual hydrogen-rich envelope via wind mass loss. Accompanying the loss of this envelope should be a substantial reduction in the star's radius, as can be seen by comparing our ``WNL'' and ``WNE'' models in Table~\\ref{tab:models}.  Based on this reasoning, hydrogen-poor WR~123 may be characterized by an appreciably smaller radius than assumed by \\citet{Dor2006}, resulting once again in the period mismatch noted originally by \\citet{Lef2005}. This leads us to advance g modes excited by the DOB as an alternative explanation for the star's variability. As Fig.~\\ref{fig:growth} illustrates, the 9.8\\,\\hr\\ period is covered by the unstable $\\ell=1$ g modes of the hydrogen-poor ``WNE'' model. While we describe this model as `early-type', we once again emphasize that the label is to be understood in an evolutionary rather than spectroscopic sense.  The detectable presence of g modes in WR~123 could offer unparalleled insights into the star's internal structure; these modes penetrate down to the boundary of the convective core, and are therefore particularly well suited to asteroseismological analyses. However, it is unclear how best to proceed with the necessary initial step of testing our interpretation of the observations. Pulsation in g modes is often diagnosed through time-series monitoring of spectroscopic line profile variations (lpv); the predominantly horizontal velocity fields generated by these modes \\citep[e.g.,][]{Unn1989} tend to concentrate variability in the wings of line profiles. Unfortunately, this approach is not feasible in the case of W-R stars, because the supersonic wind washes out any lpv arising from (subsonic) pulsational velocity fields. It remains to be seen whether other approaches can be devised that are able to discriminate g modes from other potential sources of variability.  Beyond the specific case of WR~123, the discovery of a new class of Wolf-Rayet instability is of great interest in itself. Future investigation can now focus on exploring the characteristics and extent of the DOB instability in greater detail. Specific topics worthy of attention include establishing red and blue edges in the HR diagram, and determining the sensitivity of the instability against changes in input physics (e.g., elemental diffusion, mass-loss rates, etc.). As we have already noted, a proper treatment of pulsation at the outflowing outer boundary is also needed. An integral part of this treatment must be a better understanding of the reverse side of the pulsation-wind interaction: how can pulsation at the wind base lead to the modulations observed in Wolf-Rayet stars? Initial steps toward addressing this question were made by \\citet{CraOwo1996}, but there remains much work to be done."
    },
    "0601/astro-ph0601288_arXiv.txt": {
        "abstract": "{Evidence has been mounting recently that activity in some radio-loud AGNs (RLAGNs) can cease shortly after ignition and that perhaps even a majority of very compact sources may be short-lived phenomena because of a lack of stable fuelling from the black hole. Thus, they can fade out before having evolved to large, extended objects. Re-ignition of the activity in such objects is not ruled out.} {With the aim of finding more examples of these objects and to investigate if they could be RLAGNs switched off at very early stages of their evolution, multifrequency VLBA observations of six sources with angular sizes significantly less than an arcsecond, yet having steep spectra, have been made.} {Observations were initially made at 1.65\\,GHz using the VLBA with the inclusion of Effelsberg telescope. The sources were then re-observed with the VLBA at 5, 8.4 and 15.4\\,GHz. All the observations were carried out in a snapshot mode with phase referencing.} {One of the sources studied, 0809+404, is dominated by a compact component but also has diffuse, arcsecond-scale emission visible in VLA images. The VLBI observations of the ``core'' structure have revealed that this is also diffuse and fading away at higher frequencies. Thus, the inner component of 0809+404 could be a compact fading object. The remaining five sources presented here show either core-jet or edge-brightened double-lobed structures indicating that they are in an active phase.} {The above result is an indication that the activity of the host galaxy of 0809+404 may be intermittent. Previous observations obtained from the literature and those presented here indicate that activity had ceased once in the past, then restarted, and has recently switched off again.} ",
        "introduction": "The activity period for radio-loud AGNs (RLAGN) can last up to $\\sim$$10^{8}$ years \\citep{al87,liu92} and, as their lobes are huge reservoirs of energy, even if the energy supply from the central engine to the hotspots and the lobes eventually cuts off, the radio sources are still observable for a substantial period of time. This so-called ``coasting phase'' of the lobes of a RLAGN can last up to $10^{8}$\\,yr \\citep{kom94,sl01} and preserves information of past nuclear activity. As the source gradually fades out, its spectrum becomes steeper and steeper because of radiation and expansion losses. Objects possessing these features are sometimes termed ``relics'' or ``faders''.  The structure and other properties of the double radio source B2\\,0924+30, identified with an E/S0 galaxy IC\\,2476, shows it to be a good example of a fader \\citep{cor87}. The projected linear size of the whole system is 270\\,$h^{-1}$\\,kpc\\footnote{For consistency with earlier papers in this field, the following cosmological parameters have been adopted throughout this paper: $H_0$=100${\\rm\\,km\\,s^{-1}\\,Mpc^{-1}}$ and $q_0$=0.5. Wherever in the text the linear sizes are referred to, $h^{-1}$ is introduced.} which indicates that it is a Large Symmetric Object (LSO). \\citet{jam04} have confirmed that B2\\,0924+30 is indeed a relic radio structure and that it switched off its activity $\\sim$$5\\times10^7$ years ago and, as such, can be labelled a ``dead'' radio galaxy. A large number of examples of fossil radio galaxies or cluster relic systems have recently been found by \\citet{coh04} as a result of their VLA observations at 74\\,MHz.  There are no obvious reasons for the existence of a lower limit to the length of the activity period in a RLAGN; it could be shorter than has already been seen for LSOs. If this is the case it could be that the growth of the radio source has been impeded, even at an early stage of its evolution. As a result, small-scale faders might exist. An attempt to test observationally if there are young faders among Medium-sized Symmetric Objects (MSOs), i.e. the objects that have linear sizes in the range 1--20 $h^{-1}$\\,kpc, has been made by \\citet{kmts05} -- hereafter Paper~II. This revealed one strong candidate --~1542+323.  \\citet{gir05} have described a class of low power compact (LPC) radio sources, their small sizes and moderate luminosities (comparable to those of low power giant FR\\,I radio galaxies) being ascribed to a number of different physical reasons: youth, low kinetic power of the jets or frustration and also the premature end of nuclear activity. 1855+37, one of the sources they investigated, could be a very good example of a compact fader.  There is no reason why one should not search for faders among the most compact of radio sources, particularly those belonging to the class of Compact Symmetric Objects (CSOs), i.e. sub-kiloparsec-scale extragalactic radio sources with symmetric radio structures. Most of these sources are triples with the central component being a radio core, or doubles with only two detectable radio lobes. Some CSOs have radio spectra that peak at a few gigahertz and have been classified as Gigahertz-Peaked Spectrum (GPS) sources. GPS radio galaxies seem to be CSOs with a simple structure \\citep{odea98}. However, most GPS quasars with a core-jet or complex milliarcsecond-scale morphology do not seem to have symmetric structures and they are thought to be a different class of object \\citep{stan03}.  The physical origin of CSOs has been explained in two ways. According to the frustration scenario \\citep{bmh84} the radio source is permanently confined to a region within the host galaxy by a dense environment from which it is impossible for it to evolve into a large source. \\citet{alex00} also claims that the observational density of sources in the power--linear size plane can be reproduced only if it is assumed that there exists a class of young frustrated objects. They would remain very weak and have diffuse hotspots and a brighter jet. Those objects, which manage to escape the high-density regions, would continue their evolution and reach the classical FR\\,II stage.  Alternatively, \\citet{pm82} and \\citet{c85} suggested that CSOs -- they were labelled compact-doubles at that time; the term ``CSO'' was introduced later by \\citet{wil94} and \\citet{r94} -- could be young radio sources that would evolve into large radio objects during their lifetimes. Based upon this scenario, \\citet{r96} proposed an evolutionary scheme unifying three classes of radio sources: CSOs, MSOs -- a subset of the Compact Steep Spectrum sources (CSS) class -- and LSOs.  At the present time this is the youth scenario that is generally accepted; see the review by \\citet{fanti00}. An argument in favour of this has been found mainly in age measurements of individual classes of radio sources: CSOs are younger than $\\sim$$10^4$~years \\citep{ocp98,gir03}, MSOs are typically $\\sim$$10^5$~years old \\citep{mur99} and LSOs can manifest their activity for up to $\\sim$$10^8$~years \\citep{al87,liu92}. \\citet{sn99,sn00} added an important ingredient to this scheme, namely that the radio luminosities of CSOs increase as they evolve, reach a maximum in the MSO phase and then gradually decrease as these objects increase in size to become LSOs.  \\citet{mar46,mks06} claim that this evolutionary track is not the only one possible. In fact, several, if not a ``continuum'' of such tracks might exist and the one shown by \\citet{sn99,sn00} just {\\em appears} as the only one simply because of selection effects. If the energy supply cuts off early, the object leaves the ``main sequence'' proposed by \\citet{sn99,sn00} and will never reach the LSO stage, at least in a given phase of activity. Thus, there should exist a class of (very) small-scale objects that resemble large-scale faders. This conforms to early predictions that many CSOs could be young objects that switch off after a short period of time \\citep{r94}. Strong support for such an idea also comes from \\citet{rb97} who proposed a model in which extragalactic radio sources are intermittent on timescales of $\\sim$$10^4$--$10^5$~years. The above considerations have recently been supported by observational results obtained by \\citet{gug05}, suggesting that many CSOs die young or are episodic in nature, and so it is likely that only a minority of them ``survive'' and further evolve.  In this paper -- the fourth of the series -- VLBA observations of 6~compact sources which are candidates for prematurely dying CSOs are presented and discussed.  \\begin{table*}[t] \\caption[]{Optical magnitudes and radio flux densities of target sources at two frequencies.} \\begin{center} \\begin{tabular}{@{}c c c c l l c l c l r l@{}} \\hline \\hline &    &     &   &     &   & \\multicolumn{1}{c}{Total} & & \\multicolumn{1}{c}{Total} &     &  &  \\\\ ~~~Source & RA & Dec & ID& \\multicolumn{1}{c}{$m_{R}$}& \\multicolumn{1}{c}{\\it z}& \\multicolumn{1}{c}{flux at}& \\multicolumn{1}{c}{log$P_{1.4\\mathrm{GHz}}$}& \\multicolumn{1}{c}{flux at}& \\multicolumn{1}{c}{$\\alpha_{1.4\\mathrm{GHz}}^{4.85\\mathrm{GHz}}$}& \\multicolumn{1}{c}{LAS}& \\multicolumn{1}{c}{LLS}\\\\ ~~~Name   & h~m~s & $\\degr$~$\\arcmin$~$\\arcsec$ &  & & &1.4\\,GHz& &4.85\\,GHz&&&\\\\ &    &     &   &     &   & \\multicolumn{1}{c}{[mJy]}& \\multicolumn{1}{c}{[W~Hz$^{-1}$]} & \\multicolumn{1}{c}{[mJy]}& & \\multicolumn{1}{c}{[mas]}& [$h^{-1}~{\\rm pc}]~~~$ \\\\ ~~~(1)& (2)& (3) &(4)& \\multicolumn{1}{c}{(5)}& \\multicolumn{1}{c}{(6)}& \\multicolumn{1}{c}{(7)}& \\multicolumn{1}{c}{(8)}& \\multicolumn{1}{c}{(9)}& \\multicolumn{1}{c}{(10)}& \\multicolumn{1}{c}{(11)}& \\multicolumn{1}{c}{(12)}\\\\ \\hline ~~~0809+404 &08 12 53.131 &40 19 00.09&Q&19.50&~0.551&1068&~~~26.63& 392& $-$0.81&14.87&~~~~54.9\\\\ ~~~0949+287 &09 52 06.090 &28 28 32.35&G&20.15&~~~~-- &1364&~~~27.38$\\ast$& 529& $-$0.76&308.79&1330.0$\\ast$\\\\ ~~~1159+395 &12 01 49.982 &39 19 11.26&G&23.32&~2.370&~~599&~~~27.78& 249& $-$0.71&41.08&~~161.3\\\\ ~~~1315+396 &13 17 18.653 &39 25 28.02&Q&18.20&~1.560&~~615&~~~27.39& 227& $-$0.80&33.77&~~143.8\\\\ ~~~1502+291 &15 04 26.715 &28 54 30.55&q&18.58&~0.056?&~~567&~~~24.29?& 261& $-$0.63&41.60&~~~~30.8?\\\\ ~~~1616+366 &16 18 23.546 &36 32 01.33&G&19.35&~0.734&~~536&~~~26.60&268& $-$0.56&$\\sim$~60.00&~~242.0\\\\ \\hline \\end{tabular} \\end{center}   \\vspace{0.5cm} {\\small Description of the columns: \\begin{itemize} \\item[---]{Col.~~(1): Source name in the IAU format;} \\item[---]{Col.~~(2): Source right ascension (J2000) extracted from FIRST;} \\item[---]{Col.~~(3): Source declination (J2000) extracted from FIRST;} \\item[---]{Col.~~(4): Optical identification: G - galaxy, Q - quasar, q - star-like object without known redshift;} \\item[---]{Col.~~(5): Red magnitude extracted from SDSS/DR4;} \\item[---]{Col.~~(6): redshift;} \\item[---]{Col.~~(7): Total flux density at 1.4\\,GHz extracted from FIRST;} \\item[---]Col.~~(8): Log of radio luminosity at 1.4\\,GHz in W~Hz$^{-1}$; \\item[---]{Col.~~(9): Total flux density at 4.85\\,GHz extracted from GB6;} \\item[---]{Col.~~(10): Spectral index between 1.4 and 4.85\\,GHz calculated using flux densities in columns (7) and (9);} \\item[---]{Col.~(11): Largest Angular Size (LAS) in milliarcseconds measured in the 1.65-GHz VLBA image --- in most cases, as a separation between the outermost component peaks; in one case (denoted with``$\\sim$'') measured in the image contour plot;} \\item[---]{Col.~(12): Largest Linear Size (LLS) in $h^{-1}$~pc.}  Values denoted with ``$\\ast$'' were computed assuming median redshift $z=1.085$.\\\\ Redshift quoted for 1502+291 is the redshift of Abell\\,2022 cluster. \\end{itemize} }  \\label{table1} \\end{table*} ",
        "conclusions": "\\subsection{CSOs with arcsecond-scale relic extensions}  According to \\citet{pol03} and references there\\-in, CSOs are young radio sources as their kinematic ages are of the order of $10^{3}-10^{4}$ years. Recently, \\citet{gug05} have investigated the ages of CSOs in a systematic manner and have shown that there is a clear cutoff in the age distribution at approximately 500~years, suggesting that CSOs may be young, not only because they are in the initial stages in an evolutionary chain, but that they are also short-lived i.e. their activity phase lasts for only a few hundred years. The most straightforward cause of this is a lack of stable fuelling. It follows that, because of a cutoff of the energy transport from the core to the lobes, not only does the luminosity of the source drop, but also diffuse radio lobes showing an absence of edge-brightening result as the hotspots fade away quickly.  It could be that low frequency observations might reveal remnants of earlier stages of activity as in the case of 0108+388, which is a compact double with an arcsecond-scale relic extension $\\sim$$20\\arcsec$ east of the nucleus \\citep{b90}. \\citet{ocp98} have shown that the bright component of 0108+388 is a very young source with a kinematic age of 367\\,years. To explain the diffuse structure of 0108+388 they have adopted a scenario of recurrent activity in the nucleus and have proposed an interpretation of the asymmetries of the extended emission as being caused by light travel time effects.  Contrary to this, \\citet{b90} have suggested that 0108+388 could be a normally aged radio galaxy in which most of the radio emitting plasma is unable to escape from the nuclear region. Such a situation might arise if the host galaxy has recently swallowed a gas rich companion, that has smothered the source. The idea of a recent merger event is also a very plausible explanation of the misalignment between the active (inner) and inactive (outer) parts of the source.  0108+388 is the first known example of a CSO with an arcsecond-scale structure. Moreover, it has also been classified as a GPS source because of a spectral turnover, which, according to \\citet{marr01}, results from free-free absorption by nonuniform gas, possibly in the form of a disk in the central tens of parsecs. Instabilities in such a disk could result in a periodic infall of gas, that would produce apparent renewed activity. \\citet{car98} have also found significant H\\,I absorption along the line of sight to the core of 0108+388, suggesting the existence of a large amount of thermal gas.  Another example of a CSO possessing a (relatively) large-scale structure is 0402+379 \\citep{man04}. It has an arcsecond-scale, core-dominated structure and has been classified as an MSO. Unlike 0108+388, the outer lobes of 0402+379 are symmetric.  \\subsection{The case of 0809+404}\\label{s-0809}  As the VLA images of 0809+404 (F2001) show it to have a very asymmetric double structure with 1\\farcs2 separation and a flux density ratio $\\sim$100:1 at 4.9\\,GHz, we suggest that, as far as the arcsecond structure is concerned, 0809+404 resembles 0108+388. The weaker western component of the 0809+404 VLA structure is a relic of previous activity and its spectral index, calculated from the 4.9 and 8.5\\,GHz VLA observations, is very steep ($\\alpha=-1.5$). This component is not present in the 15-GHz VLA image. There is also no indication of a radio core located elsewhere in this image. However, on a milliarcsecond scale, the differences between the images of 0108+388 \\citep[see e.g.][]{tay96}, and ours of 0809+404 become apparent: 0108+388 is a triple with mini-FR\\,II-like lobes and a core, whereas no clear FR\\,II-like structure is present in 0809+404, although there is a hint of the existence of two lobes in the 1.65-GHz image. Neither a core nor hotspots are observable at any frequency and the whole milliarcsecond-scale structure fades away towards higher frequencies, indicating that activity in the nucleus has switched off.  It can be seen that there is an appreciable ($\\sim$$40\\degr$) misalignment between the axis of the inner structure with respect to the outer, relic structure. Perhaps, as in the case of 0108+388, a recent merger event in 0809+404 is also a likely scenario.  The eastern bright component shows little Faraday rotation but is strongly depolarized \\citep[F2001,][]{fanti04}. The polarization asymmetry is very common among small-scale CSS sources and can be caused by differences in the gas density in the surrounding medium \\citep{tsm03}. We have begun a programme using the WSRT to investigate the gaseous medium of 23 sources from the parent sample (Paper~I) with known redshifts. So far three of them have been observed during WSRT service time at UHF-high frequencies (700--1200\\,MHz) and 0809+404 was one of them. Approximately 3 hours of observing time were spent on each object, but none of them has shown any H\\,I absorption. For 0809+404, a 2$\\sigma$ upper limit of $\\tau$$\\sim$0.005 for the optical depth has been set based upon a noise level of $\\sim$3.7\\,mJy/beam. The results of these observations will be described in detail in a future paper.  \\citet{stan03} and \\citet{stan05} have discussed in detail the properties of a few GPS objects with extended arcsecond-scale emission and have shown that the extended emission around CSO/GPS galaxies is well explained as the co-presence of past and new activity. GPS quasars without a CSO morphology, but having arcsecond-scale emission, are likely to be intrinsically large, old and active radio sources seen along the radio axis, whereas CSO morphologies in quasars, or at least mini-lobe dominated structures, may be signatures of a small and young radio source.  0809+404 is classified spectroscopically as a quasar in SDSS/DR4. More precisely, \\citet{zakam03} have found this object to be one of 291 type-II quasar candidates. The spectra of these objects are dominated by narrow emission lines, so their broad-line emission region as well as the UV-continua are completely obscured at optical wavelengths by a dusty torus. It can be assumed that the axes of such objects are perpendicular to the line of sight. This assumption and the fact that the integrated spectrum of 0809+404 is not of a GPS type (M.~Murgia, priv. comm.) is in agreement with the conjecture that this object is not beamed to the observer. Besides, based on his findings, it is plausible that the inner structure of 0809+404 is younger than the arcsecond-scale emission.  The radio spectrum of 0809+404 is well described by a continuous injection (CI) model (which refers to the source as a whole) with $\\alpha_{inj}=-0.4$ and a break frequency $\\nu_{br}$= 1.56\\,GHz (M.~Murgia, priv. comm.), suggesting that the radio source is continuously replenished by a constant flow of fresh relativistic particles. The initial spectral index $\\alpha_{inj}$, which is the spectral index of the synchrotron radiation in the part of the spectrum not affected by the evolution, implies a power law for the energy distribution of the injected electrons of $\\delta\\approx$1.8. During the CI phase, the electrons lose energy by synchrotron emission and inverse Compton scattering of the cosmic microwave background photons.  The continuous injection model does not contradict the possibility of intermittent activity, since, during the lifetime of the extended radio emission, the nucleus appears ``on average'' active. Assuming minimum energy conditions we have estimated a source age $t_{sync}\\sim 10^{5}$ years using the formulae from \\citet{miley} (eq.\\,[7]), where $\\theta_{x}$ and $\\theta_{y}$ correspond to the beam widths of the 4.9-GHz VLA map (produced from the data extracted from the VLA archive). The minimum energy magnetic field was calculated by using eq.\\,[2] from \\citet{miley} and the standard assumptions about the lower and upper cutoff frequencies (10\\,MHz and 100\\,GHz respectively), the uniform filling factor ($\\eta=1$), the ratio between protons and electrons ($k=1$), the angle between the magnetic field and the line of sight (90\\degr) and an equivalent field of the cosmic background radiation ($B_{CMB}=3.25(1+{\\it z})^2~[{\\rm \\mu G}]$). The calculations were made for 4.9\\,GHz and the corresponding flux density was taken from F2001.  A number of theories trying to explain the episodic activity of radio sources exist in the literature. According to one, activity could be initiated as a result of a merger event. Torques and shocks during the merger can remove angular momentum from the gas in the merging galaxies and this provides injection of substantial amounts of gas/dust into the central nuclear regions \\citep{mh96}. It is, therefore, likely that in the initial phase of an AGN, this gas still surrounds (and possibly obscures) the central regions. Such activity can perhaps cease when there is no more matter to be accreted.  \\citet{tin03} have also suggested that at least some GPS sources are limited in their development by the effects of merger activity and the resulting likely sporadic fuelling of the central black holes and accretion disks. \\citet{bar96}, who tracked the evolution of both gas and stars in the merger of two disk galaxies, show that in the final state of such a merger, 60\\% of the gas is driven into the inner part of the galaxy to within 100\\,pc of the nucleus. In such an environment, a black hole may undergo many fuelling events and each event may completely disrupt and/or restart the jet. After each renewal, the jet may need to force its way through the dense nuclear environment anew.  Alternatively, to interpret the nature of the inner structure of 0809+404 a theory of the Super\\-mass\\-ive Black Hole (SMBH) accretion disk instabilities \\citep[][and references therein]{hat01,jan04} could be adopted. It takes into account a non-stationary accretion due to hydrogen ionization that can develop in the disk. These instabilities can cause a sudden accretion of the material surrounding the source and accumulated in the outer torus. It also predicts that galaxies spend the greater part of their lifetime, say $\\sim$70\\%, in a ``quiescent'' state and $\\sim$30\\% in an active state with the length of the active phase of an AGN as well as the timescale of the re-occurrence of activity being determined by the mass of the SMBH. Specifically, if the SMBH mass is assumed to be of the order of $10^7$ M$_\\odot$ -- and such an assumption is plausible for RLAGNs \\citep{wo02,osh02} -- the length of the activity period may be as low as $\\sim$$10^3$ years. This means that the transition to the fader phase can also happen at a very early stage of evolution i.e. at the CSO stage. Note also that such an interpretation is in agreement with an early hypothesis on the episodic nature of CSOs given by \\citet{r94} and recent findings of \\citet{gug05}. However, there is some inconsistency if such an interpretation is applied to 0809+404, namely that a misalignment between the inner and outer parts of the source is observed. This indicates that the previous period of activity may have been linked to a recent merger. It is to be noted that misalignments between the inner and outer parts of the source can be observed regardless of the source scale. 0809+404 and 0108+388 are good examples of very compact sources accompanied by arcsecond-scale relics. In 1245+676, the inner CSO part shows a modest misalignment with respect to the outer megaparsec-scale structure \\citep{mar16} whereas in a few core-dominated sources shown by \\citet{mtm06} large misalignments of the kiloparsec-scale structures with respect to the outer large-scale ones are quite common.  Another interpretation of the nature of this source is also conceivable, namely that the brighter eastern component of the arcsecond-scale structure is a radio lobe. However, even in this scenario, the lack of a visible hotspot in the eastern, dominating component clearly suggests that the source as a whole is a fader.  \\subsection{The other five sources}  The morphology of the edge-brightened lobes of 0949+287 as seen at 1.65\\,GHz suggests this source is an FR\\,II like object, with a steep integrated spectrum \\citep{bol04}. To calculate its linear size and power we adopted a redshift of $z=1.085$, which is the median value of all the available redshifts of the 60 sources from our primary sample. The estimated linear size and the calculated power at 1.4\\,GHz (Table~\\ref{table1}) indicate that 0949+287 could be a bright CSS object.  1159+395 is a highly redshifted galaxy with a double structure. It is the most powerful object in our sample (Table~\\ref{table1}). It appears to be unresolved in the VLA images at 4.9 and 8.5\\,GHz (F2001) so it can be assumed that it has no extended emission. Its double structure with compact features identified as hotspots as well as a non-GPS spectrum \\citep{mur99} suggest that it is a typical ``active'' CSO. According to \\citep{mur99}, 1159+395 is a very young source with a synchrotron age of the order of $\\sim10^3$\\,years.  The remaining three sources from the subsample presented here (1315+396, 1502+291, 1616+366) have very compact core-jet structures clearly indicating that these sources are not CSOs so that, even if potentially they could be intermittent, they are in an active phase at the current epoch."
    },
    "0601/astro-ph0601241_arXiv.txt": {
        "abstract": "\\psrb{} is in a highly eccentric 3.4 year orbit with a Be star and crosses the Be star disc twice per orbit, just prior to and just after periastron. Unpulsed radio, X-ray and gamma-ray emission observed from the binary system is thought to be due to the collision of pulsar wind with the wind of Be star.  We present here the results of new \\xmm\\ observations of the \\psrbsp system during the beginning of 2004 as the pulsar approached the disc of Be star. We combine these results with earlier unpublished X-ray data from \\sax\\ and \\xmm\\ as well as with \\asca\\ data. The detailed X-ray lightcurve of the system shows that the pulsar passes (twice per orbit) through a well-defined gaussian-profile disk with the half-opening angle (projected on the pulsar orbit plane) $\\Delta\\theta_{disk}\\simeq 18.5^\\circ$. The intersection of the disk middle plane with the pulsar orbital plane is inclined at $\\theta_{disk}\\simeq 70^\\circ$ to the major axis of the pulsar orbit.  Comparing the X-ray lightcurve to the TeV lightcurve of the the system we find that the increase of the TeV flux some 10--100 days after the periastron passage is unambiguously related to the disk passage.  At the moment of entrance to the disk the X-ray photon index hardens from $\\Gamma\\simeq 1.8$ up to $\\Gamma\\simeq 1.2$ before returning to the steeper value $\\Gamma\\ge 1.5$. Such behaviour is not easily accounted for by the model in which the X-ray emission is synchrotron emission from the shocked pulsar wind. We argue that the observed hardening of the X-ray spectrum is due to the inverse Compton or bremsstrahlung emission from 10-100 MeV electrons responsible for the radio synchrotron emission. ",
        "introduction": "\\psrb{} is a $\\sim$48 ms radio pulsar in a highly eccentric (e$\\sim$0.87), 3.4 year orbit with a Be star SS 2883 \\citep{johnston92}. The system's distance from Earth is $d\\sim 2$ kpc \\citep{johnston94}, note, however, that $d$ is uncertain by a factor of $\\sim$2 based on radio measurements (see, e.g., Manchester, Taylor, \\& Lyne 1993).  Be stars are well-known to be sources of strong highly anisotropic matter outflow. Both a dilute polar wind and a denser equatorial disc have been invoked to reconcile models for infra-red, ultra-violet and optical observations \\citep{waters88}. Timing analysis of the \\psrb{} system shows that the disc of Be star is tilted with respect to the orbital plane \\citep{wex98,wang04}. The properties of the radio emission from the system are very different close to and far from the periastron.  Radio observations of the 1994, 1997 \\citep{johnston99}, 2000 \\citep{connors02}, and the 2004 \\citep{johnston05} periastron passages show that when the pulsar is far from periastron the observed radio emission is comprised entirely in highly linearly polarized pulsed emission from the pulsar itself with an intensity practically independent of the pulsar orbital position \\citep{McClure98}.  But, starting about 100 days before periastron, depolarization of the pulsed emission occurs, the dispersion measure and the absolute value of the rotation measure increase while the flux density decreases. The pulsed radio emission then disappears entirely as the pulsar enters the disc and is hidden behind it (relative to the observer) approximately 20 days before the periastron passage.  Shortly before the disc crossing unpulsed radio emission appears and within several days sharply rises to a peak which is more than ten times higher than the intensity of the pulsed emission far from the periastron.  Afterward the unpulsed flux slightly decreases, as the pulsar passes through periastron before reaching a second peak just after the pulsar crosses the disc for the second time.  The unpulsed radio emission is detected until at least 100 days after the periastron passage. \\citep{johnston99,johnston05,connors02}.  The first detection of the X-rays from \\psrb\\ system was done by \\textit{ROSAT} in 1992 \\citep{cominsky94}. The most recent published observations of the system in X-rays were carried out with the \\ascasp satellite in 1994 and 1995 \\citep{kaspi95,hirayama99}. These observations show that X-ray emission is approximately twice as high at the time of the disc crossing than at periastron. In soft gamma-rays the system was observed only after the periastron passage with CGRO/OSSE and \\intgr/ISGRI \\citep{grove95,shaw04}.  The combined soft gamma- and X-ray spectrum is consistent with a power law of photon index $\\sim 1.7$.  This index also coincides with the radio spectral index of the unpulsed emission. No pulsed X-ray emission was detected from the system. With the standard epoch-folding method the upper limit to a pulsed component in 2 -- 10 Kev energy range was estimated to be about 10\\% of the unpulsed component close to the periastron, and to about 40\\% during the apastron \\citep{kaspi95,hirayama99}. Note that search for the pulsations in P - \\.P space allowed \\citet{kaspi95} to reduce an upper limit close to the periastron to less than 2\\% under a certain assumptions on the shape of pulse profile.   The 2004 periastron passage was observed in TeV gamma-rays by HESS \\citep{aharon05}. TeV lightcurve of the source shows an unexpected behaviour. Contrary to the prediction that the TeV flux should be maximum at the periastron \\citep{kirk02}, the flux apparently has a local minimum around this moment.  Several theoretical models were put forward to explain the behaviour of the X-ray lightcurve of the system. In the model of \\citet{tavani97}, the enhancement of {\\it synchrotron} X-ray emission during the disk passage is attributed to enhancement of magnetic field at the position of the contact surface of pulsar and stellar winds. In the model of \\citet{cher99,cher00} the enhancement of the {\\it inverse Compton} X-ray emission during the disk passage is supposed to be due to the effect of the macroscopic mixing of the stellar and pulsar winds in the disk which effectively enhances the escape time of high-energy electrons responsible for the observed emission.  Typical energies of electrons responsible for the X-ray emission are $E_e\\sim$~TeV in \\citet{tavani97} model while in \\citet{cher99} model $E_e\\sim 10$~MeV.  One can, in principle, distinguish between the two models by comparing the details of radio, X-ray and TeV lightcurve of the system.  We have organized a cycle of five \\xmmsp observations as the pulsar approached the 2004 periastron. In this paper we present the results of our observations along with other \\xmm\\ observations of the system not previously published. For the completeness we also include the analysis of the previously unpublished \\sax\\ data. The 2004 periastron passage was extensively observed in radio \\citep{johnston05} and TeV \\citep{aharon05} bands. Combination of the radio, X-ray and TeV data opens up a possibility, to study the evolution of the spectrum of \\psrb\\ system over some 12 decades in energy.  Our analysis of large X-ray data set enables, for the first time, to constrain the geometrical parameters of the disk of Be star, such as the opening angle and orientation. Comparing the TeV and X-ray lightcurves we find that the increase of the TeV flux some 10 to 100 days after the periastron passage  is unambiguously related to the pulsar passage through the Be star disk for the second time. This challenges the conventional interpretation of the observed TeV emission as the inverse Compton emission from the pulsar wind electrons.  Another surprising result of our analysis that is not predicted in the theroetical models cited above is the persistently hard X-ray spectrum (photon index $\\Gamma\\simeq 1.2-1.3$) observed during the 10-day period when the pulsar enters the disk.  This paper is organized as follows: in Section 2 we describe the details of the \\xmmsp and \\sax\\ data analysis.  The results are presented in Section 3, and discussed in Section 4.  \\section[]{Observations and Data Analysis}  In Figure \\ref{xmm_ellips} we show a schematic drawing of the pulsar orbit along with the approximate location of the pulsar position during \\xmm, \\sax, and \\asca\\ observations. The X-ray observations are more-or-less homogeneously distributed over the whole range of pulsar orbital phase $0\\le\\theta<360^\\circ$. The figure also shows the location of the Be star disk (parameters of the disk are found below from the analysis of the X-ray data). One can see that the set of observations denoted as X6 -- X10 enables to study the entrance to the disk in great details. \\begin{figure} \\includegraphics[width=8.5cm,angle=0]{PSRB_scheme.eps} \\caption{Schematic representation of  PSR B1259-63 binary system with the locations of the pulsar during \\xmm\\, \\sax\\, and \\textit{ASCA} observations. Shaded area shows the geometry of the disk inferred from the X-ray data. Darker and lighter shading corresponds to 1 and 2 half-opening angles of the disk, respectively. Angle $\\theta$ is the orbital phase of the pulsar referred in the text.} \\label{xmm_ellips} \\end{figure}   \\subsection{\\xmm\\ observations}  The log of the \\xmm\\ data analyzed in this paper is presented in Table 1.  Observations X6 -- X10 were obtained as part of the 2004 periastron campaign and we have also used all other \\xmm\\ observations of the system, which are now public (X1 -- X5).  In the Table negative $\\tau$ denotes the number of days before the date of 2004 periastron passage (2004 March 7), and positive values correspond to the number of days after the 2000 periastron (2000 October 17). $\\theta$ is the orbital phase counted as shown in Fig. \\ref{xmm_ellips}. \\begin{table} \\caption{Journal of \\xmm\\ observations of \\psrb \\label{data}} \\begin{tabular}{c|c|c|c|c|c} \\hline Data& Date & MJD& $\\tau$&$\\theta$&Exposure \\\\ Set&      &    &  (days)&(deg)&(ks)         \\\\ \\hline X1&  2001-01-12 & 51921.73 & 87.2    & 320.8 &  11.3 \\\\ X2&  2001-07-11 & 52101.31 & 266.7   & 342.9 &   11.6 \\\\ X3&  2002-07-11 & 52467.24 & $-$604.1& 0.5   &  41.0 \\\\ X4&  2003-01-29 & 52668.27 & $-$403.1& 9.31  & 11.0 \\\\ X5&  2003-07-17 & 52837.53 & $-$233.9& 19.5  &  11.0 \\\\ X6&  2004-01-24 & 53028.79 & $-$42.6 & 57.4  &   9.7 \\\\ X7&  2004-02-10 & 53045.43 & $-$25.0 & 73.4  &   5.2  \\\\ X8&  2004-02-16 & 53051.39 & $-$20.0 & 83.1  &    7.7 \\\\ X9&  2004-02-18 & 53052.02 & $-$18.4 & 86.5  &   5.2 \\\\ X10& 2004-02-20 & 53055.82 & $-$15.6 & 93.2  &    6.9 \\\\ \\hline \\end{tabular} \\end{table}   The \\xmmsp Observation Data Files (ODFs) were obtained from the on-line Science Archive\\footnote{http://xmm.vilspa.esa.es/external/xmm\\_data\\_acc/xsa/index.shtml}; the data were then processed and the event-lists filtered using {\\sc xmmselect} within the Science Analysis Software ({\\sc sas}) v6.0.1. In all observations the source was observed with the MOS1 and MOS2 detectors.  The X1 -- X5 observations were done in the Full Frame Mode, while the 2004 observations were performed in the Small Window Mode, to minimize the pile-up problems. For the X6 -- X10 observations PN data are also available.  In all observations a medium filter was used.  The event lists for spectral analysis were extracted from a 45$^\\prime$$^\\prime$ radius circle at the source position for the X1 -- X5 observations, and from a 22.5$^\\prime$$^\\prime$ radius circle for MOS 1,2 observations in Small Window Mode (X6 -- X10). For the PN instrument a region of 35$^\\prime$$^\\prime$ around the source position was chosen.  Background photons were collected from a region located in the vicinity of the source with the same size area as the one chosen for the source photons.  For the spectral analysis, periods of soft proton flares need to be filtered out. To exclude them we have built light curves with a hundred second binning and excluded all time bins in which the count rate above 10 keV was higher than 1.5 cnt/s.   The source countrate in 1 -- 10 keV energy range  varied from about 0.07  to 1.2 cts/s for MOS instruments, and from 0.4 to 4 cts/s for PN. Data from MOS1, MOS2 and, when available, PN detectors were combined in the spectral analysis to achieve best statistics.    \\subsection{\\sax\\ observations} Table \\ref{saxobs} lists the observations of \\psrb{} made with \\sax\\ around the 1997 periastron passage.  In the Table $\\tau$ denotes the number of days before or after the 1997 periastron passage.  We used data of MECS instruments aboard BeppoSAX observatory which provides the coverage in the energy range $\\sim$1.5-11.0 keV. Reduction of MECS data was done with the help of standard tasks of LHEASOFT/FTOOLS 5.2 package according to recommendations of BeppoSAX Guest Observer Facility\\footnote{http://heasarc.gsfc.nasa.gov/docs/sax/abc/saxabc/saxabc.html, http://www.asdc.asi.it/bepposax/software/cookbook/}. Spectrum of the source was obtained extracting events from a circle of $3^\\prime$ radius around the source position. Spectrum of the instrument background was estimated considering events from a circle of the same radius $6^\\prime$ away from the source. Typically the source provided several hundreds of counts per observation. Response matrices (RMF and ARF) of MECS detectors were taken from HEASARC archive\\footnote{http://heasarc.gsfc.nasa.gov/FTP/sax/cal/responses}. Corrections of photons arrival time to the Solar system barycenter for subsequent usage of the event lists in timing analysis were done with the help of tasks from SAXDAS package\\footnote{http://www.asdc.asi.it/bepposax/software/saxdas/index.html}.  \\begin{table} \\caption{Journal of \\sax\\ observations of \\psrb \\label{saxobs}} \\begin{tabular}{c|c|c|c|c|c} \\hline Data& Date & MJD& $\\tau$&$\\theta$&Exposure \\\\ Set&      &    &  (days)&(deg)&(ks)         \\\\ \\hline S1& 1997-03-22  & 50529.63 & $-$68.31 & 44.9& 28.622\\\\ S2& 1997-09-02  & 50693.75 &  95.81 &  322.9& 16.985\\\\ S3& 1997-09-08  & 50699.44 & 101.50 &  324.1& 53.555\\\\ S4& 1997-09-17  & 50708.08 & 110.14 &  325.9& 51.301\\\\ S5& 1997-09-25  & 50716.05 & 118.11 &  327.3& 28.719\\\\ \\hline \\end{tabular} \\end{table} We have checked that all the spectra obtained in the S2 -- S5 observations are consistent with each other, and combined them in order to improve statistics. This combined spectrum we denote as SC hereafter. ",
        "conclusions": "In this paper we have presented the \\xmm\\ observations of the pulsar PSR B1259-63 during its 2004 periastron passage near the companion Be star SS 2883. Combining our observations with the previous X-ray observations from \\xmm, \\sax\\ and \\asca\\ we  produced the detailed lightcurve of the system over the wide range of the orbital phases.  Using the large X-ray data sample we were able to constrain the geometry of the Be star disk. In particular, we have found that the half-opening angle of the disk (projected on the pulsar orbital plane) is $\\Delta\\theta_{disk}\\simeq 18.5^\\circ$ and that the line of intersection of the disk with the pulsar orbital plane is inclined at about $70^\\circ$ with respect to the major axis of the pulsar orbit.  Inspection of the X-ray and TeV lightcurves of the system has revealed correlation of the X-ray and TeV flux variations. Our analysis shows that the variations of the TeV flux some 10-100 days after the periastron passage can be well fit with the gaussian curve whose parameters coincide with the parameters of the disk found from the X-ray analysis (center at $\\theta_{disk}=289.1^\\circ$, half-width $\\Delta\\theta=18.5^\\circ$) (see Fig. \\ref{gauss}). If the observed behaviour of the flux is indeed due to the influence of Be star disk, it can not be explained within the inverse Compton scattering model of TeV emission from the system.  Our \\xmm\\ observations which were done during the period of first pulsar entrance to the disk (some 40 to 20 days before the periastron) revealed quite unexpected evolution of the spectrum of X-ray emission. In particular, the X-ray spectral slope hardened from $\\Gamma\\simeq 1.8$ down to $\\Gamma\\simeq 1.2$ roughly at the onset of the rapid growth of the X-ray luminosity. This behaviour can not be explained within the synchrotron model of X-ray emission from the system. We have shown that the observed hardening can be the result of the hardening of electron spectrum due to the Coulomb losses in the disk in the model where X-ray emission is the inverse Compton emission from the population of electrons responsible for the radio synchrotron emission. Otherwise, the hardening can be caused by the increase of the bremsstrahlung emission from these electrons. However, the bremsstrahlung problem faces a serious energy balance problem."
    },
    "0601/astro-ph0601594_arXiv.txt": {
        "abstract": "Weak gravitational lensing has several important effects on the cosmic microwave background (CMB): it changes the CMB power spectra, induces non-Gaussianities, and generates a $B$-mode polarization signal that is an important source of confusion for the signal from primordial gravitational waves. The lensing signal can also be used to help constrain cosmological parameters and lensing mass distributions. We review the origin and calculation of these effects. Topics include: lensing in General Relativity, the lensing potential, lensed temperature and polarization power spectra, implications for constraining inflation, non-Gaussian structure, reconstruction of the lensing potential, delensing, sky curvature corrections, simulations, cosmological parameter estimation, cluster mass reconstruction, and moving lenses/dipole lensing. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601077_arXiv.txt": {
        "abstract": "We present 2~\\micron\\ polarization measurements of positions in the BN region of the Orion Molecular Cloud (OMC-1) made with NICMOS Camera 2 ($0.2''$ resolution) on {\\it Hubble Space Telescope}. Our goals are to seek the sources of heating for IRc2, 3, 4, and 7, identify possible young stellar objects (YSOs), and characterize the grain alignment in the dust clouds along the lines-of-sight to the stars. Our results are as follows: BN is $\\sim 29$\\% polarized by dichroic absorption and appears to be the illuminating source for most of the nebulosity to its north and up to $\\sim 5''$ to its south. Although the stars are probably all polarized by dichroic absorption, there are a number of compact, but non-point-source, objects that could be polarized by a combination of both dichroic absorption and local scattering of star light. We identify several candidate YSOs, including an approximately edge-on bipolar YSO $8.7''$ east of BN, and a deeply-embedded variable star. Additional strongly polarized sources are IRc2-B, IRc2-D, and IRc7, all of which are obviously self-luminous at mid-infrared wavelengths and may be YSOs. None of these is a reflection nebula illuminated by a star located near radio source I, as was previously suggested. Other IRc sources are clearly reflection nebulae: IRc3 appears to be illuminated by IRc2-B or a combination of the IRc2 sources, and IRc4 and IRc5 appear to be illuminated by an unseen star in the vicinity of radio source I, or by Star n or IRc2-A. Trends in the magnetic field direction are inferred from the polarization of the 26 stars that are bright enough to be seen as NICMOS point sources. Their polarization ranges from $\\lesssim 1$\\% (all stars with this low polarization are optically visible) to $> 40$\\%. The most polarized star has a polarization position angle different from its neighbors by $\\sim 40^\\circ$, but in agreement with the grain alignment inferred from millimeter polarization measurements of the cold dust cloud in the southern part of OMC-1. The polarization position angle of another highly-polarized, probable star also requires a grain alignment and magnetic field orientation substantially different from the general magnetic field orientation of OMC-1. ",
        "introduction": "The Orion Molecular Cloud 1, OMC-1, centered $1'$ northwest of the Trapezium stars of the Orion Nebula \\ion{H}{2} Region, is the closest and best-studied site of massive star formation. The brightest near-infrared (NIR) and mid-infrared (MIR) source in the region, the Becklin-Neugebauer object BN (Becklin \\& Neugebauer 1967), is possibly a runaway B star from the Trapezium cluster (Plambeck et al.\\ 1995; Tan 2004). It is still surrounded by an optically thick envelope containing silicate dust and ices, which are identified via their infrared spectral features (e.g., Lee \\& Draine 1985; Hough et al.\\ 1996). Genzel \\& Stutzki (1989) list the indicators of star formation in their review of the literature to 1989: At MIR and far-infrared (FIR) wavelengths, BN is outshone by the extended nebulosity $\\sim 10''$ south called the Kleinmann-Low Nebula (KL, Kleinmann \\& Low 1967), which has a total luminosity of $6 \\times 10^4$ to $1.2 \\times 10^5 L_\\odot$ (Genzel \\& Stutzki 1989). In addition to the high luminosity, indicators of massive star formation in the KL region include OH, H$_2$O, and CH$_3$OH masers, and a dense, warm molecular cloud called the ``hot core'', which emits an extremely rich spectrum arising from an extensive number of complex molecules (see Fig.~6 of Genzel \\& Stutzki 1989 for a cartoon illustrating the locations of these sources).  At higher spatial resolution, KL is seen to consist of numerous components, all extended, whose coordinates and finding charts are given by Rieke et al.\\ (1973), Gezari et al.\\ (1998), and Shuping et al.\\ (2004). Controversy has arisen over how many of these ``IRc'' sources, if any, contain massive stars that are powering the FIR luminosity. At L band (3.8 \\micron) and MIR wavelengths, the most luminous of these, IRc2, is seen to consist of at least four compact sources (Chelli et al.\\ 1984; Dougados et al.\\ 1993; Shuping et al.\\ 2004), none of which is luminous enough to power the FIR source by itself (Gezari et al.\\ 1998). At NIR wavelengths, the other IRc sources (IRc2 is not readily visible at 2 \\micron, probably because it is behind the edge of the hot core) are all strongly polarized, with centro-symmetric polarization vectors indicating scattered light; that is, the IRc sources are illuminated by either BN (IRc1) or some star in the vicinity of IRc2 (Werner et al.\\ 1983; Minchin et al.\\ 1991; Chrysostomou et al.\\ 1994). Most of the fainter IRc sources are probably not self-luminous, but are reflection nebulae at NIR wavelengths and warm dust heated by external luminous sources at MIR wavelengths.  On the other hand, radio observations by Menten \\& Reid (1996) have identified two small ionized (possibly \\ion{H}{2}) regions: (1) between IRc2 and the hot core and (2) at the location of Star ``n'' (a very red star first identified in the NIR survey by Lonsdale et al.\\ 1982, LBLS). The star exciting the ionized region, located between IRc2 and the hot core and known as radio source ``I'' (Churchwell et al.\\ 1987), is not detected at any infrared wavelength, probably because it is behind too much extinction from the hot core. Radio source I and Star n are both intriguing because they are closer to the centroids of the masers and outflows than the candidate massive stars in IRc2. Greenhill et al.\\ (2004b) and Menten \\& Reid (1995) discuss the SiO masers surrounding radio source I, which they measured with the VLBA and VLA with 0.2 mas and $0.25''$ resolution, respectively. Greenhill et al.\\ (2004b) suggest that source I is excited by a high-mass young stellar object (YSO) and that the SiO masers arise in the bipolar wind from the surface of its accretion disk. Star n, about $3''$ southwest of source I, also has a disk seen at MIR wavelengths (Shuping et al.\\ 2004; Greenhill et al.\\ 2004a) and is near the center of the NH$_3$-line-emitting molecular cloud that includes the hot core and another cloud northwest of Star n (Shuping et al.\\ 2004). To summarize, both the star exciting radio source I and Star n appear to have disks --- with axes aligned more or less in the northeast-southwest direction --- that may be connected with the OH and H$_2$O masers in OMC-1.  In addition, there is also a strong outflow, seen in shocked H$_2$, in the northwest-southeast direction (Allen \\& Burton 1993; Schultz et al.\\ 1999; Kaifu et al.\\ 2000) that appears to have its centroid in the IRc2/Star n/source I location. It had been speculated that one of these objects is the origin of the outflow; however, it is now seen that the outflow does not line up with the outflows that appear to be associated with sources I and n (e.g., Shuping et al.\\ 2004). We note that the origin of this poorly-collimated H$_2$ outflow is not well-defined because of its large extent and wide opening angle, and so deserves further study. Bally \\& Zinnecker (2005) suggest that the outflow is the result of a stellar merger less than 1000 yr ago; if so, the stars in the region might not currently be ejecting mass in the northeast and southwest directions.  As discussed above, polarized scattered light has been used to infer the spatial locations of sources otherwise invisible due to foreground extinction. Polarization studies are also useful for tracing the component of the magnetic field that lies in the plane of the sky. Elongated interstellar dust grains in quiescent regions of space are aligned with their spin axes (normal to their major axes) parallel to the local magnetic field (see e.g., Weingartner \\& Draine 2003, Lazarian 2003, and Draine 2004 for reviews of the relevant physics). At FIR to mm wavelengths, the grains emit preferentially along their long axes, with the result that the emitted light is polarized perpendicular to the local magnetic field. Houde et al.\\ (2004) and Schleuning (1998), at 350 \\micron\\ and 100 \\micron, respectively, show that the magnetic field in OMC-1 is generally oriented northwest-southeast, with the field pinched along the northeast-southwest direction on a scale of several arcmin. At the location of BN/KL, the polarization percentage is low compared to elsewhere, but that may be due to small scale variations undetectable in their $12'' - 35''$ resolution maps.  Light can also be polarized by absorption by aligned grains if the wavelength of the light is similar to somewhat larger than the grain size (e.g., Kim \\& Martin 1995). In this case, the elongated grains absorb light whose electric vector lies preferentially along their long axes, so that the transmitted light is polarized parallel to the short axis of the grains. This is known as dichroic absorption. Because the short axes of the grains are preferentially aligned with the magnetic field, the visible and NIR polarization position angles are parallel to the component of the magnetic field in the plane of the sky. Polarization due to dichroic absorption has long been used to map out the orientation of the magnetic field in the Milky Way and other galaxies at optical and NIR wavelengths (see the review of Milky Way polarization by Heiles \\& Crutcher 2005). The high degree of polarization of BN, proportional to the extinction (Aitken et al.\\ 1985, 1989; Capps et al.\\ 1978; Minchin et al.\\ 1991) and in angular agreement with the FIR-inferred magnetic field directions, has been used to deduce that BN is polarized by dichroic absorption at all NIR and MIR wavelengths where it has been measured (e.g., Lee \\& Draine 1985; Hough et al.\\ 1996; Aitken et al. 1997; Smith et al.\\ 2000).  Emission from diffuse nebulosity can also be affected by dichroic absorption. Hough et al.\\ (1986), Burton et al.\\ (1991) and Chrysostomou et al.\\ (1994) measured the polarization of the H$_2$\\ $v = 1 - 0$ S(1) line at 2.12 \\micron. They attribute the polarization in the vicinity of IRc2 to dichroic absorption, and particularly, the twist in the polarization vectors $5''$ east of IRc2 to a twist in the magnetic field direction. On the other hand, they infer that the polarization at large distances from IRc2 and BN is due to scattering in an H$_2$ reflection nebula (RN).  In this paper we describe measurements made with the Near-Infrared Camera and Multi-Object Spectrometer (NICMOS) on the {\\it Hubble Space Telescope} ({\\it HST}) to measure the 2~\\micron\\ polarization of the region at $0.2''$ resolution. Our region extends $\\sim 20''$ north and south of BN; in the rest of this paper our references to OMC-1 refer only to this region. We describe the observations in \\S 2, our results on the polarization of both the stars, and the diffuse nebulosity in \\S 3, and summarize our conclusions in \\S 4. We will discuss the illuminating sources of the diffuse nebulosity, several of which are probably YSOs, and we will discuss the orientation of the magnetic field in OMC-1, as best we can infer it from the polarization of the stars. The locations of the illuminating sources and the magnetic field directions both have relevance to the question of the origin of the massive outflows in OMC-1. ",
        "conclusions": "This paper presents 2 \\micron\\ polarization measurements of OMC-1 made with NICMOS on {\\it HST}. Thanks to the unprecedented spatial resolution and sensitivity over the whole field of view, we are able to measure the polarization of 25 stars and 6 YSOs and candidate YSOs not previously measured at 2 \\micron. Individual objects discussed in this paper are summarized in Table~4.  The field of view can be divided into two sections by an approximately east-west line $\\sim 5''$ south of BN. BN illuminates the northern region but has no apparent influence on the polarization of the scattered light south of $-5.5''$. From this we conclude that BN is separated from the southern region, possibly by a substantial distance in the line of sight. Another possibility is a dust lane preventing light from BN from reaching the reflection nebulae to the south. Another dust lane crosses in front of BN at a position angle of $\\sim 300^\\circ$, almost in line with BN's polarization position angle of $114^\\circ$.  The southern region includes the bright reflection nebulae IRc3, IRc4, and IRc5, and the controversial sources IRc2 and IRc7. IRc2 is mostly obscured by the OMC-1 hot core at 2 \\micron, but components IRc2-B and IRc2-D are both visible in our images. This is the shortest wavelength at which either has been detected. Both sources are substantially more polarized than their immediate surroundings; however, their polarization position angles are similar to the polarization position angles of the surrounding nebulosity, the polarization of which is due to dichroic absorption. We conclude that the position angles of IRc2-B and IRc2-D are consistent with either dichroic absorption by foreground aligned dust grains or with scattered light with Star n being the illuminating source. {\\it Neither IRc2-B nor IRc2-D is a dust cloud illuminated by a star near radio source I.} Considering the MIR luminosities of all the sources in the region, we suggest that both are deeply embedded, self-luminous objects with optically thick envelopes obscuring any central stars. IRc7 is probably also a self-luminous object with an optically thick envelope scattering and polarizing the light from its central star/YSO.  The only dust clouds whose polarization vectors are consistent with illumination by a star near radio source I are IRc4 and IRc5, although this identification is not statistically significant. These reflection nebulae could also be illuminated by components of IRc2, as is IRc3. We conclude that there is no strong evidence at 2 \\micron\\ for the object that excites radio source I illuminating any other object in our field of view.  In fact, there is no evidence for any illuminating source dominating the region south of BN. The lack of evidence for such a dominating source suggests that no single object currently powers the strong outflows seen in CO and H$_2$. Evidence of a dominating source, if such existed, would contradict theories like that of Bally \\& Zinnecker (2005) that the outflows originated in an explosion some 500 -- 1000 yr ago.  The optically obscured stars are strongly polarized, indicating that they are deeply embedded in the cloud of dust grains aligned by the OMC-1 magnetic field. We use the polarization position angles of the stars to investigate the OMC-1 magnetic field, assuming that the angles are indeed parallel to the magnetic field, as is generally supposed. As such, we infer that the magnetic field in OMC-1 changes direction from $\\sim 120^\\circ$ in the region north of IRc2 to $\\sim 140^\\circ$ in the region south and east of IRc2. The objects exhibiting these position angles may be relatively foreground of the hot core --- south of the hot core there appears to be an additional cold cloud, producing the polarization seen in the mm by Rao et al.\\ (1998) and in our 2~\\micron\\ observations of the deeply embedded Star 25 at position angle $\\sim 2^\\circ$. From the observed polarization position angles we infer that the magnetic field in this cloud must have a north-south orientation (or east-west if the grains have recently been aligned by an outflow). Either orientation is greatly different from the overall magnetic field orientation foreground to the cloud. Another location with a magnetic field orientation significantly different from the overall field is IRc3, where again the magnetic field may be in either a north-south direction or an east-west direction, but $\\sim 90^\\circ$ different from that of the region in front of Star 25. These variations in magnetic field direction may be connected to the strong outflows seen in OMC-1 and centered near radio source I, IRc2, and Star n."
    },
    "0601/astro-ph0601131_arXiv.txt": {
        "abstract": "The launch of Chandra and XMM-Newton has led to important new findings concerning the X-ray emission from supernova remnants. These findings are a result of the high spatial resolution with which imaging spectroscopy is now possible, but also some useful results have come out of the grating spectrometers of both X-ray observatories, despite the extended nature of supernova remnants. The findings discussed here are the evidence for slow equilibration of electron and ion temperatures near fast supernova remnant shocks, the magnetic field amplification near remnant shocks due to cosmic ray acceleration, a result that has come out of studying narrow filaments of X-ray synchrotron emission, and finally the recent findings concerning Fe-rich ejecta in Type Ia remnants and the presence of a jet/counter jet system in the Type Ib supernova remnant Cas A. ",
        "introduction": "Supernovae are the most important sources of kinetic energy and chemical enrichment of galaxies. By studying supernova remnants (SNRs) we hope to learn about supernova explosion properties and  chemical yields. Moreover, because of their energy and large extent, SNRs are thought to be the principal sources of cosmic rays of energies up to $\\sim 10^{15}$~eV. Note that the large sizes of SNRs are also an important ingredient for their ability to accelerate cosmic rays, because astrophysical sources cannot accelerate particles beyond energies for which their gyro-radii are larger than the sources themselves. This means that for cosmic ray acceleration high magnetic fields and/or large objects sizes are required.  Their large sizes, several parsecs, make SNRs also rewarding targets for X-ray imaging spectroscopy. X-ray imaging spectroscopy made a great leap forward with the launch of the \\chandra\\ and \\xmm\\ satellites. It is now possible to obtain spectra with a spectral resolution of $E/\\Delta E \\sim 50$ at 6~keV, isolating individual regions with an accuracy ranging from $\\sim$0.5\\arcsec\\ (\\chandra) to $\\sim$5\\arcsec (\\xmm). Before 1999 \\asca\\ and \\sax\\ already provided imaging spectroscopy, but on arcminutes scales rather than arcseconds scales. The \\einstein\\ and \\rosat\\ imagers on the other hand, had imaging resolution of  $\\sim$5\\arcsec, but without any appreciable spectral resolution.  \\chandra\\ and \\xmm\\ also have dispersive high resolution spectrometers on board, three transmission gratings for \\chandra's, and two Reflective Grating Spectrometers (RGSs) for \\xmm. Due to their dispersive nature, these instruments are not ideal for spectrometry of extended objects. Nevertheless, for objects of modest extent, $< 1$~\\arcmin, the RGS is still able to obtain spectra with a resolution of $\\lambda/\\Delta \\lambda > 160$ at 20~\\AA.  So what have these instrumental advances brought us, as far as our knowledge of SNRs is concerned? The answer is that we have learned substantially about SNR shocks, concerning both the shock heating process and cosmic ray acceleration. Moreover, imaging spectroscopy has also revealed regions with metal-rich, pure ejecta plasma, and has provided us with better means to measure SNR kinematics through X-ray proper motion and Doppler shift studies.  Another important aspect is the study of supernova explosion through the analysis of kinematics of fresh explosive nucleosynthesis products. Finally, the high spatial resolution of in particular \\chandra\\ has enabled the discovery of many young neutron stars inside SNRs, such as the still enigmatic point source in \\casa\\ that was discovered in the first light image of \\chandra\\ \\citep{tananbaum99} (see also R. Petre these proceedings). However, in these proceedings I limit myself to three important topics 1) collisionless shock physics, 2)  cosmic ray acceleration, and 3) supernova explosions and nucleosynthesis, ",
        "conclusions": "An overview like this is not much more than summing up the conclusions of many individual studies. So what can I add to that, except reiterating that the successful launch of \\chandra\\ and \\xmm\\ has given us many new insights into the shock heating, cosmic ray acceleration, and composition of SNRs. Many old questions seem to have been (partially) answered, such as the question of electron-ion temperature equilibration at the shock front (electron-ion equilibration is not efficient in fast shocks), or the question whether cosmic rays can be accelerated up to the ``knee'' by supernova remnant shocks (yes they can). New questions have been raised by the new observational capabilities, such as the nature of the jet/counter jet system in \\casa.  In that respect the investigation of SNRs is very similar to other topics discussed at this meeting. So let me therefore conclude by thanking the organizers for having invited me to this very interesting and stimulating meeting.   \\itemsep -2mm"
    },
    "0601/hep-ph0601224_arXiv.txt": {
        "abstract": "In the standard model, the weak gauge bosons and fermions obtain mass after spontaneous electro-weak symmetry breaking, which is realized through one fundamental scalar field, namely Higgs field. In this paper we study the simplest scalar cold dark matter model in which the scalar cold dark matter also obtains mass through interaction with the weak-doublet Higgs field, the same way as those of weak gauge bosons and fermions. Our study shows that the correct cold dark matter relic abundance within $3\\sigma$ uncertainty ($ 0.093 < \\Omega_{dm} h^2 < 0.129 $) and experimentally allowed Higgs boson mass ($114.4 \\le m_h \\le 208$ GeV) constrain the scalar dark matter mass within $48 \\le m_S \\le 78$ GeV. This result is in excellent agreement with that of W.~de Boer et.al. ($50 \\sim 100$ GeV). Such kind of dark matter annihilation can account for the observed gamma rays excess ($10\\sigma$) at EGRET for energies above 1 GeV in comparison with the expectations from conventional Galactic models. We also investigate other phenomenological consequences of this model. For example, the Higgs boson decays dominantly into scalar cold dark matter if its mass lies within $48 \\sim 64$ GeV. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601657_arXiv.txt": {
        "abstract": "Two massive binary systems of the microquasar type, LS 5039 and LSI +61$^{\\rm o}$ 303, have been suggested as possible counterparts of EGRET sources. LS 5039 has been also recently detected in TeV gamma-rays. Since the massive stars in these binary systems are very luminous, it is expected that high energy gamma-rays, if injected relatively close to the massive stars, should be strongly absorbed, initiating inverse Compton $e^\\pm$ pair cascades in the anisotropic radiation from stellar surfaces. We investigate influence of the propagation effects on the spectral and angular features of the $\\gamma$-ray spectra emerging from these two binary systems by applying the Monte Carlo method. Two different hypothesis are considered: isotropic injection of primary gamma-rays with the power law spectrum due to e.g. interaction of hadrons with the matter of the wind, and the isotropic injection of electrons, e.g. accelerated in the jet, which comptonize the radiation from the massive star. It is concluded that propagation effects of $\\gamma$-rays can be responsible for the spectral features observed from LS 5039 (e.g. the shape of the spectrum in the GeV and TeV energy ranges and their relative luminosities). The cascade processes occurring inside these binary systems significantly reduce the gamma-ray opacity obtained in other works by simple calculations of the escape of gamma-rays from the radiation fields of the massive stars.  Both systems provide very similar conditions for the TeV $\\gamma$-ray production at the periastron passage. Any TeV $\\gamma$-ray flux at the apastron passage in LSI +61$^{\\rm o}$ 303 will be relatively stronger with respect to its GeV flux than in LS 5039. If $\\gamma$-rays are produced inside these binaries not far from the massive stars, i.e. within a few stellar radii, then clear anticorrelation between the GeV and TeV emission should be observed, provided that primary $\\gamma$-rays at GeV and TeV energies are produced in the same process by the same population of relativistic particles. These $\\gamma$-ray propagation features can be tested in the near future by the multi-wavelength campaigns engaging the AGILE and GLAST telescopes ($>$30 MeV) and the Cherenkov telescopes ($>$100 GeV, e.g. MAGIC, HESS, VERITAS and CANGAROO). ",
        "introduction": "Massive binary systems provide very promising conditions for acceleration of particles to relativistic energies and also well defined conditions for their possible interaction. Therefore, they have been considered for a long time as possible sources of high energy $\\gamma$-rays and neutrinos in which acceleration processes can be tested. In fact, some GeV $\\gamma$-ray sources observed by EGRET detector ($>100$ MeV) have been proposed to be related to well known massive binaries in which non-thermal processes were evident at lower energies, e.g. LSI +61$^{\\rm o}$ 303 (2EG J0241+6119; Thompson et al.~1995), Cen X-3 (Vestrand, Sreekumar \\& Mori~1997), Cyg X-3 (2EG J2033-4112; Mori et al.~1997), and LS 5039 (3EG J1824-1514; Paredes et al.~2000). Early searches of the TeV $\\gamma$-ray signals from massive binaries have not been very convincing, with the exception of Cen X-3 binary system (containing slowly rotating neutron star) which has been reported as TeV $\\gamma$-ray source by the Durham group (Chadwick et al.~1998, 1999, Atoyan et al.~2002). The turning-point came recently with the observations of TeV $\\gamma$-ray signals from two massive binaries: PSR B1259-63/SS 2883, containing a radio pulsar with the period of 47.8 ms (Aharonian et al.~2005a), and LS 5039, so called microquasar possibly containing a solar mass black hole (Aharonian et al.~2005b).  The $\\gamma$-ray emission from massive binaries is usually interpreted in terms of the inverse Compton scattering (ICS) model in which thermal radiation coming from the stellar surface is scattered by electrons accelerated in the pulsar wind shock (e.g. Maraschi \\& Treves~1981) or by electrons moving highly anisotropically in the form of beams or jets (e.g. Bednarek et al.~1990). The binary systems which are considered as an example in this paper, LS 5039 and LSI +61$^{\\rm o}$ 303, have been recently discussed in terms of the microquasar IC model by Bosch-Ramon \\& Paredes~(2004a,b) and Paredes, Bosch-Ramon \\& Romero~(2005). Primary $\\gamma$-rays might be also produced in the microquasar scenario of the binary system as a result of the interaction of hadrons, accelerated in the jet, with the matter of the stellar wind (Romero, Christiansen \\& Orellana~2005). For review of the $\\gamma$-ray production in microquasars see e.g. Romero (2004) or Paredes~(2005).  The problem of importance of the propagation of TeV $\\gamma$-rays inside compact massive binary systems appeared at the early 80's after first reports on possible observation of such $\\gamma$-rays (see e.g. Weekes~1988). The optical depths for $\\gamma$-rays in the radiation field of the accretion disk around the compact object inside the binary system have been calculated by e.g. Carraminana~1992 and Bednarek~(1993) and in the radiation field of the massive star by e.g. Protheroe \\& Stanev~1987 and Moskalenko, Karaku\\l a \\& Tkaczyk~1993). In the case of very compact binaries, the TeV $\\gamma$-rays injected close to the surface of the massive star initiate the Inverse Compton $e^\\pm$ pair cascade. The conditions for which the cascade processes should become important have been considered by Bednarek (1997, Sect. 2) and Dubus (2005, Sect 6). In general, the product of the massive star luminosity and the square of its surface temperature should be much greater than $\\sim 10^{45}$ erg s$^{-1}$ K$^2$ (see Eg. 1 and Fig. 3 in Bednarek 1997). The $\\gamma$-ray spectra which emerge toward the observer from such compact binary systems strongly depend on the phase of the injection place of the primary relativistic particles in respect to the observer. At some phases significant TeV $\\gamma$-ray fluxes are expected but at others only photons with energies extending to a few tens GeV are able to escape (see detailed calculations of such anisotropic cascades and their application to specific sources, e.g. Cen X-3 and Cyg X-3, in Bednarek~1997 (B97), Bednarek~2000 (B00), and Sierpowska \\& Bednarek~2005 (SB05)). Recently, optical depths of TeV $\\gamma$-rays have been calculated for other TeV $\\gamma$-ray sources, LS 5039, PSR 1259-63 and LSI +61$^{\\rm o}$ 303, without taking into considerations the effects of the IC $e^\\pm$ pair cascading (B\\\"ottcher \\& Dermer~2005 and Dubus~2005).  Note that only calculations by Bednarek and collaborations (B97, B00, SB05) and Dubus~(2005) take into account dimensions of the massive star. Therefore, they can be applied to very compact binaries in which injection distance of the TeV $\\gamma$-rays from the center of massive stars is be comparable to their radii, e.g. at the periastron passage of the compact objects in LS 5039 and LSI +61$^{\\rm o}$ 303, or in the case of propagation of $\\gamma$-rays injected at larger distances but passing close to the surface of the massive stars. The comparison of the exact calculations with the approximate ones, which neglect dimensions of the massive stars, can be found in Dubus~(2005).  In this paper we concentrate on details of the propagation of high energy $\\gamma$-rays injected close to the surface of the massive star taking into account the effects of cascades initiated through the Inverse Compton and $e^\\pm$ pair production processes. We calculate the $\\gamma$-ray spectra emerging to the observer from such anisotropic cascades for two compact massive binaries LS 5039 (already reported in GeV and TeV $\\gamma$-rays) and LSI +61$^{\\rm o}$ 303 (reported only at GeV energies). Specific cases with the injection of primary $\\gamma$-rays (e.g. by relativistic hadrons) or electrons (accelerated in the jet or the shock) with simple power law spectra and spectral indexes equal to 2 are considered in order to better understand the basic features of such anisotropic cascade processes.  \\begin{figure} \\vskip 4.2truecm \\special{psfile=fig1.eps voffset=0 hoffset=0 hscale=48 vscale=45} \\caption{Schematic picture of a compact binary system composed of a massive star and a compact object (a black hole or an neutron star) on an orbit around the massive OB star. The observer is located at the inclination angle $\\theta$ and the azimuthal angle $\\omega$. The matter accreting onto a compact object from the massive star creates an accretion disk. Particles (electrons or protons) are accelerated  inside the jet launched from the inner part of an accretion disk. Primary electrons and/or $\\gamma$-rays, injected at the distance $z$ from the base of the jet, initiate an anisotropic inverse Compton $e^\\pm$ pair cascade in the radiation field of the massive star. A part of the primary $\\gamma$-rays and secondary cascade $\\gamma$-rays escape from the binary system toward the observer.} \\label{fig1} \\end{figure} ",
        "conclusions": "We have performed Monte Carlo simulations of the propagation of high energy $\\gamma$-rays inside compact massive binaries of the microquasar type, applying as an example parameters of LS 5039 (recently observed in  TeV $\\gamma$-rays and suggested as a counterpart of the EGRET source) and  LSI +61$^{\\rm o}$ 303 (similar parameters to LS 5039 and also suggested as a counterpart of the EGRET source). In both sources the optical depths for TeV $\\gamma$-rays are greater than unity for most of the propagation directions, provided that the injection place of primary particles ($\\gamma$-rays or electrons) is relatively close to the base of the jet launched from the inner part of the accretion disks around compact objects. Therefore, these primary particles initiate IC $e^\\pm$ pair cascades in the anisotropic radiation of the massive stars. We have calculated the $\\gamma$-ray spectra produced in such cascades and investigated their basic features. It is concluded that the absorption dips in TeV $\\gamma$-ray spectra emerging toward the observer are not so drastic as predicted in the earlier papers, e.g. by B\\\"ottcher \\& Dermer~(2005) and  Dubus~(2005). This is clearly seen by comparing $\\gamma$-ray spectra calculated with pure absorption (e.g Figs.~\\ref{fig3}a, \\ref{fig4}a, \\ref{fig6}a, and \\ref{fig7}a) with the corresponding total $\\gamma$-ray spectra produced in the cascade processes inside the binary system (Figs.~\\ref{fig3}c, \\ref{fig4}c, \\ref{fig6}c, and \\ref{fig7}c). This less pronounced dips are due to the re-distribution of energy in the primary spectrum of $\\gamma$-rays from the region above $\\sim 1$ TeV  to $0.1-1$ TeV energy range. Since the soft radiation field created by the massive stars in these two binaries are quite similar during periastron passages of their compact objects, the $\\gamma$-ray spectra produced in the cascade processes do not differ significantly in the case of LS 5039 and LSI +61$^{\\rm o}$ 303. However, the propagation effects are responsible for essential differences at the apastron passage, since LSI +61$^{\\rm o}$ 303 is more extended (the apastron distance in LSI +61$^{\\rm o}$ 303 is a factor of $\\sim$2 larger than in LS 5039).  Let us concentrate at first on the features of $\\gamma$-ray emission from LS 5039. The $\\gamma$-ray luminosity from this source observed in the GeV energy range (assuming that the identification with the EGRET source by Paredes et al.~2000 is correct) is about two orders of magnitude larger than observed in TeV energy range (Aharonian et al.~2005b). However, spectra reported in these two energy ranges are quite similar (spectral index close to 2). These spectral features can be naturally explained by the propagation effects considered in this paper, i.e. IC e$^\\pm$ pair cascading processes, inside the binary system. The primary $\\gamma$-rays or electrons could be injected with a simple power law spectrum (and spectral index close to 2), relatively close to the base of the jet. The $\\gamma$-ray luminosity expected in this propagation model at GeV energies ($1-10$ GeV) should vary by a factor of $\\sim 2-3$, and at TeV energies ($>$100 GeV) by a factor of $\\sim 10$ with the orbital period of the binary system LS 5039. In fact, possible variation of the TeV signal by a factor of $\\sim 3$ has been recently suggested by Casares et al. (2005), although not statistically significant (Aharonian et al. 2005b). Moreover, we show that, due to the propagation (IC $e^\\pm$ pair cascade) effects, the maximum in TeV $\\gamma$-ray light curve should occur at phase $\\sim 0.6$ (measured as an azimuthal angle from the periastron passage, see Fig.~\\ref{fig1}). This seems to be inconsistent with the recent suggestion by Casares et al.~(2005b) who, reanalyzed the HESS data with the new orbital parameters and, concluded that maximum of TeV emission occurs at the phase $\\sim 0.9$. This discrepancy does not allow us to put any definitive conclusions since the error bars in TeV $\\gamma$-ray light curve, shown by Casares et al.~(2005b), are very large. However, if real, it can give some information on the efficiency of acceleration process of the primary particles occurring in the jet (possibly linked with the efficiency of the accretion process) with the phase of the compact object on its orbit around the massive star.  Based on the propagation calculations it is concluded that $\\gamma$-ray spectra produced in such IC $e^\\pm$ pair cascades should show clear anticorrelation between the fluxes in the GeV and TeV energy ranges (see the light curves in Fig.~\\ref{fig9}). We also investigated the dependence of the propagation effects on the distance of the injection place from the base of the jet. If primary particles are injected already at the distance of a few stellar radii from the base of the jet (the case of $z = 5r_\\star$ is discussed), then $\\gamma$-ray fluxes expected from LS 5039 do not vary strong enough to explain relative luminosities in the GeV and TeV energy ranges. Either a more complex shape for the spectrum of primary particles injected into the jet is required, e.g composed of at least three different power laws, or two different populations of primary particles are needed, or the origin of different parts of $\\gamma$-ray spectra in different regions of the jet has to be postulated. Moreover, the shape of the $\\gamma$-ray spectra escaping toward the observer depends also on the inclination angle of the binary system which is not at present well known (estimated on $25^{\\rm o}$ for the case of a solar mass black hole and on $60^{\\rm o}$ for the case of a neutron star). The level of variability of the TeV emission for the injection place of primary particles at the base of the jet is clearly larger for the inclination angle $\\theta= 60^{\\rm o}$ than for $\\theta= 25^{\\rm o}$. However these differences seems to be to low in order to put constraints on the inclination angle of binary system LS 5039 based only on the observations with the present sensitivity Cherenkov telescopes.  If the identification of the binary system LSI +61$^{\\rm o}$ 303 with the EGRET source 3EG J0241+6103 is correct (see introduction), then average $\\gamma$-ray luminosity ($>100$ MeV), $\\sim 8\\times 10^{34}$ erg s$^{-1}$, is only about a factor of three lower than in the case of LS 5039. This suggests that $\\gamma$-ray properties of these two sources should be quite similar. Our simulations of the propagation of primary $\\gamma$-rays show that absorption of $\\gamma$-rays is similar in both binaries at the periastron passage of the compact object (if primary particles are injected at the base of the jet) but is less important at the apastron passage of LSI +61$^{\\rm o}$ 303. The TeV fluxes from  LSI +61$^{\\rm o}$ 303 predicted by the propagation effects should be relatively larger in respect to its GeV fluxes than observed in the binary system LS 5039. This feature might help in detection of LSI +61$^{\\rm o}$ 303 at TeV energies. The maximum in TeV $\\gamma$-ray light curve predicted by the propagation effects in LSI +61$^{\\rm o}$ 303 should occur at the phase $\\sim 0.2-0.4$, measured from the periastron. The TeV $\\gamma$-ray signal is expected to be modulated by an order of magnitude with the maximum for the case of injection of primary electrons injected farther along the jet (at  a distance of a few stellar radii, see the case of $z = 3r_\\star$ in Fig. \\ref{fig11}). Also strong variability of the GeV flux is possible (even by an order of magnitude) for the case of injection of primary electrons farther from the base of the jet (see e.g. the case of $z = 3r_\\star$ in Fig. \\ref{fig11}b). Injection of electrons at the base of the jet predict modulation by a factor of $\\sim 2$ with the minimum when the compact object is in front of the massive star. Re-analysis of the EGRET data indicates probable modulation of the GeV signal from the source toward LSI +61$^{\\rm o}$ 303 (Tavani et al.~1998, Wallace et al.~2000) with the maximum emission near periastron (Massi 2004). This is inconsistent with the predictions of studied here propagation effects for the case of injection of primary electrons which show deficit of GeV emission close to periastron. Therefore, if this modulation is real then the mechanism responsible for such modulation of GeV emission should be even much more efficient since considered here propagation effects work in opposite direction. For the case of isotropic injection of primary $\\gamma$-rays, the GeV signal is only weakly modulated with the period of the binary system due to the domination of primary $\\gamma$-rays at such low energies.  In this paper we have only analyzed the effects of propagation of $\\gamma$-rays inside these two binary systems assuming that efficiency of particle acceleration and spectra of injected particles do not depend on the location of the injection place of primary particles along the jet, i.e. they are uniformly distributed along the jet, i.e. from $z = 0r_\\star$ to at $z = 5r_\\star$ (LS 5039) or $3r_\\star$ (LSI +61$^{\\rm o}$ 303). Proper analysis of $\\gamma$-ray emission from these sources require, accept considering the IC $e^\\pm$ pair cascade effects, a detailed model for efficiency of accretion process, conversion of accretion energy into particles, acceleration of particles (their spectra), and production of $\\gamma$-rays. Such models, which unfortunately do not take into account the cascade effects analyzed in this paper but discuss only production of $\\gamma$-rays in the Inverse Compton process far away from the base of the jet, have been recently discussed by Bosch-Ramon, Romero \\& Paredes~(2005) and Dermer \\& B\\\"ottcher~(2005)."
    },
    "0601/astro-ph0601182_arXiv.txt": {
        "abstract": "We present {\\it Swift}-UVOT data on the optical afterglow of the X-ray flash of 2005 April 6 (XRF~050406) from 88~s to $\\sim 10^5$~s after the initial prompt $\\gamma$-ray emission. Our observations in the {\\it V}, {\\it B} and {\\it U} bands are the earliest that have been taken of an XRF optical counterpart. Combining the early-time optical temporal and spectral properties with $\\gamma$- and simultaneous X-ray data taken with the BAT and XRT telescopes on-board {\\it Swift}, we are able to constrain possible origins of the XRF. The prompt emission had a FRED profile (fast-rise, exponential decay) with a duration of $\\mathrm{T}_{90}=5.7~\\mathrm{s}\\pm 0.2~\\mathrm{s}$, putting it at the short end of the long-burst duration distribution. The absence of photoelectric absorption red-ward of $4000$~\\AA ~in the UV/optical spectrum provides a firm upper limit of $\\mathrm{z}\\leq 3.1$ on the redshift, thus excluding a high redshift as the sole reason for the soft spectrum. The optical light curve is consistent with a power-law decay with slope $\\alpha = -0.75\\pm 0.26~(F_{\\nu}\\propto t^{\\alpha})$, and a maximum occurring in the first 200~s after the initial $\\gamma$-ray emission. The softness of the prompt emission is well described by an off-axis structured jet model, which is able to account for the early peak flux and shallow decay observed in the optical and X-ray bands. ",
        "introduction": "\\label{sec:intr} X-ray flashes (XRFs) are a softer class of gamma-ray bursts (GRBs) with a $\\nu F_{\\nu}$ spectrum that peaks at energies $E_{\\mathrm{p}} \\lesssim 30$~keV~\\citep{hik+01}. They make up around a third of the GRB population and their defining characteristic is a high X-ray fluence ($S_{\\mathrm{X}}$) relative to their gamma-ray fluence ($S_{\\gamma}$) such that log[$S_{\\mathrm(X)}$/$S_{\\gamma}$]$ >0$. Classical GRBs have log[$S_{\\mathrm{X}}$/$S_{\\gamma}$]$ <-0.5$, and bursts with intermediate fluence ratios are referred to as X-ray rich (XRR) bursts.  All three classes of burst are distributed isotropically on the sky and share many temporal and spectral properties. Their prompt emission light curves are similar in duration and structure, and their spectra are well fit by the `Band' function~\\citep{bmf+93}. The differentiating feature in the context of the Band model is the spectral break energy. Most GRBs have observed peak energies in the range $E_{\\mathrm{p}}\\sim 135-350$~keV~\\citep{pbm+00}, whereas XRFs have $E_{\\mathrm{p}}\\ll 100$~keV that can be as low as 3~keV~\\citep{kwh+01,bol+03,slk+05}.  Other similarities exist in a number of empirical laws that are satisfied by all three classes of bursts. The burst populations form a continuum in the [$S(2-400$~keV), $E_{\\mathrm{p}}$]-plane~\\citep{stk+03}, and the Amati relation~\\citep{aft+02} is satisfied by at least three XRFs~\\citep{ldg05,sbb+05}, which relates a GRB's peak energy $E_{\\mathrm{p}}$ to its isotropic equivalent energy $E_{\\mathrm{iso}}$ such that $E_{\\mathrm{p}}\\propto E_{\\mathrm{iso}}^{1/2}$~\\citep{ldg05}.  Theoretical models that aim to explain the origins of XRFs include (1) GRBs at high redshift~\\citep{hei03}, which in this context is at redshifts $\\mathrm{z} > 5$, (2) a baryon loaded fireball (dirty fireball model)~\\citep{dcb99}, (3) a clean fireball ($\\Gamma\\gg 300$) where the difference in Lorentz factor between colliding shells is small~\\citep{zm02,bdm+05}, (4) a photosphere-dominated model~\\citep{rl02,mrr+02}, and (5) classical GRBs seen off-axis, either within the context of a uniform jet model~\\citep{yin02} or a universal structured jet model~\\citep{zdl+04}.  The broad range of properties shared by XRFs, XRR bursts and GRBs suggests that all three classes are due to the same underlying phenomenon and that there should be a single model to explain their origins. \\citet{rl02} and \\citet{mrr+02} suggested a dominating baryonic photosphere with an extended optically thick region caused by pair-producing shocks as a source of excess X-ray emission. However, such a model requires additional considerations to account for those bursts with $E_{\\mathrm{p}} = 3-5$~keV, as is the case for XRF~020427~\\citep{acf+02}, XRF~020903~\\citep{skb+04} and XRF~010213~\\citep{ssk+04}. Models that require the GRB jet to be observed off-axis have been particularly successful at explaining observations of XRFs (e.g. Fynbo et al. 2004; Soderberg et al. 2005), and their ability to describe both GRBs and XRFs in a unified way have made them increasingly popular.  Afterglows have now been detected for a number of XRFs from the X-ray down to radio wavebands~\\citep{lg03}. The first X-ray and radio afterglows were discovered in 2001 (XRF~011030; Taylor et al. 2001; Harrison et al. 2001), and in 2002 the first XRF optical afterglow was observed (XRF~020903; Soderberg et al. 2004). Since then only a handful of XRFs have had associated optical counterparts (e.g. XRF~030429, XRF~030528, XRF~030723, XRF~040916, XRF~050315, XRF~050416a), with the earliest observation of an optical afterglow detection taking place 31.3~mins after the initial burst (XRF~030723; Smith et al. 2003).  In this paper we describe {\\it Swift} observations of XRF~050406, focusing in particular on the optical/UV band data (see~\\citet{rmb+05} for a detailed analysis of the X-ray data). These are the earliest observations taken of an XRF optical counterpart to date, beginning only 88~s after the initial emission. In section~\\ref{sec:obs} we review the observations made by all 3 instruments on-board {\\it Swift} followed by an analysis of the temporal and spectral behaviour of the XRF in the UV, optical and X-ray energy bands in sections~\\ref{sec:lc} and \\ref{sec:spec}. In section~\\ref{sec:disc} we discuss the results and investigate their implications for XRF theoretical models. The conclusions are summarised in section~\\ref{sec:concs}. ",
        "conclusions": "\\label{sec:concs} We have presented early-time optical data for XRF~050406, beginning 88~s after the burst trigger. This constitutes the earliest detection of an X-ray Flash in the optical band to date. The properties of the light curve and spectrum constrain the possible models that explain the soft spectrum from the prompt emission. The lack of Lyman absorption down to $4000$~\\AA~in the UV/optical spectrum provides a firm upper limit of $\\mathrm{z}\\leq 3.1$ on the redshift, and broadband spectral modelling yields a best-fit redshift $\\mathrm{z} = 2.44^{+0.30}_{-0.35}$, and spectral index $\\beta = -0.84^{+0.11}_{-0.06}$. A fit to the optical light curve provides a decay rate of $\\alpha = -0.75\\pm 0.26$. Combining this information with the prompt $\\gamma$-ray emission and X-ray afterglow we conclude that the soft spectrum observed in the initial emission from XRF~050406 is most likely the result of a GRB with a structured jet observed off-axis. This model is able to account for the lack of a peak in the X-ray or optical afterglow in the order of $10^3$~s or later after the prompt emission, as well as the shallow decay rate observed in both energy bands. It also predicts a smooth prompt emission light curve.  In order to continue to further constrain the possible models that describe the origins of XRFs, further prompt observations of these objects are necessary. It is only during the early stages of emission, when the internal mechanisms that power these bursts are probed, that many of these models can be differentiated. It is also important to increase the sample of XRFs with spectroscopic redshifts to gain information on their rest-frame properties. The classification of bursts as X-ray Flashes is based on observer frame properties. However, only knowledge of the rest-frame properties will reveal the underlying differences between XRFs, XRR bursts and GRBs. At a redshift of $\\mathrm{z}=2.44$ XRF~050406 would have a peak energy $E_{\\mathrm{p}}\\sim 52$~keV, at which stage there begins to be an overlap between the intrinsic properties of XRFs that are at redshifts $\\mathrm{z} > 1-2$ and nearby GRBs. Continued prompt and deep observations of XRFs attainable with {\\it Swift} together with an increase in spectroscopic redshifts acquired from rapid follow-up ground-based observations promise to provide the data necessary to further constrain the possible origins of XRFs."
    },
    "0601/astro-ph0601461_arXiv.txt": {
        "abstract": "{The photospheres of about 5--10\\,\\% of the upper main sequence stars exhibit remarkable chemical anomalies. Many of these chemically peculiar (CP) stars have a global magnetic field, the origin of which is still a matter of debate.} {We present a comprehensive statistical investigation of the evolution of magnetic CP stars, aimed at providing constraints to the theories that deal with the origin of the magnetic field in these stars.} {We have collected from the literature data for 150 magnetic CP stars with accurate Hipparcos parallaxes. We have retrieved from the ESO archive 142 FORS1 observations of circularly polarized spectra for 100 stars. From these spectra we have measured the mean longitudinal magnetic field, and discovered 48 new magnetic CP stars (five of which belonging to the rare class of rapidly oscillating Ap stars). We have determined effective temperature and luminosity, then mass and position in the H-R diagram for a final sample of 194 magnetic CP stars.} {We found that magnetic stars with $M > 3 \\,M_\\odot$ are homogeneously distributed along the main sequence. Instead, there are statistical indications that lower mass stars (especially those with $M \\le 2\\,M_\\odot$) tend to concentrate in the centre of the main sequence band.  We show that this inhomogeneous age distribution cannot be attributed to the effects of random errors and small number statistics. Our data suggest also that the surface magnetic flux of CP stars increases with stellar age and mass, and correlates with the rotation period. For stars with $M > 3\\,M_\\odot$, rotation periods decrease with age in a way consistent with the conservation of the angular momentum, while for less massive magnetic CP stars an angular momentum loss cannot be ruled out.} {The mechanism that originates and sustains the magnetic field in the upper main sequence stars may be different in CP stars of different mass.} ",
        "introduction": "\\label{intro} Observations suggest that magnetic fields are ubiquitous in late-type stars, and that a correlation exists between magnetic activity and stellar rotation. The magnetic field of late-type stars is typically localised in spots, and evolves on relatively short time scales. Although not yet fully understood, a \\textit{dynamo} mechanism is commonly invoked to explain the presence of a magnetic field in these kinds of stars.  Early-type stars show a completely different magnetic phenomenology. Magnetic fields appear organised on a large-scale at the stellar surface, and do not change within a time scale shorter than several decades. Instead, a periodic field variability is observed, which is commonly interpreted in terms of the so-called \\textit{Oblique Rotator Model}: the magnetic field geometry is not symmetric about the rotation axis, and the observer sees a magnetic configuration that changes as the star rotates. The field strength (typically a few hundreds up to a few tens of thousands of Gauss) does not seem correlated to the star's rotational velocity. Only a minority (about 5\\,\\%) of early-type stars is magnetic. Practically all known magnetic stars of the upper main sequence are classified between late F- and early B-type, and belong to the category of the so-called \\textit{chemically peculiar} (CP) stars, i.e., stars that exhibit distinctive peculiarities in the element abundance of their atmospheres. Most CP stars (hence most magnetic stars) show also abnormally slow rotation.  The origin of the magnetic fields in CP stars is a matter of debate (Moss \\cite{DM04}). The dynamo hypothesis can hardly explain the observed high field strengths and the lack of a correlation with rotation. A more promising approach is offered by the \\textit{fossil field} theory (Cowling \\cite{C45}; Moss \\cite{DM89}; Braithwaite \\& Spruit \\cite{BS04}), according to which the observed fields in the upper main sequence magnetic stars are the remnants of fields present during earlier stages of stellar evolution. The fact that no correlation is observed with stellar rotation, and the fact that only a small percentage of the upper main sequence stars are magnetic, are naturally explained in terms of variations in the amount of magnetic flux trapped during star formation. However, it not yet clear (observationally nor theoretically) if and how these fields evolve during the main sequence phase.  To provide constraints to the theory of the origin of the magnetic field, it is important to study the \\textit{evolutionary state} of magnetic CP stars. This has been done by several authors, and with conflicting conclusions. Based on a small sample of about 30 magnetic stars, Hubrig et al. (\\cite{HNM00}) suggested that magnetic fields appear at the surface of CP stars with $M\\la 3\\,M_\\odot$ only after they have spent considerable fraction of their life on the main sequence. Most low mass stars studied by Hubrig et al. (\\cite{HNM00}) were drawn from the group of very slowly rotating magnetic stars with detectable Zeeman splitting in their spectra (Mathys et al. \\cite{MHL97}). This selection procedure results in a statistic sample small in size, that may also be intrinsically biased toward older stars (should the star's angular momentum be lost or/and the surface magnetic field increase during the main sequence evolutionary phase). Both Gomez et al.~(\\cite{GLG98}) and P\\\"{o}hnl et al.~(\\cite{PPM05}) investigated much larger samples of CP stars, but instead of selecting objects for which magnetic field was detected, they utilised chemical peculiarity as a proxy of magnetism. They have found that CP stars with chemical anomalies similar to those observed in magnetic CP stars occupy all the regions of the main sequence.  In an attempt to clarify this puzzling situation, we have carried out a new study of the evolution of magnetic stars of the upper main sequence that is based on a very large sample of observed magnetic CP stars, collected through a thoroughfull investigation of the literature and of the ESO archive. Our study includes all known objects for which precise parallaxes were measured by the Hipparcos mission, and for which the presence of magnetic field at the surface could be asserted via direct field detections. We have also investigated whether there are observational evidences for angular momentum losses occurring during the star's evolution in the main sequence.  This paper is organised as follows. Section~\\ref{sample} describes the observations collected from the literature, and presents new magnetic measurements obtained from the analysis of data collected with FORS1 at the ESO VLT. In Sect.~\\ref{teff-lum} we determine the temperature and luminosity of the selected stellar sample, and in Sect.~\\ref{tracks} we derive the position of the selected stars in the H-R diagram. Results and statistical analysis are presented in Sect.~\\ref{results} and discussed in Sect.~\\ref{discuss}. ",
        "conclusions": "\\label{discuss}  We have carried out a detailed statistical investigation of the evolutionary state of the upper main sequence magnetic CP stars. The sample of 194 objects has included all CP stars whose magnetic status could be confirmed with a direct detection of the surface magnetic field and for which precise parallax is available in the Hipparcos catalogue. The literature data on the magnetic observations of CP stars were complemented with the analysis of the archival spectropolarimetric data acquired with the FORS1 instrument at ESO VLT. This allowed us to detect magnetic field in 53 A and B-type CP stars, of which only five were previously known to be magnetic. Using the medium-band (Geneva or Str\\\"omgren) photometry of the program stars, we have determined \\te\\ and then have placed stars on the H-R diagram. We were able to obtain stellar masses, radii, and ages from the comparison of the observed luminosity and temperature with the predictions of the theoretical evolutionary tracks. Our investigation has included a comprehensive non-linear error analysis, that has permitted us to quantify uncertainty of the derived stellar properties for each period of the main sequence life. Numerical simulations were applied to assess the properties of the resulting age distributions of the magnetic CP stars of different masses. We have also employed correlation analysis to study dependence of the surface magnetic field, magnetic flux, and rotation period on the stellar mass and to probe possible evolutionary changes of the stellar rotation and magnetic field.  The key results of our study can be summarised as follows. \\begin{enumerate} \\item The most massive magnetic CP stars ($M>3M_\\odot$) are distributed homogeneously in the main sequence band. On the other hand, stars with $M\\le3M_\\odot$ show the tendency to cluster in the middle of the main sequence. The relative shortage of young and very old stars is especially pronounced in the group of low mass ($M\\le2M_\\odot$) stars. This uneven age distribution cannot be attributed to the effect of random errors in determination of the stellar parameters.\\\\ \\item We have found 22 young ($\\tau\\le0.3$) magnetic stars among the objects with $M\\le3M_\\odot$, thereby rejecting the proposal by Hubrig et al. (\\cite{HNM00}) that all observably magnetic low mass CP stars have completed significant fraction of their main sequence evolution. At the same time, our data for the least massive ($M\\le2M_\\odot$) stars is not inconsistent with a population of stars older than $\\tau=0.3$.\\\\ \\item The average surface field observed in magnetic CP stars decreases with time. However, for stars with $M\\le3M_\\odot$ this decrease is slower than the field weakening computed under the assumption of the magnetic flux conservation. Consequently, we suggest that the surface magnetic flux of the low mass CP stars increases with time.\\\\ \\item Comparison of the average magnetic fluxes of the CP stars from different mass ranges shows that massive CP stars are substantially more magnetic. At the same time, we found a correlation between magnetic field and rotation period for the intermediate mass stars.\\\\ \\item Rotation period of magnetic stars increases with time, but for stars with $M>3M_\\odot$ the stellar structure changes influencing the stellar moment of inertia can fully account for the observed period increase. On the other hand, our results do not rule out the possibility that angular momentum losses occur during the main sequence evolution of stars with $2\\,M_\\odot<M\\le3\\,M_\\odot$. \\end{enumerate}  The conspicuous inhomogeneous age distribution of the low mass magnetic CP stars which has emerged from our statistical analysis requires careful verification and interpretation. Our results raise the question whether the anomalous H-R diagram distribution is the property of stars with significant surface magnetic field or it is intrinsic to the whole group of the SrCrEu and less massive Si-type stars. No significant anomalies in the evolutionary state of these subclasses of CP stars were reported by Gomez et al. (\\cite{GLG98}). However, the adopted upper threshold of the relative parallax uncertainty and the methods used in the Gomez et al. study differ substantially from the analysis procedure employed in the present investigation. Methodologically more similar analysis by P\\\"{o}hnl et al.~(\\cite{PPM05}) included only 15 stars with $M\\le2\\,M_\\odot$, and most of these objects were found in the centre of the main sequence band, very similar to our results.  An indirect hint that the SrCrEu and Si chemical peculiarity is closely related to the presence of magnetic field at the stellar surface comes from the survey of bright Ap stars carried out by Auri\\`{e}re et al. (\\cite{ASW04}). Using sensitive spectropolarimetric field diagnostic methods, these authors were able to detect magnetic field in essentially every peculiar star they have observed and have established a possible lower threshold of $\\approx$\\,250~G for the strength of the dipolar field component. In the light of these findings and taking results of our study into account, an investigation of the evolutionary status of \\textit{all} SrCrEu and Si stars with $M\\le3\\,M_\\odot$ would be of great importance and could benefit from the upcoming revision of the Hipparcos data reduction (van Leeuwen \\& Fantino \\cite{LF05}).  For many decades the problem of the evolution of global magnetic field in A and B stars was approached in the framework of the analysis of the Ohmic decay of the poloidal fossil field, assumed to exist in stellar interior (Cowling \\cite{C45}; Moss \\cite{DM84}; Landstreet \\cite{L87}). Analytical estimates predict appreciable field decay only after  $\\sim$\\,10$^{10}$--10$^{11}$~yr, which exceeds the main sequence lifetime of even the least massive magnetic CP stars. Consequently, any observation of the possible secular evolution of magnetic field was often considered as a challenge for the fossil field theory. However, this classical assessment may be fundamentally flawed due to neglect of the toroidal field component, which must exist in the interior of magnetic CP stars in order to ensure dynamical stability of their global fields (Prendergast \\cite{P56}; Tayler \\cite{T80}). Recent numerical MHD simulations of the fossil field dynamics in radiative stellar interiors by Braithwaite \\& Spruit (\\cite{BS04}, see also Braithwaite \\& Nordlund \\cite{BN05}) have established that instability of the poloidal  field component can be suppressed by the presence of the interior toroidal magnetic component of similar strength. The diffusive evolution of such twisted fields in the simulations by Braithwaite \\& Nordlund (\\cite{BN05}) is determined essentially by the toroidal field. In their model the total magnetic energy decreases with time, but the surface field strength is expected to increase, until the toroidal component emerges at the surface and the field decays rapidly. The time scale of this process is $\\sim$\\,$2\\times10^{9}$~yr for a $2\\,M_\\odot$ star (Braithwaite \\& Nordlund \\cite{BN05}), but the precise value is rather uncertain due to a very schematic treatment of  the atmospheric magnetic reconnection and because of the sensitivity to the assumed initial field structure. These difficulties notwithstanding, the recent numerical work has emphasized limitations of the traditional analytical studies of the global field evolution and suggested that an observation of the long-term systematic change of the surface field is not necessarily incompatible with the fossil field hypothesis. Our observation of the uneven distribution of the low mass magnetic CP stars in the H-R diagram and of the possible increase in their surface magnetic flux with time may be plausibly interpreted as a signature of the Ohmic diffusion of the twisted fossil field. The relative shortage of the old low mass stars may indicate that in these objects we are observing the final stage of the field emergence and rapid decay. If this scenario is correct, the field starts to emerge in stars with $M\\la2\\,M_\\odot$ after $\\sim$\\,$4\\times10^{8}$~yr of the main sequence evolution and completes its decay after $\\sim$\\,10$^{9}$~yr.  The difference in the age distribution and magnetic field properties of the low and high mass CP stars may be related to their pre-main sequence (PMS) evolutionary history and, especially, to the behaviour of the interior and envelope convective zones. One can argue (e.g., Tout et al. \\cite{TWF04}) that a large-scale coherent field typical of magnetic CP stars can be frozen in, or undergo a slow diffusive evolution as envisaged by Braithwaite \\& Nordlund (\\cite{BN05}), in the fully radiative parts of the stellar interiors. In contrast, the turbulence in the convective zones leads to rapid reconnection of the field lines and thus contributes to the dissipation of the magnetic flux. Consequently, primordial field can survive only in stars that do not pass through a long fully convective phase during their approach to the ZAMS and in the subsequent main sequence evolution. Theoretical calculations of the PMS stellar evolution by Palla \\& Stahler (\\cite{PS93}) have shown that, although the low mass stars are fully convective during a large fraction of their PMS life, the duration of this phase decreases rapidly with an increasing stellar mass. A $1.5\\,M_\\odot$ star is expected to spend roughly $10^5$ years in the fully convective phase, which is already a factor of ten shorter compared to a solar-type star and is probably too short to destroy the fossil magnetic field completely. For masses of $\\ga$\\,$2.4\\,M_\\odot$ the PMS star is never fully convective, which facilitates survival of the primordial field. Thus, the observed difference in the age distribution of the low and high mass magnetic CP stars may reflect the history of the field dissipation in the convectively unstable regions of stellar interior. The presence of more extended and long-lived PMS convective zones in stars with $M \\la 2.4-2.0\\,M_\\odot$ suggests that these objects are more likely to possess weak or no field in the outer regions when they reach the ZAMS. Subsequent increase in the field strength results from the outward expansion of magnetic field into the newly formed radiative regions. In contrast to this scenario of the field behaviour in the low mass stars, the fossil field in more massive magnetic CP stars is probably not significantly altered by the feeble PMS convective zones and appears close to the stellar surface even in young stars.  In the context of the study of the origin and evolution of the magnetic field in CP stars it is helpful also to investigate the evolution of the stellar angular momentum. Magnetic CP stars have generally longer rotation periods than normal A and B stars. The bulk of their rotation rates forms a separate Maxwellian distribution with an average value 3--4 times lower than that found in normal A and B-type stars (St\\c{e}pie\\'{n} \\cite{S00}), but there are also groups of CP stars with rotation period of years (Mathys et al. \\cite{MHL97}) or decades (e.g., $\\gamma$ Equ, see Leroy et al. \\cite{LBL94}). Several works suggest that neither field CP stars nor cluster CP stars undergo any significant magnetic braking during their life on the main sequence (North~\\cite{N98}; this study). Therefore, angular momentum must be lost before the star reaches the ZAMS.  St\\c{e}pie\\'{n} (\\cite{S00}) explains the slow rotation as the result of an interaction of the stellar magnetic field with the circumstellar environment during the PMS phase. And if a magnetised wind still persists after the dissipation of the circumstellar disk, a PMS CP star may further slow down, reaching the ZAMS with an extremely long rotation period. In this respect, the very existence of slow rotating stars with strong magnetic fields suggests that magnetic field was already present when the stars were in the pre-main sequence phase. For stars with $M>2\\,M_\\odot$ this hypothesis is supported by the discovery of a magnetic field in NGC\\,2244~334, a $4\\,M_\\odot$ star that has spent only 2\\,\\% of its life in the main sequence (Bagnulo et al. \\cite{BHL04}), and in HD~66318, a star with a mass of $2.1\\,M_\\odot$ and $\\tau = 0.16$, that belongs to the open cluster NGC~2516. More direct confirmation that magnetic fields are present during the pre-main sequence phase comes from the discovery of magnetic field in Herbig Ae/Be stars (considered the progenitors of main sequence early-type stars) by Wade et al. (\\cite{WDB05}).  The most puzzling evolutionary properties are observed for the group of stars with $M \\le 2\\,M_\\odot$, which contains very few young stars and no stars approaching the terminal age main sequence. Observational evidence for lack of young lower mass ($M \\la 2\\,M_\\odot$) Ap stars comes also from open cluster studies. Abt (\\cite{A79}) suggested that low-mass Ap stars are found only in clusters that are at least $10^8$\\,yr old, and a study of four nearby young clusters by P\\\"{o}hnl et al. (\\cite{PMP03}) is also consistent with this finding. Assuming that magnetic breaking during pre-main sequence phase is the reason why these stars rotate slower than normal late A-type stars, then we face a scenario in which magnetic field appears and disappears from the stellar surface several times during the star life. Magnetic field was present at the surface of the star during the pre-main sequence phase, then disappeared when the star reached the ZAMS, to appear again at a more evolved state, and disappear toward the end of the star's life in the main sequence. At the same time, one can argue that magnetic field responsible for the PMS angular momentum loss is not directly related to the fossil field seen at the surfaces of main sequence low mass magnetic CP stars. The former field may be generated by the dynamo processes or represent the outer part of the fossil field tangled by the envelope convection. This complex field decays by the time the star reaches the ZAMS and, after some evolution on the main sequence, the interior fossil field which has retained its global organisation appears at the surface.  Our study is the first investigation which included enough very low mass magnetic CP stars to identify interesting characteristics of this stellar group, yet it is clear that the number of stars in the corresponding mass range is still rather small. This circumstance does not permit us to draw definite conclusions about the nature of the anomalies in the age distribution observed for these stars. Thus, we call for a more detailed analysis of the rotation, magnetic, and evolutionary characteristics of CP stars with $M\\le2\\,M_\\odot$. Such an investigation would represent the most interesting follow up of the present study."
    },
    "0601/astro-ph0601711_arXiv.txt": {
        "abstract": "We describe results of modeling the effects on Earth-like planets of long-duration gamma-ray bursts (GRBs) within a few kiloparsecs.  A primary effect is generation of nitrogen oxide compounds which deplete ozone.  Ozone depletion leads to an increase in solar UVB radiation at the surface, enhancing DNA damage, particularly in marine microorganisms such as phytoplankton.  In addition, we expect increased atmospheric opacity due to buildup of nitrogen dioxide produced by the burst and enhanced precipitation of nitric acid.  We review here previous work on this subject and discuss recent developments, including further discussion of our estimates of the rates of impacting GRBs and the possible role of short-duration bursts. ",
        "introduction": "More attention has been paid in recent decades to the impact of astrophysical events on the Earth and on Earth-like planets which may exist elsewhere.  Impacts by asteroids and comets have been studied and connected with at least one major mass extinction \\cite{alv80}.  Radiation events, such as large solar flares, supernovae, etc. have also been considered.  Gamma-ray bursts are another event in this category and were suggested more than a decade ago \\cite{thor95} as a possible source of danger, and later targeted \\cite{smith04} as a possible accelerating mechanism for evolution by intermittent smaller jolts of radiation.  Knowledge of the effects of a relatively nearby GRB on the Earth are of interest for understanding the history and future of life on Earth, as well as in addressing questions of where and when life may exist elsewhere in the Galaxy and beyond.  GRBs may have played a role in at least one mass extinction in the Earth's history \\cite{mel04}.  Quantifying concepts such as the galactic habitable zone \\cite{gonz01} require information on the rates and effective distances of GRBs throughout the Galaxy.  An important direct impact of a GRB on an Earth-like planet is on atmospheric chemistry due to ionization and dissociation of $\\mathrm{O_2}$ and $\\mathrm{N_2}$, leading to formation of nitrogen oxide compounds. Subsequent effects include destruction of stratospheric ozone which permits solar UVB (290 -- 315 nm) to penetrate more effectively to the ground.  This UVB flux enhances DNA damage, especially in simple life forms such as phytoplankton.  In addition, atmospheric darkening may result due to greatly increased levels of $\\mathrm{NO_2}$.  Precipitation of nitrogen oxide compounds generated by the burst as nitric acid rain may have complicated effects, including even a short-term enhancement of plant productivity.  Each of these effects has been explored in some detail by the authors and collaborators.  Full details may be found in \\cite{mel04, thom05a,thom05b,mel05}.  We summarize here these results and discuss implications and future work. ",
        "conclusions": "It is clear from results reviewed here, combined with the likelihood of a significant GRB event in the last Gy that these events must be considered when looking at the long-term history of life on Earth.  In addition, in investigating probable hazards elsewhere in the Galaxy, GRBs must be taken into consideration, along with supernovae, impactors and other such events.  As we seek to improve our understanding of where and when life may exist beyond the confines of the Earth, these considerations will become increasingly important."
    },
    "0601/astro-ph0601527_arXiv.txt": {
        "abstract": "Irradiation of a stellar atmosphere by an external source (e.g. an AGN) changes its structure and therefore its spectrum. Using a state-of-the-art stellar atmosphere code, we calculate the infrared spectra of such irradiated and transformed stars. We show that the original spectrum of the star, which is dominated by molecular bands, changes dramatically when irradiated even by a low-luminosity AGN ($L_{\\rm X}  = 10^{33}$ erg s$^{-1}$), becoming dominated by atomic lines in absorption. We study the changes in the spectrum of  low-mass carbon- and oxygen-rich giant stars as they are irradiated by a modest AGN, similar to the one at the Galactic center (GC). The resulting spectra are similar to those of the faintest S-cluster stars observed in the GC. The spectrum of a star irradiated by a much brighter AGN, like that powered by a tidally disrupted star, is very different from that of any star currently observed near the GC. For the first time we have discovered that the structure of the atmosphere of an irradiated giant changes dramatically and induces a double inversion layer. We show that irradiation at the current level can explain the observed trend of CO band intensities decreasing as a function of increasing proximity to Sg $A^{*}$. This may indicate that (contrary to previous claims) there is no paucity of old giants in the GC, which coexist simultaneously with young massive stars. ",
        "introduction": "The fact that UV radiation and X-rays can alter the atmospheres of stars has been recognized for more than thirty years \\citep{DavidsonOstriker73,BaskoSunyaev73,Arons73,Basko+74,Fabian79}. Irradiation by a source like an active galactic nucleus (AGN) will produce an increase in the atmospheric temperature and in the mass loss rate \\citep{Edwards80,VoitShull88,ChiuDraine98}. Even a modest level of irradiation from a low-luminosity AGN, like the one currently at the center of the Milky Way, can be sufficient to destroy molecules formed in the atmosphere of cool giant stars, thus transforming their spectrum without inducing significant mass loss. Recently, \\citet{Barman+2004} carried out detailed computations of the atmospheric structure of an M dwarf irradiated by a hot stellar companion (a pre-cataclysmic variable). However, up to now, and despite recent advances in the calculation of molecular opacities \\citep{Jorgensen05} and in stellar modeling algorithms, no detailed computations of the atmosphere of a cool giant star that is irradiated by an external source have been performed.  In this paper, for the first time, we compute the stellar spectrum and atmospheric structure of a cool ($\\sim 4000$ K) giant (both carbon and oxygen rich) star in the presence of an AGN. Our stellar atmosphere code includes a complete frequency dependent description of the atomic and molecular lines that dominate the infrared (IR) spectrum. Our computational approach is general, but we focus here on the spectrum of a star that is irradiated by a source at the Galactic center (GC). We are motivated by the recent discovery \\citep{Revnivtsev+04} that the GC may have been a low-luminosity AGN ($L\\approx 10^{39}$ erg s$^{-1}$) as recently as a couple of hundred years ago. In addition, recent estimates of the rate of stellar tidal disruptions by the GC supermassive black hole (SBH) \\citep{merritt04,merritt05} suggest a rate of order one event per $10^4$ yr or higher for solar-mass stars. Tidal disruption of a star by a SBH is expected to produce an extremely luminous event, $L\\approx 10^{44}$ erg s$^{-1}$, with a duration of weeks or months (e.g. \\citet{Komossa02}).  The GC SBH was recently found to be surrounded by a cluster of apparently young stars. For one of these stars, S0-2, an IR spectrum showed no CO absorption, and possible HeI ($2.1$ $\\mu$m) absorption \\citep{Ghez+2003}. The latter indicates that the star cannot have an effective temperature less than about $15,000$ K \\citep{Hanson+1996}. The orbit of S0-2 has a pericenter distance of $\\sim 100$ AU and an apocenter distance of $\\sim 1000$ AU. The presence of the HeI line and the absence of CO point to S0-2 being young, with spectral class in the domain O and B. The case for hot stars so close to the GC has been further reinforced by \\citet{Eisenhauer+05} who obtained high S/N spectra of 17 S-cluster stars. For the brightest stars in this sample, spectra clearly show the presence of the HeI line at $2.1127 \\mu$m. However for stars with $K > 15$ the line is not present. Hence, the spectral properties of the S-cluster stars are not uniform. If these stars are actually young (a few Myr), their formation so close to GC SBH is a serious theoretical challenge \\citep{Phinney89}. One possibility is that these stars are not young but old, and that their atmospheres have been modified by some physical mechanism: stellar collisions, tidal stripping or external radiation (e.g.  \\citet{genzel+03,alexandermorris03,HM03}, see also the review by \\citet{alexander05} and references therein).  In this work, we compute the detailed spectrum of giant stars irradiated by an AGN in the luminosity range $10^{33} - 10^{44}$ erg s$^{-1}$ integrated over the  range $2-200$ keV (hereafter we denote the luminosity in this range by $L_{\\rm X}$.  In particular, we show how these old giants are affected by low-luminosity AGN ($L_{\\rm X} \\approx 10^{39}$erg s$^{-1}$) like the one that may have been present at the GC in the recent past, and also by the high luminosity ($L_{\\rm X} \\approx 10^{44}$ erg s$^{-1}$) that is believed to accompany the tidal disruption of a star by the SBH. We find that the spectrum of  an irradiated old giant star is similar to that of the faintest S-cluster stars ($K > 15$) observed by \\citet{Eisenhauer+05}. However, we are unable to transform the spectra of cool giants such that they resemble those of the brightest S-cluster stars for a low luminosity AGN ($ L_{\\rm X} < 10^{39}$ erg ${\\rm s^{-1}}$). Higher luminosities though, as the one from a tidally disrupted star are sufficient to heat up the stellar atmosphere above $15,000$ K, and therefore produce the HeI line. However, the total destruction of molecular lines in the spectra of these giants seems to point out that the deficiency of giants in the GC \\citep{Sellgren+1990} might not be such. We show that, within our model, irradiation explains the decreasing CO band intensity observed as a function of radius from Sgr A$^{*}$. The other main result of our study is that we find a double layer inversion in the atmospheres of irradiated giants for irradiation fluxes $f > 10^{2}$ erg s$^{-1}$ cm$^{-2}$ on the surface of the star.  We demonstrate that even if the AGN was shining at or near the Eddington luminosity for $10^6$ years, due to the high eccentricity of the S-cluster stars orbits, the amount of irradiation is most likely not sufficient to make the HeI line appear. On the other hand, stars on more circular orbits could potentially be sufficiently transformed to reproduce the spectral properties of the brightest S-cluster stars (including the HeI line). ",
        "conclusions": "It has been argued in the literature (e.g. \\citet{Ghez+2003}) that the observed spectra of the S-cluster stars, are in agreement with standard spectra of type B8 or earlier, indicating that the stars are young, which is a puzzle because at such distances the tidal force by the central BH is far to great to be overcome by densities in normal molecular clouds. However, effects of the radiation field due to Sg $A^{*}$ on stellar atmospheres has hitherto not been taken into account. The upper layers of stars at the distance of the S-cluster will be strongly affected by this irradiation. We have therefore computed fully self-consistent stellar atmospheres where this irradiation is take into account. The result is a substantial heating of the upper atmosphere. The heating of the upper layers of the atmosphere reduce the intensity of the CO bands as well as all other molecular bands, hereby making even stars of quite late type look fairly much like the observed S-cluster stars.  In particular, in the spectra from our model atmospheres of irradiated giant stars with T$_{\\rm eff} \\approx$ 4000K, the intensity of the CO bands is decreasing when the distance to Sgr A is decreased, in qualitative agreement with the observations by \\citet{Sellgren+1990}. This suggests that contrary to previous claims, there is no dearth of old giants near the galactic center; their molecular signatures have simply been wiped out by the radiation field from Sg$A^{*}$. However, some of the observed S-cluster stars have a strong HeI line in their spectra. We have not been able to reproduce this line for irradiated low-mass stars  for realistic values of the GC luminosity, suggesting that that some other mechanism (perhaps recent star formation of massive star) is responsible for their presence.  Our illuminated atmospheric code is static. In fact we have only been able to obtain converged models for values of $f < 100 f_o$. We have not been able with the present static code to predict the structure and spectrum of a star illuminated with $f > 100 f_o$. For doing this a dynamic model is needed. It is not inconceivable that when dynamics effects are included and models are converged for $f > 100 f_o$, the spectra of these irradiated stars will look even more extreme than the models presented in this work. In particular it will be interesting to investigate if the HeI line can be obtained at higher illuminations for dynamical models."
    },
    "0601/astro-ph0601705_arXiv.txt": {
        "abstract": "{We study the long-term evolution of massive black hole binaries (MBHBs) at the centers of galaxies using detailed full three-body scattering experiments. Stars, drawn from a distribution unbound to the binary, are ejected by the gravitational slingshot. We quantify the effect of secondary slingshots -- stars returning on small impact parameter orbits to have a second super-elastic scattering with the MBHB -- on binary separation. Even in the absence of two-body relaxation or gas dynamical processes, very unequal mass binaries of mass $M=10^7\\,\\msun$ can shrink to the gravitational wave emission regime in less than a Hubble time, and are therefore a target for the planned {\\it Laser Interferometer Space Antenna (LISA)}. Three-body interactions create a subpopulation of hypervelocity stars on nearly radial, corotating orbits, with a spatial distribution that is initially highly flattened in the inspiral plane of the MBHB, but becomes more isotropic with decreasing binary separation.  The mass ejected is $\\gtrsim 0.7$ times the binary reduced mass, and most of the stars are ejected in an initial burst lasting much less than a bulge crossing time. ",
        "introduction": "It is now widely accepted that the formation and evolution of galaxies and massive black holes (MBHs) are strongly linked: MBHs are ubiquitous in the nuclei of nearby galaxies, and a tight correlation is observed between hole mass and the stellar mass of the surrounding spheroid or bulge (e.g. Magorrian et al. 1998; Gebhardt et al. 2000; Ferrarese \\& Merritt 2000; Haring \\& Rix 2004). If MBHs were also common in the past (as implied by the notion that many distant galaxies harbor active nuclei for a short period of their life), and if their host galaxies experience multiple mergers during their lifetime, as dictated by popular cold dark matter (CDM) hierarchical cosmologies, then close MBH binaries (MBHBs) will inevitably form in large numbers during cosmic history. A MBHB model for the observed bending and apparent precession of radio jets from active galactic nuclei was first proposed by Begelman, Blandford, \\& Rees (1980). The coalescence of two spinning black holes in a radio galaxy may cause a sudden reorientation of the jet direction, perhaps leading to the so-called ``winged'' or ``X-type'' radio sources (Merritt \\& Ekers 2002). Recently, observations with the {\\it Chandra} satellite have revealed two active MBHs in the nucleus of NGC 6240 (Komossa et al. 2003), and a MBHB is inferred in the radio core of 3C 66B (Sudou et al. 2003).  BH pairs that are able to coalesce in less than a Hubble time will give origin to the loudest gravitational wave (GW) events in the universe. In particular, a low-frequency space interferometer like the planned {\\it Laser Interferometer Space Antenna (LISA)} is expected to have the sensitivity to detect nearly all MBHBs in the mass range $10^4-10^7\\, \\msun$ that happen to merge at any redshift during the mission operation phase (Sesana et al. 2005). The coalescence rate of such ``{\\it LISA} MBHBs'' depends, however, on the efficiency with which stellar and gas dynamical processes can drive wide pairs to the GW emission stage.  Following the merger of two halo$+$MBH systems of comparable mass (``major mergers''), it is understood that dynamical friction will drag in the satellite halo (and its MBH) toward the center of the more massive progenitor (see, e.g., the recent numerical simulations by Kazantzidis et al. 2005): this will lead to the formation of a bound MBH binary in the violently relaxed core of the newly merged stellar system. As the binary separation decays, the effectiveness of dynamical friction slowly declines because distant stars perturb the binary's center of mass but not its semi-major axis (Begelman, Blandford, \\& Rees 1980). The bound pair then hardens by capturing stars passing  in its immediate vicinity and ejecting them at much higher velocities (gravitational slingshot). It is this phase that is considered the bottleneck of a MBHB's path to coalescence, as there is a finite supply of stars on intersecting orbits and the binary may ``hung up'' before the back-reaction from GW emission becomes important. This has become known as the ``final parsec problem'' (Milosavljevic \\& Merrit 2003).  Recently, in collaboration with Piero Madau and Monica Colpi, the group in Como started a study of MBHB dynamics from two different perspectives, namely, the role of gas in the dynamical friction regime (the 100-to-1 pc decay), and the role of 3--body encounters with bulge stars. The interested reader can find exahustive descriptions of computational tools, and detailed discussions of the results in Dotti, Colpi \\& Haardt (2006), and in Sesana, Haardt \\& Madau (2006).  We have tried to answer a number of questions, i.e.,  i) Is stellar slingshot able to drive MBHBs to the GW regime in an Hubble time? What is the role of different mass ratios and eccentricities in setting the orbital decay time scale?  ii) What are the cinematical properties of the stellar population after the interaction with the MBHB? Hyper velocity stars can be an observable large scale signature of a past MBHB merging.  iii) How do eccentric orbits evolve in a gaseous circum nuclear disk? Do they become circular?  This issue may be relevant in establishing the initial conditions for the braking of the binary due to the slingshot mechanism and/or gaseous gravitational torque.  iv) During the sinking process, do the BHs collect substantial amounts of gas?  This is a query related to the potential activity of a MBH during a merger and its detectability across the entire dynamical evolution. This issue is related to the search of EM counterparts to GW signals.  In the following, I briefly review our main results concerning the hardening in a time evolving stellar background, i.e., I will discuss our answers to questions i) and ii) above. ",
        "conclusions": "We have made a first attempt of modeling the time dependent dynamical evolution of MBH pairs in stellar backgrounds using an hybrid approach, i.e., using results of three-body scattering experiment to follow the time evolution of the stellar distribution. Despite several approximations and limitations, we have obtained a number of results we briefly summarize in the following.  A net outcome of the MBHB-stellar interaction is a significative mass ejection, that we found scaling as $M_{\\rm ej}\\simeq 0.7 \\mu$, nearly independent on the binary eccentricity. This implies that, during every major merger, a stellar mass in the range $0.5-1 M$ is  ejected from the bulge, a possible clue to explain the observed mass deficit in ellipticals (Haenhelt \\& Kauffmann 2002; Ravindranath, Ho \\& Filippenko 2002; Volonteri et al. 2003; Graham 2004). The properties of the ejected stars are well defined. They typically corotate with the MBHB, and show a large degree of flatness, as they are, preferentially, ejected in the binary plane. The ejected sub-population, if detected in nearby galaxies, would be an unambiguous sign of a relatively recent major merger. Though the total ejected stellar mass does not depend on the exact value of the binary eccentricity, the velocity function of the population does. In fact, an highly eccentric binary tends to produce a more pronounced tail of hypervelocity stars.  The relevant quantity to asses coalescence is the ratio $a_f/a_{\\rm GW}$, where the final separation $a_f$ is the orbital separation after loss cone depletion. In general, $a_f$ is reached in a short time, $t \\sim 10^7$ yrs. Though rapid, slingshot is, in general, not sufficient to drive binaries to the gravitational wave emission regime, unless the binary is very massive ($M \\gtrsim 10^8 \\msun$), and/or very eccentric ($e \\gtrsim 0.7$), and/or very unequal mass ($q\\lesssim 0.01$, in fact the gap $a_h-a_{\\rm GW}$ is lower for unequal mass binaries, though the shrinking factor is small). Our conlusion is supported by recent N--body simulations (Makino \\& Funato 2004; Berczik et al. 2005). In terms of {\\it LISA} binaries ($10^4-10^7 \\msun$), it is then important to study in details the possible role of other mechanisms, such as non sphericity of the stellar bulge and presence of circum--nuclear gaseous disks."
    },
    "0601/gr-qc0601118_arXiv.txt": {
        "abstract": "We investigate the effect of deviations from general relativity on approach to the initial singularity by finding exact cosmological solutions to a wide class of fourth-order gravity theories. We present new anisotropic vacuum solutions of modified Kasner type and demonstrate the extent to which they are valid in the presence of non-comoving perfect-fluid matter fields. The infinite series of Mixmaster oscillations seen in general relativity will not occur in these solutions, except in unphysical cases. ",
        "introduction": "There have been numerous studies in the literature on generalisations of the usual Einstein-Hilbert action of general relativity to more complicated functions of the curvature (see refs. in \\cite{Bar83,schmidt}). The motivation for such studies comes from a variety of different sources ranging from attempts to include quantum effects in the gravitational action \\cite{Stelle} to the investigation of phenomena that are presently inadequately explained in the standard model, such as dark energy \\cite% {dark1,dark2}, and violations of the cosmic no hair theorem \\cite{bher,bar}. Of particular interest is the behaviour of these modified theories of gravity in the high-curvature limit, where quantum corrections to general relativity are expected to become important and their influence on the occurrence of singularities and possible bounces of the universe at high curvatures can be studied \\cite{ruz}. It is the investigation of these modified theories of gravity on approach to the initial cosmological singularity that concerns us here. Previous studies have focussed on the behaviour of isotropic cosmologies but, as we know from the situation in general relativity, their behaviour can be misleading. Anisotropies diverge faster than isotropic densities at high curvatures and will generally dominate the behaviour of the cosmology at early (and in some cases even late) times. The most important anisotropic solutions of general relativistic cosmologies are the vacuum Kasner solutions \\cite{kas} and their fluid-filled counterparts that form Type I of the Bianchi classification of three-dimensional homogeneous spaces. These universes are geometrically special, in vacuum (or perfect-fluid) cases they are defined by just one (or two) free constant(s) compared to the four (eight) that specify the most general spatially homogeneous vacuum (or perfect fluid) solutions. However, they have proved to provide an excellent dynamical description of the evolution of the most general models over finite time intervals. The chaotic vacuum Mixmaster universe of Bianchi Type IX undergoes a infinite sequence of chaotic space-time oscillations on approach to its initial or final singularities which is well approximated by a sequence of different Kasner epochs which form a Poincar\\'{e} return mapping for the chaotic dynamical system \\cite% {bkl, misner, jb1, jb2, chern, rend}.  The way in which the Kasner metric has played a central role in the elucidation of the existence and structure of anisotropic cosmological models and their singularities in general relativity makes it the obvious starting point for an extension of that understanding to cosmological solutions of higher-order gravity theories. In an earlier Letter \\cite{Bar06} we have reported the discovery of a new class of exact Kasner-like solutions for gravity theories derived from Lagrangians that are an arbitrary power of the scalar curvature. In this paper we will generalise that study to a much wider class of higher-order Lagrangian theories. We will determine the conditions for the existence of Kasner solutions and find their exact forms. In some cases these solutions are required to be isotropic and correspond to exact Friedmann-Robertson-Walker (FRW) vacuum solutions with zero spatial curvature. By studying gravity theories whose Lagrangians are derived from arbitrary powers of the curvature invariants we are able to find simple exact solutions. Past studies have usually focussed on the addition of higher-order curvature terms to the Einstein-Hilbert Lagrangian. This results in enormous algebraic complexity and exact solutions cannot be found. The results presented here provide a tractable route into understanding the behaviours of anisotropic cosmological models in situations where the higher-order curvature corrections are expected to dominate. The solutions we present are vacuum solutions but we provide a simple analysis which determines when the introduction of perfect fluids with non-coming velocities has a negligible effect on the cosmological evolution at early times. ",
        "conclusions": "We have investigated some anisotropic cosmological solutions to higher-order Lagrangian theories of gravity. Whist the standard general relativistic theory, defined by the Einstein-Hilbert action, appears to be consistent with all weak-field tests, there is less reason to think that it should be valid in high curvature regimes, such as in the vicinity of a possible initial cosmological curvature singularity. In fact, it is in the high-curvature limit that quantum effects should become important and we should expect to see deviations from the standard theory. Without knowing the exact form of such deviations, we have approached this problem by considering a general class of theories that can be derived from an arbitrary analytic function the three curvature invariants $R$, $R_{ab}R^{ab} $ and $R_{abcd}R^{abcd}$. Expanding this function as a power series in these variables we then expect the dominant term in the Lagrangian to be of the form $R^{n}$, $(R_{ab}R^{ab})^{n}$ or $(R_{abcd}R^{abcd})^{n}$ as the singularity is approached. We have found all of the solutions to these theories that can be expressed in terms of the Bianchi type I line--element (% \\ref{kasner}) in vacuum. These solutions provide simple testing grounds for the exploration of quantum cosmological effects in higher-order gravity theory. We have also found the homogeneous, isotropic and spatially flat solutions to these theories in the presence of a perfect fluid.  Exact vacuum anisotropic solutions of the form (\\ref{kasner}) were found for the theories $\\mathcal{L}=R^{n}$ and $\\mathcal{L}=(R_{ab}R^{ab})^{n}$, whilst it was shown that no such solutions exist for theories defined by $% \\mathcal{L}=(R_{abcd}R^{abcd})^{n}$. The properties of these solutions, with respect to their relation to the more general Bianchi type VIII and IX cosmological behaviour, has been investigated. We have argued that for all of the physically relevant new solutions of these theories, the universe will not experience an infinite number of Mixmaster oscillations as the singularity is approached. This is an extension of what was found in the earlier work \\cite{Bar06}. The extent to which these vacuum solutions can be considered as realistic in the presence of a perfect fluid has also been investigated. It has been shown in the velocity-dominated approximation that all the anisotropic vacuum solutions found for plausible theories are valid in the limit $t\\rightarrow 0$. The case of stiff fluids that do not satisfy the velocity-dominated approximation as $t\\rightarrow 0$ has also been investigated. For the solutions to the theories $\\mathcal{L}=R^{n}$ it was found that for a fluid with equation of state $\\gamma <5/3$ the vacuum solutions are good approximations in the vicinity of the singularity. For the theories $\\mathcal{L}=(R_{ab}R^{ab})^{n}$ it was found that the vacuum solutions are good approximations for all $\\gamma <2$.  \\ack  T. Clifton acknowledges a PPARC studentship."
    },
    "0601/astro-ph0601475_arXiv.txt": {
        "abstract": "{} { A model to describe cosmic ray spectra in the energy region from $10^{10}$ to $10^{17}$ eV is suggested based on the assumption that Galactic cosmic ray flux is a mixture of  fluxes accelerated by shocks from nova and supernova  of different types. } { We analyze  recent experimental data on cosmic ray spectra obtained in direct measurements above the atmosphere and data obtained with ground Extensive Air Shower arrays. } { The model of the three classes of cosmic ray sources  is consistent with direct experimental data on cosmic ray elemental spectra and gives a smooth transition from  the all particle spectrum measured in the direct experiments to the all particle spectrum measured with EAS.} {} ",
        "introduction": "It is thought that cosmic rays below $10^{17}$ eV are accelerated by  shocks of supernova explosions, with all components having the same rigidity spectra. However,  recent experimental data show that spectral indices of elemental spectra are different. Therefore  we suggest  that the concepti of a single  population of  sources with the same index of rigidity spectra for all elements is inaccurate We  describe the experimental data using the idea that the measured cosmic ray flux is generated  by three different classes of  sources. It is suggested that each class of sources generates the spectrum that is  power-law by rigidity  with its specific spectral index and maximal rigidity. We  analyze  cosmic ray spectra, measured by  direct methods, and  all-particle spectra, measured in the extensive air showers (EAS). The data used are shown in Table 1.   \\begin{table*} \\caption{Experiments used in this paper}             % \\centering                          % \\begin{tabular}{l l l l l}        % \\hline\\hline                 % Experiment & Technique &Site & Reference \\\\    % \\hline                        %  AMS & Magnetic spectrometer & Spacecraft&  Alcaraz et al., \\cite{ams} \\\\ CAPRICE & Magnetic spectrometer & Balloon&  Boezio et al., \\cite{caprice} \\\\ BESS-TEV& Magnetic spectrometer & Balloon&  Haino et al., \\cite{caprice} \\\\ ATIC-2 & Calorimeter & Balloon & Wefel et al., \\cite{atic}; Panov et al., \\cite{panov} \\\\ SOKOL&Calorimeter & Spacecraft&  Ivanenko et al., \\cite{sokol} \\\\ JACEE& Emulsion chamber & Balloon&  Asakimori et al., \\cite{jacee}; Takahashi et al., \\cite{jacee_nucl} \\\\ MUBEE& Emulsion chamber& Balloon& Zatsepin et al.,\\cite{zatsepin1}; \\cite{mubee}\\\\ RUNJOB& Emulsion chamber& Balloon&  Derbina et al., \\cite{runjob} \\\\ HEAO& \\t Cerenkov counter& Spacecraft&  Engelmann et al., \\cite{heao} \\\\ CRN& Transition radiation & Spacecraft&  M\\\"uller et al., \\cite{crn} \\\\ TRACER& \\t Transition radiation& Balloon&  M\\\"uller et al., \\cite{tracer} \\\\ TIC& Calorimeter & Balloon&  Adams et al., \\cite{tic} \\\\ KASCADE& EAS& Ground based&  Roth et al., \\cite{kascade} \\\\ HEGRA-AIROBIC& EAS+\\t Cerenkov light&Ground based& Arqueros et al.,\\cite{hegra}\\\\ CASA-BLANCA& EAS+ \\t Cerenkov light& Ground based&  Fowler al., \\cite{casa-blanca} \\\\ DICE& \\t Cerenkov light& Ground based&  Kieda et al., \\cite{dice} \\\\ TUNKA& \\t Cerenkov light& Ground based&  Budnev et al., \\cite{tunka} \\\\ \\hline                                   % \\end{tabular} \\end{table*} ",
        "conclusions": "We proceed from the assumption that the difference between proton and helium spectral indices found in the JACEE experiment above 10 TeV (Asakimori et al.\\cite{jacee}), and supported by the preliminary ATIC-2 analysis (Wefel et al., \\cite{atic}) in the energy region above 100 GeV is real. We  suppose  that this difference results from the fact that  cosmic ray fluxes  are a mixture of fluxes from sources with different spectral indices and different maximal energy. We checked a simplest model assuming two classes of sources. We showed that elemental cosmic ray spectra pointed  to a 'kink' at the rigidity of about 50 TV, and connected this with the assumption that one class of sources terminates its effective acceleration  at this rigidity. This source class may be identified with  the supernova explosions into the ISM.  The second source class, presumably supernovae within the local superbubble, accelerates cosmic rays up to rigidity of 4 PV, producing the classic knee. However, the assumption of these two classes of sources is inadequate to fit the energy spectra below 300 GeV per nucleon along with the spectra in the high energy region, even taking into account minimal reacceleration of cosmic rays in the Galaxy. We assume that  the contribution  of nova stars is  essential in the energy region below $\\sim$ 300 GeV per nucleon."
    },
    "0601/astro-ph0601643_arXiv.txt": {
        "abstract": "{We derive single-epoch radial velocities for a sample of 56 B-type stars members of the subgroups Upper Scorpius, Upper Centaurus Lupus and Lower Centaurus Crux of the nearby Sco-Cen OB association. The radial velocity measurements were obtained by means of high-resolution echelle spectra via analysis of individual lines. The internal accuracy obtained in the measurements is estimated to be typically 2-3 km\\,s$^{-1}$, but depends on the projected rotational velocity of the target. Radial velocity measurements taken for 2-3 epochs for the targets HD120307, HD142990 and HD139365 are variable and confirm that they are spectroscopic binaries, as previously identified in the literature. Spectral lines from two stellar components are resolved in the observed spectra of target stars HD133242, HD133955 and HD143018, identifying them as spectroscopic binaries. ",
        "introduction": "The Scorpius-Centaurus association is the nearest association of young OB stars to the Sun. Blaauw (1960, 1964) divided this association into three stellar subgroups: Upper Scorpius (US), Upper Centaurus Lupus (UCL) and Lower Centaurus Crux (LCC). LCC and UCL have roughly similar ages of about $16-20$ Myr, while US is younger with an estimated age of $\\sim$ 5 Myr (Mamajek et al.2002; Sartori et al. 2003). This complex OB association of unbound stars is of great interest because, as recently shown, it is related to the origins of nearby  moving groups of low mass post-T Tauri stars with ages around 10 Myr: the $\\beta$ Pictoris Moving Group, the TW Hydra association, and the $\\eta$ and $\\epsilon$ Chamaleonis groups (Mamajek et al. 2000; Ortega et al. 2000, 2004; Jilinski et al. 2005). In addition, the Scorpius-Centaurus association also appears to be the source of a large bubble of hot gas in which the Sun is plunged. All these structures are believed to have been possibly triggered by supernova explosions taking place in UCL and LCC during the last ~13 Myr (Ma\\'iz-Apell\\'aniz 2001).  The technique adopted for investigating the origins of the $\\beta$ Pictoris Moving Group, for example, consists of tracing back the 3-D stellar orbits of the members of these moving groups until their main first orbits confinement was found, as well as the past mean positions of LCC and UCL. This enabled, investigator not only to determine the dynamical age of this moving group, but also to investigate properties of their birth clouds (Ortega et al. 2002, 2004). It is also possible to find the past positions of the possible supernovae that triggered the formation of these groups by tracing back the orbit of a runaway OB star, which could have been the result of a supernova explosion in LCC or UCL (see, for example, Hoogerwerf et al. 2001 and Vlemmings et al. 2004).  While the past evolution of these moving groups of low mass stars appears to be a relatively simple problem  (as the dynamical ages are not so old), the dynamical evolution of the older and more numerous subgroups LCC and UCL appears to be more difficult. There is the possibility of the presence of several generations of hot stars during the mainstream of the OB association evolution (Garmany 1994). Substructure in LCC and UCL was found by de Bruijne (1999), based on Hipparcos data. The formation of the younger US subgroup could have been triggered by UCL some 6-8 Myr ago (Preibisch et al. 2001). All these studies require reliable radial velocities in order to calculate space velocities. In this paper we present single-epoch radial velocity (RV) measurements for 56 B-type stars members of LCC, UCL and US subgroups, to contribute to studies of their dynamics so as to unravel their origins. ",
        "conclusions": "The cross-correlation technique, which is used for precise RV determinations in later type stars, when applied to the hotter OB stars can be problematic as early type star spectra show few absorption lines. These lines are in many cases, intrinsically broad (up to a few hundreds km\\,s$^{-1}$) due to stellar rotation. In addition, there is also the possibility of line variability affecting their line profiles (Steenbrugge et al. 2003). Therefore, the cross-correlation peak that defines the value of radial velocity can be very broad and contain important sub-structures caused by blending of spectral lines that appear to have different widths. In addition to having high $v \\sin i$ values many OB stars are binary and it is not straightforward to apply the cross-correlation method and to identify them as double-lined binaries; in order to obtain the orbital solution a long set of observations is needed. Detailed cross-correlation technique analyses applied to determinations of radial velocities of early-type stars has been presented in a number of recent publications (see, for example, Verschueren et al. 1997; Verschueren et al. 1999a; Griffin et al. 2000). Griffin et al. (2000), in particular, discuss in detail the difficulties in obtaining accurate RV measurements from cross-correlation in early-type stars spectra.  In this study, having high-resolution observations covering a large spectral range, radial velocity values for the target stars were obtained from measurements of the positions of individual spectral lines of He I, C II, N II, O II, Mg II, Si II and Si III, relative to their rest wavelengths. Radial-velocity standard stars were not observed. (The adopted linelists can be found in Daflon et al. 2001, 2003.) We inspected and identified all unblended lines visible in the spectral range between 3798 \\AA\\ (H\\,{\\sc i}) and 7065 \\AA\\ (He\\,{\\sc i}) in each target star: the number of measurable lines varied between 10 and 74, depending on the star spectral type, rotation velocity, possible multiplicity, but also on the signal-to-noise of the obtained spectra. (We note, however, that for the double lined binary HD 133242, it was possible to measure positions only for 4 lines in component A and 6 lines in component B.) Mean radial velocities using all measurable lines (${\\rm RV}$) and respective dispersions were calculated for the individual target stars.  In Table~1 we assemble our RV results as well as results from the literature. In the two first columns of this table we list the HD numbers of the observed stars with the respective spectral types; in the columns 3 and 4 we list the heliocentric Julian Date (HJD) and the measured radial velocities, plus the number of measured lines in brackets. In the other columns we list results from the literature: columns 5 and 6 list the RV$_{\\rm GCRV}$ and associated error or quality, from the General Catalogue of Radial Velocities (GCRV; quality flags A to E, or I for insufficient data); column 7 presents the projected rotational velocity from Brown \\& Verschueren (1997) and, when not available in this source, the $v\\,sin i$ was taken from the compilation of Glebocki \\& Stawikowski (2000); in column 9 we list, when available, literature references where information about duplicity can be found for the stars. The stars are separated according to the different subgroups in the Scorpius-Centaurus association, following the membership probabilities P(m) listed by de Zeeuw et al. (1999).  The internal precision of our RV determinations can be represented by the scatter obtained from the RV measurements line-by-line, which is listed in column 4 of Table 1. These are typically smaller than $\\sim$2.0 ${\\rm km\\,s}^{-1}$ for stars with estimated $v\\,sin i$ smaller than 100 ${\\rm km\\,s}^{-1}$. We note, however, that when the target $v\\,sin i$ are large, the uncertainties in the derived RVs can be significantly larger due to uncertainties in defining the line center. This can be seen in Figure~2, where we show the obtained line-to-line scatter versus the target projected rotational velocity (as taken from Brown \\& Verschueren 1997 and Glebocki \\& Stawikowski 2000).  In order to evaluate possible systematic effects that line selection could have on the RV results, we selected a  homogeneous set of 28 spectral lines of  H\\,{\\sc i}, He\\,{\\sc i}, Si\\,{\\sc iii} and Mg\\,{\\sc ii}, that could be measured in most of the studied spectra, and recalculated the mean radial velocity values for all possible stars. A comparison of the mean radial velocities ${\\rm RV}$ with ${\\rm RV}_{\\rm 28lines}$ (obtained using only the selected 28 lines) indicates that there are non-significant systematic differences between the two determinations ${\\rm RV} - \\, {\\rm RV}_{\\rm 28lines}= -1.0 {\\rm km\\,s}^{-1}$, with $\\sigma  = 3.1 {\\rm km\\,s}^{-1}$.  \\subsection{Membership}  Table~1 lists the target stars according to their membership the 3 subgroups as assigned by de Zeeuw et al. (1999). Most of the stars in the Lower Centaurus Crux subgroup are flagged as binaries in the literature, except for HD103079, HD106490 and HD108483. For these 3 stars we measured radial velocities of $19.3 {\\rm km\\,s}^{-1}$, $15.3 {\\rm km\\,s}^{-1}$ and $12.8 {\\rm km\\,s}^{-1}$, respectively, with RV$_{mean}$=15.8 $\\pm$3.2 ${\\rm km\\,s}^{-1}$. This mean value is in general agreement with the mean radial velocity calculated by de Zeeuw et al. (1999) for LCC, which is of $12 {\\rm km\\,s}^{-1}$. For the subgroup UCL, our sample has 2 non-binary stars (HD121790 and HD128345) and  RV$_{mean}$=$9.6 {\\rm km\\,s}^{-1}$ which is $\\sim 5 {\\rm km\\,s}^{-1}$ higher than the de Zeeuw et al. (1999) mean value of $4.9 {\\rm km\\,s}^{-1}$. For the Upper Scorpius subgroup all the sample stars have been flagged as binaries in the literature.  For the five target stars that had not been identified as members of any of the three subgroups in the Sco-Cen association (listed as \"others\" in Table~1) and for which we have no information on duplicity, we can attempt to discuss their membership status based on the comparison of the radial velocities measured here and in the literature. We find that the measured radial velocities for HD109026 (RV=4.0 ${\\rm km\\,s}^{-1}$) and HD110335 (RV= 4.8 ${\\rm km\\,s}^{-1}$) are consistent with the mean radial velocity for UCL of 4.9 ${\\rm km\\,s}^{-1}$s. For the target star HD109026 we have RV$_{GCRV}$= 2.5 ${\\rm km\\,s}^{-1}$, therefore it could be considered initially as having constant RV (within the uncertainties) and possibly a member of the Upper Centaurus Lupus subgroup. For HD110335, we find a larger discrepancy between our measurement (RV=4.8 ${\\rm km\\,s}^{-1}$) and the RV value in the GCRV (RV$_{GCRV}$=12.5 ${\\rm km\\,s}^{-1}$). The RVs,  however, are marginally consistent given the expected uncertainty brackets that affect the 2 determinations. If this is really the case, HD110335 could be also considered as a possible member of the UCL subgroup. In addition, the target star HD120640 (RV$_{mean}$=-1.8 ${\\rm km\\,s}^{-1}$ from this study and RV$_{GCRV}$=-4.7 ${\\rm km\\,s}^{-1}$) can be assumed here to have constant radial velocity. These measurements are consistent with the mean RV value of -4.6 ${\\rm km\\,s}^{-1}$ listed by de Zeeuw et al. (1999) for this subgroup.  The two other stars in our sample of 'others' (HD115846 and HD105937) for which we derived RV= -21.8 ${\\rm km\\,s}^{-1}$ and RV=22.7 ${\\rm km\\,s}^{-1}$, respectively, have values in the GCRV of RV$_{GCRV}$=3.0 ${\\rm km\\,s}^{-1}$ and RV$_{GCRV}$=15.0 ${\\rm km\\,s}^{-1}$. We found no information in the literature about these stars  being confirmed binary stars, but the variation in RV for HD115846 exceeds the expected uncertainties: this target probably has a non-constant RV, which prevents further considerations about it belonging to any of the Sco-Cen subgroups. HD105937 has an RV only marginally constant within the uncertainties, but its mean RV is not compatible with any of the subgroups.  \\subsection{Duplicity}  Results from a search for duplicity information for the targets stars (column 9; Table~1) indicate that a large number of stars in our sample are flagged as binaries in the literature. For most of these targets we have only one single-epoch RV measurement and our results alone cannot be used to infer duplicity. However, the RVs derived in this study can be added to RV databases and contribute to long term studies of their orbits. Only a small number of stars had not been previously flagged as binaries in the different studies in the literature. For this subsample of 10 stars, considered a priori as RV constants, it is possible to compare our RV values with the averaged radial velocities assembled in the GCRV. This comparison is shown in Figure~3. Our RV determinations compare favorably with the RVs from the GCRV with a scatter of the order of the estimated uncertainties. ${\\rm RV}_{\\rm GCRV} - {\\rm RV} = 0.7$ and $\\sigma = 4.9 {\\rm km\\,s}^{-1}$. (This was calculated excluding one discrepant star, HD115846, which could be a binary system.) Taking into account the mean precision of RV determinations from the GCRV as $\\pm  5 {\\rm km\\,s}^{-1}$, the external precision of our RV determinations may be evaluated as approximately $\\pm  5 {\\rm km\\,s}^{-1}$.   For those targets with more than one epoch RV measurement in our study, a subsample showed radial velocity variations larger than the expected uncertainties: HD120307, HD142990 and HD139365. Since these had been previously identified as SBs in the literature (Levato et al. 1987 and Batten et al. 1989), our results confirm their duplicity. Four other stars with multiple epoch observations in this study showed a constant RV within the uncertainties: HD116087, HD130807, HD132200 and HD120640.  For 3 targets in our sample (HD133242, HD133955 and HD143018) we were able to separate and identify lines of two stellar components, classifying them as double lined spectroscopic binaries. Their combined spectra showing spectral lines from two stars are shown in Fig~4. Two of these stars (HD143018 and HD133955) were previously identified in Batten et al. (1999) as a spectroscopic binaries."
    },
    "0601/astro-ph0601125_arXiv.txt": {
        "abstract": "We present observations of the early X-ray emission for a sample of 40 gamma--ray bursts (GRBs) obtained using the \\swift\\ satellite for which the narrow-field instruments were pointed at the burst within 10 minutes of the trigger. Using data from the Burst Alert and X-Ray Telescopes, we show that the X-ray light curve can be well described by an exponential that relaxes into a power law, often with flares superimposed. The transition time between the exponential and the power law provides a physically defined timescale for the burst duration. In most bursts the power law breaks to a shallower decay within the first hour, and a late emission ``hump'' is observed which can last for many hours. In other GRBs the hump is weak or absent.  The observed variety in the shape of the early X-ray light curve can be explained as a combination of three components: prompt emission from the central engine; afterglow; and the late hump. In this scenario, afterglow emission begins during or soon after the burst and the observed shape of the X-ray light curve depends on the relative strengths of the emission due to the central engine and that of the afterglow.  There is a strong correlation such that those GRBs with stronger afterglow components have brighter early optical emission. The late emission hump can have a total fluence equivalent to that of the prompt phase. GRBs with the strongest late humps have weak or no X-ray flares. ",
        "introduction": "Gamma-ray bursts (GRBs) are identified as a brief flash of gamma-rays seen at a random location on the sky. For that instant the GRB becomes the intrinsically brightest single object in the Universe.  The duration of a GRB in terms of its prompt emission (i.e., the ``burst'') is usually defined in terms of the timescale over which 90\\% of the gamma-rays were detected -- the T$_{90}$ parameter.  It is conventional to describe those GRBs for which T$_{90}$ $> 2$s as ``long/soft bursts'' and those with shorter duration as ``short/hard bursts'' (e.g. Kouveliotou et al. 1993).  It is now generally accepted that long-duration GRBs result from the death of a rapidly-rotating massive star (Paczy\\'nski, 1998; MacFadyen \\& Woosley, 1999).  The stellar core collapses inwards to form a black hole surrounded by an accreting disk or torus.  The accreting material liberates gravitational potential energy either in the form of neutrinos or via magneto-hydrodynamic processes. These generate a relativistic jet, oriented along the rotation axis of the stellar core, which eventually escapes the star. The jet contains a relatively modest amount of baryonic material moving at high Lorentz factor.  Prior to the \\swift\\ era, almost all of our knowledge about the emission from GRBs beyond a few seconds came from the study of long bursts. Short bursts were, quite literally, too short to be localized by observatories prior to \\swift . The \\swift\\ satellite (Gehrels et al. 2004) has changed that situation. The data for the first short bursts detected by \\swift\\ (Gehrels et al. 2005; Barthelmy et al. 2005b; Hjorth et al. 2005; Bloom et al. 2006) strongly support the idea that the gamma-ray emission seen from short GRBs arises from a jet powered by a merger of two compact objects, most likely two neutron stars or a neutron star and a black hole. The \\swift\\ data also show that the X-ray emission from some short bursts can be detected long after T$_{90}$, allowing for a direct comparison of the early X-ray emission between short and long GRBs. Some short bursts may have collimated flow (e.g. GRB051221A, Burrows et al. 2006) although this is less clear in others (e.g. GRB050724, Barthelmy et al. 2005b).  For both types of GRB, it is thought that the jet flow is inhomogeneous, leading to internal shocks caused by a variable Lorentz factor (Rees \\& M\\'esz\\'aros 1994; Sari \\& Piran 1997). These produce the variable gamma-ray emission seen, when viewing within the jet beam, as a GRB.  The gamma-ray emission typically lasts a few tens of seconds before fading below detectability with the current generation of gamma-ray instruments, including the \\swift\\ Burst Alert Telescope (BAT; Barthelmy et al. 2005a).  The sensitivity of BAT is a complicated function of burst duration and spectral shape (see Band et al. 2006). In practice the BAT detects bursts with 15--150 keV fluences as low as $\\sim 10^{-8}$ erg cm$^{-2}$. The \\swift\\ satellite, however, has the capability to rapidly slew and point its X-ray Telescope (XRT; Burrows et al. 2005b) and Ultraviolet Optical Telescope (UVOT; Roming et al. 2005) at the burst location. The XRT can detect a source at fainter flux levels in the (observed) 0.3--10 keV band than are possible using extrapolated BAT data. For a Crab-like spectrum and a Galactic column of $1\\times 10^{20}$ cm$^{-2}$ the XRT detects $\\sim 1$ count per 100s for a source with a 0.3--10 keV flux of $10^{-12}$ erg cm$^{-2}$ s$^{-1}$. Thus, \\swift\\ can routinely follow the evolution of the earliest X-ray emission from GRBs with only modest gaps in the observed light curve.  The GRB flux can be represented as a function of time and frequency using a function $f_\\nu\\propto \\nu^{-\\beta} t^{-\\alpha}$, where $\\beta$ is the spectral index and $\\alpha$ is the temporal index\\footnote{The photon index $\\Gamma$ is related to $\\beta$ by $\\Gamma = \\beta + 1$.}.  Analysis of some of the first bursts observed with \\swift\\ shows that the early X-ray light curves are complex. In some cases (e.g. Tagliaferri et al. 2005; Hill et al. 2006; Cusumano et al. 2006; Vaughan et al. 2006a) the early X-ray emission (observed with the XRT within a few hundred seconds of the burst) can decline rapidly in the first few minutes, with $\\alpha \\sim 3$ or greater. The light curve can then show a break to a shallower decay component, which we will henceforth refer to as the late emission ``hump''. Analysis of GRB samples from the early phase of the \\swift\\ mission (Nousek et al. 2006) confirm that this pattern of a rapidly decaying light curve followed by a late emission hump is common. However, in some bursts the earliest observed X-ray flux appears to decline relatively slowly ($\\alpha \\sim 1$) (e.g. Campana et al. 2005; Blustin et al. 2006).  The fading X-ray emission could be due to a number of components, including high-latitude emission from the fading burst (Kumar \\& Panaitescu 2000), the interaction of the jet with the surroundings --- the afterglow emission produced by an external shock (e.g. M\\'esz\\'aros \\& Rees 1997), and thermal emission from a photosphere around the outflow (e.g. M\\'esz\\'aros \\& Rees 2000) or from a hot cocoon associated with the jet (e.g. M\\'esz\\'aros \\& Rees 2001). A significant fraction of GRBs also show X-ray flares (e.g. Burrows et al. 2005a) superimposed on the declining light curves.  If the BAT and the XRT are initially detecting the prompt emission from the jet, the most rapidly-decaying X-ray light curves could be due to viewing photons at high-latitudes (i.e. large angles to the line-of-sight) as the prompt emission fades (Kumar \\& Panaitescu 2000; Tagliaferri et al. 2005; Nousek et al. 2006; Zhang et al., 2006; Panaitescu et al. 2006). In at least one case (GRB050219a, Tagliaferri et al. 2005), the BAT and XRT light curves do not appear to join and in other cases the observed X-ray rapid decay rate is higher than expected from the high-latitude model (e.g. Vaughan et al. 2006a). Alternatives such as structured jets (e.g.  Zhang \\& M\\'esz\\'aros 2002), multiple jets or patchy jets (e.g. Burrows et al. 2005c) in which we see varying emission as our light cone expands are also possibilities.  These models have difficulties, however, explaining those bursts that decline relatively slowly for which other emission components, such as the afterglow, may contribute at the earliest times.  Both the rapid-decay and afterglow models have difficulties explaining the late emission hump.  This component may be due to forward shock emission, which is refreshed with energy either due to continued emission from the central engine or because the ejecta have a range in initial Lorentz factor (Rees \\& M\\'esz\\'aros 1998; Sari \\& M\\'esz\\'aros 2000; Zhang \\& M\\'esz\\'aros 2001; Nousek et al. 2006; Zhang et al. 2006; Granot \\& Kumar 2006).  To disentangle the relative contribution of emission from the central engine and that due to afterglow, and hence understand the origin of the early gamma-ray and X-ray emission, requires a systematic analysis of the temporal and spectral properties of a large GRB sample combining data from the BAT and XRT.  The previous studies of GRB samples from \\swift\\ (e.g. Nousek et al. 2006) include a relatively small number of GRBs with early (few minutes from trigger) XRT observations and do not include a combined temporal and spectral analysis of data from the BAT and XRT.  The aim of this paper is to determine the shape of the early X-ray light curve and spectrum for a large sample of GRBs for which observations were obtained early with the \\swift\\ narrow field telescopes.  We use the combined BAT and XRT data to determine whether an extrapolation of the early gamma-ray emission detected by the BAT joins smoothly to that seen by the XRT, whether or not there is always a rapid decline phase seen by either instrument and to investigate the relative importance of emission from the central engine and the afterglow.  Our GRB sample is summarized in section 2. The analyses of the BAT and XRT data are presented in sections 3 and 4 and we explain how the data from the two instruments were combined in section 5. The observed temporal and spectral shapes are presented in section 6. In section 7 we describe a light curve fitting technique which we use to study the various components contributing to the early X-ray emission. The discussion and conclusions are given in sections 8 and 9 respectively. ",
        "conclusions": "We have analysed data for the 40 GRBs observed by \\swift\\ prior to 2005 October 1 for which XRT observations began within 10 minutes of the BAT trigger. We have combined data from the BAT and the XRT to investigate the form of the X-ray emission (0.3--10 keV) during the first few hours following the burst.  The initial XRT spectral index is slightly steeper than that seen in the BAT, showing that spectral evolution occurs as the GRB ages. Combining the BAT and XRT data, the raw light curves show that the initial X-ray emission seen in the XRT is consistent with being a continuation of the emission seen by the BAT.  Some two-thirds of the GRBs display a light curve which shows a steeply declining component that breaks to a shallower decay rate -- the late emission hump -- usually within an hour of the trigger. The remaining bursts decline fairly continuously. At least half of the GRBs in the sample display late X-ray flares, probably due to continued central engine activity, but in only a few GRBs does the fluence in the flares rival that of the burst.  To investigate the early X-ray emission, we used an automatic fitting procedure to align the light curves. Allowing for flares, this procedure works well for the entire sample. The resultant light curves display a prompt phase, mostly observed by the BAT, followed by a decline. The light curve can be described by an exponential which relaxes into a power law whose decay rate varies considerably from burst to burst. The transition time between the exponential and power law provides a well determined measure of the burst duration.  Comparing the temporal and spectral indices of the power law decline, the distribution is consistent with a simple model in which the early emission is a combination of emission from the central engine (parameterised by high latitude emission; Kumar \\& Panaitescu (2000)) and afterglow. Those GRBs in which the afterglow is weak early on decay fast during the power law phase and their X-ray light curves are consistent with the high latitude model. Some are dominated by afterglow, while the majority require a significant contribution from both components. The likelihood of an early optical detection strongly correlates with the strength of the X-ray afterglow component.  The late emission hump component may be present in all objects but can be masked by a strong afterglow component. The late hump can last for many tens of thousands of seconds and may also be due to continued central engine activity (Nousek et al. 2006; Zhang et al. 2006). Interestingly, the strongest humps seen have a total fluence which matches that of the prompt phase.  There appears to be a correlation such that bursts with the most visible humps do not have strong, late X-ray flares.   \\paragraph*"
    },
    "0601/astro-ph0601639_arXiv.txt": {
        "abstract": "We present selected results from integral-field spectroscopy of 48 early-type galaxies observed as part of the {\\tt SAURON} survey. Maps of the H$\\beta$, Fe5015, Mg\\,$b$\\/ and Fe5270 indices in the Lick/IDS system were derived for each of the survey galaxies. The metal line strength maps show generally negative gradients with increasing radius roughly consistent with the morphology of the light profiles. Remarkable deviations from this general trend exist, particularly the Mg\\,$b$\\/ isoindex contours appear to be flatter than the isophotes of the surface brightness for about 40\\% of our galaxies without significant dust features. Generally these galaxies exhibit significant rotation. We infer from this that the fast-rotating component features a higher metallicity and/or an increased Mg/Fe ratio as compared to the galaxy as a whole.  We also use the line strengths maps to compute average values integrated over circular apertures of one effective radius, and derive luminosity weighted ages and metallicities. The lenticular galaxies show a wide range in age and metallicity estimates, while elliptical galaxies tend to occupy regions of older stellar populations. ",
        "introduction": "survey} \\label{secc:sauron}  We are carrying out a survey of the dynamics and stellar populations of 72 representative nearby early-type galaxies and spiral bulges based on measurements of the two-dimensional kinematics and line strengths of stars and gas with {\\tt SAURON}, a custom-built panoramic integral-field spectrograph for the William Herschel Telescope, La Palma (Bacon et al. \\cite{bac01}). The goals and objectives of the {\\tt SAURON} survey are described in de~Zeeuw et al. \\cite{deZ02}, which also presents the definition of the sample. The full maps of the stellar kinematics for the 48 elliptical (E) and lenticular (S0) galaxies are given in Emsellem et al. \\cite{em04}. The morphology and kinematics of the ionised gas emission are presented in Sarzi et al. \\cite{sar05}. Here we summarize selected results of the absorption line strength measurements, which are described more fully in Kuntschner et al. \\cite{kun06}. ",
        "conclusions": "\\label{secc:conclusion}  The results presented here are the outcome of the first comprehensive survey of the line strength distributions of nearby early-type galaxies with an integral-field spectrograph. This data set demonstrates that many nearby early-type galaxies display a significant and varied structure in their line strength properties. This structure can be of subtle nature as seen in the deviations of Mg\\,$b$ isoindex contours compared to the isophotes of galaxies.  The 2D coverage of the line strength allows us to connect the stellar populations with the kinematical structure of the galaxies and thus further our knowledge of the star-formation and assembly history of early-type galaxies."
    },
    "0601/astro-ph0601313_arXiv.txt": {
        "abstract": "We report 0.8$-$4.5~$\\mu$m SpeX spectroscopy of the narrow-line Seyfert 1 galaxy Mrk~1239. The spectrum is outstanding because the nuclear continuum emission in the near-infrared is dominated by a strong bump of emission peaking at 2.2~$\\mu$m, with a strength not reported before in an AGN. A comparison of the Mrk~1239 spectrum to that of Ark~564 allowed us to conclude that the continuum is strongly reddened by E(B-V)=0.54. The excess of emission, confirmed by aperture photometry and additional NIR spectroscopy, follows a simple blackbody curve at T$\\sim$1200~K. This suggest that we may be observing direct evidence of dust heated near to the sublimation temperature, likely produced by the putative torus of the unification model. Although other alternatives are also plausible, the lack of star formation, the strong polarization and low extinction derived for the emission lines support the scenario where the hot dust is located between the narrow line region and the broad line region. ",
        "introduction": "Unified schemes of active galactic nuclei (AGN) invoke the presence of an obscuring dusty torus around the central engine, giving rise to type~1 objects for for pole-on viewing and type~2 objects in edges-on sources. This obscuring structure would also absorbs a significant fraction of the optical/UV/X-ray continuum of the central source and should radiate back this energy at IR wavelengths. In fact, dust reprocessing is regarded as the most likely source of the strong near- and mid-infrared (NIR and MIR, respectively) continuum emission in radio-quiet quasars and Seyfert galaxies. Observational evidence favoring models in which the IR continuum between 1 and 10~$\\mu$m is predominantly or entirely dominated by emission from heated dust is abundant \\citep{em86,ba87,als03}. Indeed, evidence of the presence of hot dust near the sublimation temperature in Seyfert~1 galaxies comes from the observation of a peak of emission with central wavelength between 2.2$-$3.5~$\\mu$m found from JHKL and L' photometry  \\citep{gla92,ma98,ma00} in a few objects. Moreover, very recently, \\citet{ro05} found in the inner 250~pc of Mrk~766, a narrow-line Seyfert~1 galaxy,  this same excess of emission by means of NIR spectroscopy. They were able to determine accurately the form of the NIR bump, confirming that it follows a simple blackbody function at T=1200~K. They found that the host dust emission accounted for up to 28\\% of the total NIR continuum flux in that object.  From above, the growing observational evidence of NIR thermal emission at temperatures close to the sublimation temperature of graphite grains in Seyfert~1 galaxies strongly supports the unified models for AGNs. Meanwhile, additional evidence is needed and many open questions have to be answered. In particular, it is necessary to investigate if the thermal emission could be related to a compact dust/molecular thick torus like the ones in the unified models of \\citet{pk92a} and \\citet{ef95}, for instance, or if it results from emission by hot dust (T$>$900~K) mixed with gas in the NLR/BLR interface region, shielded from the intense UV radiation field \\citep{ma00}.  Here, we contribute to this discussion by presenting the most outstanding evidence of a NIR bump reported to date. It corresponds to the one displayed by Mrk~1239, a compact galaxy, classified as narrow-line Seyfert 1 (NLS1) by \\citet{op85}. Data collected over the past years point out that Mrk~1239 is indeed dusty. \\citet{good89}, for instance, reports that it is one of the three galaxies that display the largest percentage of polarization, both in the line and continuum, in the sample of 18 NLS1 he studied. \\citet{sm04} modeled the polarization nature of this object and found that it was one of the rare cases of Seyfert~1 galaxies that appear to be dominated by scattering in an extended region along the poles of the torus. According to their results, the line-of-sight to the nucleus would pass through the relatively tenous upper layers of the torus, extinguishing the continuum and BLR emission. This radiation, would be polarized further off by the dust located above or below the torus.   X-rays observations also confirm the dusty nature of Mrk~1239. XMM-Newton EPIC PN data suggest two light paths between the continuum source and the observer, one indirect scattered one, which is less absorbed, and a highly absorbed direct light path, in agreement to the wavelength-dependent degree of polarization in the optical/UV. Moreover, Mrk~1239 is classified as a 60~$\\mu$m peaker \\citep{hdr99} because of its ``warm'' far-infrared colours and spectral energy distribution peaking near 60~$\\mu$m. They attributed these properties to dust-obscured active galactic nuclei \\citep{ke94,he95}, the obscuring material likely associated to the putative torus of the unified model.  This letter is organized as follows. In Section~\\ref{obs} we describe the observations, data reduction and resulting spectrum. Section~\\ref{red} determines the internal reddening affecting the nuclear spectrum of Mrk~1239 and compares the observed NIR SED with that of Ark~564. It also analyzes the different components that contribute to the observed continuum. Section~\\ref{dust} examines the hot dust hypothesis for Mrk~1239 in the light of the strong thermal NIR excess of emission detected. Conclusions are in Section~\\ref{final}. Throughout this work we adopt a Hubble constant of H$_o$=75~km\\,s$^{-1}$~Mpc$^{-1}$. ",
        "conclusions": ""
    },
    "0601/astro-ph0601580_arXiv.txt": {
        "abstract": "Hypervelocity stars (HVSs) travel with velocities so extreme that dynamical ejection from a massive black hole is their only suggested origin. Following the discovery of the first HVS by Brown and collaborators, we have undertaken a dedicated survey for more HVSs in the Galactic halo and present here the resulting discovery of two new HVSs:  SDSS J091301.0+305120 and SDSS J091759.5+672238, traveling with Galactic rest-frame velocities at least $+558\\pm12$ and $+638\\pm12$ km s$^{-1}$, respectively.  Assuming the HVSs are B8 main sequence stars, they are at distances $\\sim$75 and $\\sim$55 kpc, respectively, and have travel times from the Galactic Center consistent with their lifetimes.  The existence of two B8 HVSs in our 1900 deg$^2$ survey, combined with the Yu \\& Tremaine HVS rate estimates, is consistent with HVSs drawn from a standard initial mass function but inconsistent with HVS drawn from a truncated mass function like the one in the top-heavy Arches cluster. The travel times of the five currently known HVSs provide no evidence for a burst of HVSs from a major in-fall event at the Galactic Center in the last $\\sim$160 Myr. ",
        "introduction": "All galaxies with bulges probably host massive black holes (MBHs)  in their centers.  \\citet{hills88} first showed that a three-body interaction involving a MBH and a stellar binary can eject one member of the binary with $>$1,000 km s$^{-1}$ velocity.  Hills called stars ejected with $>$1,000 km s$^{-1}$ velocities ``hypervelocity stars.'' Hypervelocity stars (HVSs) are thus a natural consequence of the presence of a massive black hole in a dense stellar environment.  \\citet{brown05} reported the first discovery of a HVS:  a $g'=19.8$ late-B type star, $\\sim$110 kpc distant in the Galactic halo, traveling with a Galactic rest-frame velocity of at least $+709\\pm12$ km s$^{-1}$ (heliocentric radial velocity +853 km s$^{-1}$).  Photometric follow-up revealed that the object is a slowly pulsating B main sequence star \\citep{fuentes06}.  Only interaction with a MBH can plausibly accelerate a 3 $M_{\\sun}$ main sequence B star to such an extreme velocity.  Our HVS discovery inspired a wealth of work from both observers and theorists.  \\citet{edelmann05} report a 8 $M_{\\sun}$ main sequence B star, $\\sim$60 kpc distant, traveling with a Galactic rest-frame velocity of at least $+548$ km s$^{-1}$ that may be a HVS ejected from the LMC.  \\citet{hirsch05} report a helium-rich subluminous O star, $\\sim$20 kpc distant, traveling with a Galactic rest-frame velocity of at least $+717$ km s$^{-1}$ that is probably a HVS ejected from the Galactic Center.  \\citet{holley06} suggest that outliers in the velocity distribution of intracluster planetary nebulae around M87 may be HVSs.  In this paper, we report the discovery of two more HVSs from our ongoing HVS survey.  With HVSs now an observed class of objects, it is important to define true HVSs.  Run-away B stars located many kpc above the Galactic plane have long been known, but their velocities are typically $\\lesssim200$ km s$^{-1}$ and they are very probably bound to the Galaxy.  HVSs, on the other hand, are unbound.  More importantly, the classical supernova ejection \\citep{blaauw61} and dynamical ejection \\citep{poveda67} mechanisms that explain run-away B stars cannot produce ejection velocities which exceed $\\sim$300 km s$^{-1}$ \\citep{leonard93, gualandris05}.  Thus we define a HVS as an unbound star with an extreme velocity that can be explained so far only by dynamical ejection associated with a MBH.  HVSs are important tools for understanding the nature and environs of MBHs: \\citet{holley06} predict that a thin torus of ejected HVSs is the signature of two MBHs forming a tight binary. \\citet{ginsburg06} suggest that the stars on highly eccentric orbits around SgrA$^{*}$ may be the former companion stars to HVSs ejected by the MBH. \\citet{levin05} shows that an intermediate mass black hole (IMBH) on a circular in-spiral into the Galactic Center produces an isotropic burst of HVSs; an IMBH on an eccentric in-spiral produces broad jets of HVSs. \\citet{gualandris05} find that HVSs produced from stellar binary encounters with single MBHs have higher ejection velocities than HVSs from binary MBHs. \\citet{gnedin05} show that the distance and full space motion of HVSs can provide significant constraints on the shape and orientation of the Galactic dark matter halo. \\citet{yu03} expand Hill's original analysis and show that single star encounters with binary MBHs produce $\\sim$10 times more HVSs than stellar binary encounters with single MBHs.  Our paper is organized as follows.  In \\S 2 we describe our survey target selection and observations.  In \\S 3 we present the new HVSs.  We conclude in \\S 4 by discussing what the observed set of HVSs implies about their origin and the nature of the Galactic Center. ",
        "conclusions": "The very existence of the HVSs places interesting limits on the stellar mass function of HVSs, the origin of massive stars in the Galactic Center, and the history of stellar interactions with the MBH.  Interestingly, all five known HVSs are moving with positive radial velocity, consistent with the picture of a Galactic Center origin.  In principle, we can constrain the stellar population from which the HVSs originate by combining predictions of HVS rates with the results of our survey.  \\citet{yu03} predict a HVS rate of $\\sim$10$^{-5}$ yr$^{-1}$ ejected by a single MBH at the Galactic Center.  A HVS moving 600 km s$^{-1}$ travels 120 kpc in 200 Myr, thus the \\citet{yu03} rate implies $\\sim$2000 HVSs of all types to a depth of 120 kpc.  Our $g'_0<19.5$ magnitude-limited survey reaches the same depth $d=120$ kpc for B8 stars and we find two B8 HVSs over our 1900 deg$^2$ region, implying $43\\pm31$ B8 HVSs over the entire sky.  We ask whether this number of B stars is consistent with the number expected from a standard initial mass function (IMF).  A Salpeter IMF \\citep{salpeter55}, integrated over the mass range 0.2-100 $M_{\\sun}$ and normalized to 2000 stars, contains 26 stars between 3 and 5 $M_{\\sun}$.  A Scalo IMF \\citep{scalo86}, similarly calculated, contains 13 stars between 3 and 5 $M_{\\sun}$.  The IMF predictions are systematically lower than the observed frequency implied by the two B stars in our survey, but because of small number statistics there is no significant inconsistency.  In the Galactic center, there is some indication that the IMF is top heavy.  For example, \\citet{stolte05} study the mass-segregated Arches cluster and argue for truncation at masses less than 6 $M_{\\sun}$.  Integrating the \\citet{stolte05} mass function $\\Gamma=-0.86$ over the mass range 3-100 $M_{\\sun}$ and normalizing it to 2000 stars results in 750 3-5 $M_{\\sun}$ stars, an order of magnitude more late-B HVSs than implied by our observations.  Thus our observations already indicate that the HVS parent population is not composed entirely of clusters like Arches or the \\citet{yu03} HVS rate is an underestimate.  HVS travel times from the Galactic Center constrain the history of stellar interactions with the MBH.  Assuming the new HVSs are B stars, travel time estimates for the known HVSs are spread rather uniformly between 30 and 160 Myr.  Thus there is as yet no evidence for a burst of HVSs from major in-fall event at the Galactic Center in the last $\\sim$160 Myr.  The current constraints are less clear if HVS4 or HVS5 are not B stars.  If new discoveries of HVSs continue to show no evidence for coherent bursts of HVSs, theories of in-falling massive star clusters \\citep{gerhard01,kim03} possibly containing an IMBH \\citep{hansen03} to transport B stars near the MBH may not be applicable to the Galaxy.  It is interesting that the northern HVSs all have similar $b\\sim40^{\\arcdeg}$ (see Fig.\\ \\ref{fig:sky}).  An in-spiraling binary black hole produces an unique distribution of HVSs on the sky \\citep{levin05, holley06}.  Detailed theoretical predictions of the distribution of HVSs on the sky ejected by a single MBH at the Galactic Center would be an important testbed for future HVS observations.  HVSs are becoming important tools for understanding MBHs.  Discovering additional HVSs in a well-defined volume will provide better constraints on the origin of HVSs and the nature of the Galactic Center.  Proper motion measurements of HVSs, measured with the {\\it Hubble Space Telescope}, the {\\it Global Astrometric Interferometer for Astrophysics}, or the {\\it Space Interferometry Mission}, may provide significant constraints on the shape of the Galaxy's dark matter potential \\citep{gnedin05}.  The distribution of HVSs in velocity and space may constrain the history of stellar interactions with the MBH.  Identifying HVSs around other galaxies \\citep{holley06} is also an exciting prospect.  We are continuing our radial velocity survey of every late B-type star in the SDSS."
    },
    "0601/astro-ph0601255_arXiv.txt": {
        "abstract": "This review attempts to describe developments in the fields of quasar and quasar host galaxies in the past five. In this time period, the Sloan and 2dF quasar surveys have added several tens of thousands of quasars, with Sloan quasars being found to $z>6$. Obscured, or partially obscured quasars have begun to be found in significant numbers. Black hole mass estimates for quasars, and our confidence in them, have improved significantly, allowing a start on relating quasar properties such as radio jet power to fundamental parameters of the quasar such as black hole mass and accretion rate. Quasar host galaxy studies have allowed us to find and characterize the host galaxies of quasars to $z>2$. Despite these developments, many questions remain unresolved, in particular the origin of the close relationship between black hole mass and galaxy bulge mass/velocity dispersion seen in local galaxies. ",
        "introduction": "\\label{sec:1}  Quasar astronomy has made some significant advances over the past few years. In this review we will focus on recent developments in the field, and in particular on the links between quasars, their black holes and their host galaxies.  The recent discovery that the mostly dormant black holes in the nuclei of nearby galaxies have masses which correlate with the luminosity and velocity dispersion of their host galaxies has directly linked the quasar phenomenon to galaxy evolution. Black hole mass estimates from host galaxy luminosities can be compared to those which use the results of reverberation mapping and the widths of broad emission lines, with generally consistent results (Laor 1998; Onken et al.\\ 2004). This gives us the possibility of understanding the black hole mass -- galaxy mass correlation through a study of the evolution of quasars and their host galaxies.  Reliable black hole mass estimates have also stimulated long-standing debates on how the observational properties of quasars, such as their emission-line spectra, radio-loudness and accretion rates, may (or may not) depend on black hole mass (McLure et al.\\ 1999; Laor 2000; Lacy et al.\\ 2001; Boroson 2002; Woo \\& Urry 2002). The correlation between black hole mass and the mass of the host galaxy have reinvigorated studies on the links between quasar activity and galaxy formation. In particular, the relationship between galaxy mergers, starbursts and AGN has come under increased scrutiny using both results from Sloan Digital Sky Survey (SDSS) (Kauffmann et al.\\ 2003) and {\\em Hubble Space Telescope (HST)} studies of quasar host galaxies (McLure et al.\\ 1999; Canalizo \\& Stockton 2000).  A further recent development, the advent of large, uniform quasar surveys in the optical (the  SDSS quasar survey and Two Degree Field (2dF) quasar survey) and the ROSAT X-ray survey, has also helped progress in the subject. One problem that has bedevilled quasar astronomy throughout its history is that selection of quasar samples has been very prone to selection effects. Optical and (soft) X-ray selection techniques are only sensitive to quasars with little dust or gas in the host galaxy to redden or absorb the quasar light. Radio selection is not sensitive to reddening, but only $\\sim 10$\\% of quasars are bright radio emitters, and the selection effects associated with radio quasar samples are only just beginning to be understood. In particular, the SDSS is able to pick objects that would be missed from traditional quasar surveys (Fan 1999; Richards et al.\\ 2001), and also objects showing only narrow lines in the optical (Zakamska et al.\\ 2003). Further red quasar selection using 2MASS, 2MASS/FIRST, hard X-ray and now Spitzer surveys is finally revealing the population of dust-shrouded type-2 quasars, whose existence has long been suspected, but of which, until recently, there were few known examples.  Detection of the faint host galaxies of bright quasars has been a long-standing problem in observational extragalactic astronomy. Although the host galaxy and the bright, unresolved quasar nucleus typically have similar total fluxes, the diffuse nature of the galaxy emission can frequently be confused with the extended wings of a poorly-characterized point spread function (PSF). Cosmological surface brightness dimming makes detecting the host galaxies of high redshift quasars particularly challenging. Breakthroughs in quasar host studies came with the advent of the {\\em HST}, with its small, stable PSF, and near-infrared array detectors. Early HST studies with the optical instruments on board allowed the detailed study of nearby quasar hosts (Disney et al.\\ 1995; Bahcall et al.\\ 1997). In parallel, ground-based studies in the near-infrared were able to study the quasars at wavelengths where the flux of the quasar was minimized with respect to the flux of the host galaxy (Dunlop et al.\\ 1993). The marriage of $HST$'s small PSF and the near-infrared NICMOS detector allowed routine discoveries of quasar hosts up to $z\\sim 2$ (Ridgway et al.\\ 2001; Kukula et al.\\ 2001).  Adaptive optics are a relatively recent addition to the available techniques for quasar host imaging. Although problematic in some respects (principally PSF variability) AO offers the ability to study larger samples than are practical with the limited observing time available with HST, and, through using 10m-class telescopes, better resolution and surface brightness sensitivity.  We end this review with a discussion of some of the remaining open questions in quasar astronomy, and how advances in telescopes, instrumentation and multiwavelength archives may be able to answer at least some of them in the years to come.    When calculating the intrinsic properties of the quasars we assume a cosmology with $H_0=70 {\\rm kms^{-1}Mpc^{-1}}$, $\\Omega_{\\Lambda}=0.7$ and $\\Omega_{\\rm M}=0.3$. ",
        "conclusions": ""
    },
    "0601/astro-ph0601512_arXiv.txt": {
        "abstract": "{} {The archival spectrum of SDSS\\,J212531.92$-$010745.9 shows not only the typical signature of a PG\\,1159 star, but also indicates the presence of a companion. Our aim was the proof of the binary nature of this object and the determination of its orbital period.} {We performed time-series photometry of SDSS\\,J212531.92$-$010745.9. We observed the object during 10 nights, spread over one month, with the T\\\"ubingen 80\\,cm and the G\\\"ottingen 50\\,cm telescopes. We fitted the observed light curve with a sine and simulated the light curve of this system with the \\texttt{nightfall} program. Furthermore, we compared the spectrum of SDSS\\,J212531.92$-$010745.9 with NLTE models, the results of which also constrain the light curve solution. } {An orbital period of 6.95616(33)\\,h with an amplitude of 0.354(3)\\,mag is derived from our observations. A pulsation period could not be detected. For the PG\\,1159 star we found, as preliminary results from comparison with our NLTE models, $T_{\\rm eff}$\\,$\\sim$\\,90\\,000\\,K, $\\log g$\\,$\\sim$\\,7.60, and the abundance ratio C/He\\,$\\sim$\\,0.05 by number fraction. For the companion we obtained with a mean radius of $0.4\\pm 0.1\\,\\rm R_\\odot$, a mass of $0.4\\pm 0.1\\,\\rm M_\\odot$, and a temperature of 8\\,200\\,K on the irradiated side, good agreement between the observed  light curve and the \\texttt{nightfall} simulation, but we do not regard those values as final. } {} ",
        "introduction": "PG\\,1159 stars are hot hydrogen-deficient (pre-)white dwarfs with effective temperatures between 75\\,000 and 200\\,000\\,K, and $\\log g$\\,=\\,5.5--8.0 (Werner 2001). They are in the transition between the asymptotic giant branch (AGB) and cooling white dwarfs. Spectra of PG\\,1159 stars are dominated by absorption lines of He\\,{\\sc ii}, C\\,{\\sc iv} and O\\,{\\sc vi}.  Current theory suggests (e.g. Werner 2001) that they are the outcome of a late helium-shell flash, a phenomenon that drives the currently observed fast evolutionary rates of three well-known objects (FG~Sge, Sakurai's object, V605 Aql). Flash-induced envelope mixing produces a H-deficient stellar surface. The photospheric composition then essentially reflects that of the region between the H- and He-burning shells in the precursor AGB star. The He-shell flash forces the star back onto the AGB. The subsequent, second post-AGB evolution explains the existence of Wolf-Rayet central stars of planetary nebulae and their successors, the PG\\,1159 stars.  Currently, 37 PG\\,1159 stars are known. Figure\\,\\ref{pg1159stars} shows their position in a log\\,$T_{\\rm eff}$-$\\log g$-diagram. Two of them have been found to be binary stars. These are NGC\\,246 (e.g. Bond \\& Ciardullo 1999), which is a resolved visual binary, and PG\\,2131+066 (Wesemael \\etal 1985). Concerning the latter, it is still unclear whether it is a close binary (Paunzen \\etal 1998) or a resolved visual binary with an M2V star as companion (Reed \\etal 2000). \\begin{figure} \\centering \\includegraphics[width=8cm]{Hl146_f1.ps} \\caption{Positions of the known PG\\,1159 stars in the log $T_{\\rm eff}$-$\\log g$-diagram. The two known binary systems are shown as black dots, the new one is shown as square. Post-AGB evolutionary tracks are taken from Sch\\\"onberner (1983, dashed lines, 0.546\\,M$_\\odot$ and 0.565\\,M$_\\odot$), Bl\\\"ocker (1995, dashed line, 0.605\\,M$_\\odot$), and Wood \\& Faulkner (1986, solid lines) (labels: mass in M$_\\odot$). The dashed-dotted lines represent the theoretical red (Quirion \\etal 2004) and blue edge (Gautschy priv.comm.) of the GW  Vir instability strip.} \\label{pg1159stars} \\end{figure} ",
        "conclusions": "\\begin{enumerate} \\item The spectrum of SDSS\\,J212531.92$-$010745.9 from DR4 of the Sloan Digital Sky Survey shows the signature of a PG\\,1159 star plus emission from a cool irradiated companion. \\item We performed time-series photometry during 10 nights with the T\\\"ubingen 80\\,cm and the G\\\"ottingen 50\\,cm telescopes and detected a period of 6.95616(33)\\,h with an amplitude of 0.354(3)\\,mag. This represents the orbital period of the binary system. Thus, SDSS\\,J212531.92$-$010745.9 is the first close PG\\,1159 binary without any doubts. \\item From a first comparison with NLTE model spectra we derived, as preliminary results, an effective temperature of 90\\,000\\,K, $\\log g\\,\\sim$\\,7.60 and the abundance ratio C/He\\,$\\sim$\\,0.05 for the PG\\,1159 component. A detailed, quantitative NLTE spectral analysis of the PG\\,1159 star and the irradiated companion has to be done next. We will report on the results in a subsequent paper. \\item We simulated the light curve of the binary system with an orbital period of 6.95616\\,h using \\texttt{nightfall}. A good agreement with the observed light curve was obtained for a mean radius of $0.4\\pm 0.1\\,\\rm R_\\odot$, a mass of \\,$0.4\\pm 0.1\\,\\rm M_\\odot$ and a temperature of the irradiated surface of about 8\\,200\\,K for the companion. \\end{enumerate} To determine the system parameters more precisely, high-resolution phase-resolved spectroscopy of SDSS\\,J212531.92$-$010745.9 is necessary. It should then be possible to derive both the companion's variable light contribution to the overall spectrum as well as dynamical masses for both components from radial velocity measurements of their distinct line systems."
    },
    "0601/astro-ph0601348_arXiv.txt": {
        "abstract": "We have simulated 2.5$\\times$10$^3$ seconds of the late evolution of a $23 \\rm M_\\odot$ star with full hydrodynamic behavior. We present the first simulations of a multiple-shell burning epoch, including the concurrent evolution and interaction of an oxygen and carbon burning shell. In addition, we have evolved a 3D model of the oxygen burning shell to sufficiently long times (300 seconds) to begin to assess the adequacy of the 2D approximation. We summarize striking new results: (1) strong interactions occur between active carbon and oxygen burning shells, (2) hydrodynamic wave motions in nonconvective regions, generated at the convective-radiative boundaries, are energetically important in both 2D and 3D with important consequences for compositional mixing, and (3) a spectrum of mixed p- and g-modes are unambiguously identified with corresponding adiabatic waves in these computational domains. We find that 2D convective motions are exaggerated relative to 3D because of vortex instability in 3D. We discuss the implications for supernova progenitor evolution and symmetry breaking in core collapse. ",
        "introduction": "\\par Numerical simulations of stellar evolution are generally based upon restrictive assumptions regarding dynamics (e.g., hydrostatic balance and mixing-length convection), because the dynamic timescales are so much shorter than the nuclear burning timescale. Neutrino cooling accelerates the last burning stages so that direct dynamic simulation is feasible \\citep{bazan1998,asida2000}, at least for the oxygen burning shell. We are extending this work to longer evolutionary times, larger computational domains, and three dimensional flow (3D). In this letter, we summarize new results with a discussion of the hydrodynamics underlying important symmetry breaking and compositional mixing processes which may significantly affect progenitor and core-collapse supernova models. A detailed discussion of these results will appear separately; we indicate the wider implications here. ",
        "conclusions": "\\par {\\em Differences in 2D and 3D.} While these 2D simulations allow us to see the interaction of carbon and oxygen shells, and show wave generation at convective boundaries, they impose an unphysical symmetry on the problem. Our 3D simulations have only been carried to $300$ seconds of stellar time, but show that the wave generation is a robust result. The flow in the convection zones, however, are qualitatively different between the 2D and 3D models: the 2D flow is dominated by vortices which span the convection zone, while the 3D flow is characterized by smaller scale plumes. Quantitatively, the amplitude of the 2D convective motions are larger than the 3D by a factor of 8, while the 3D motions are larger than mixing length values by a factor of 1.5. {\\em Two-dimensional simulations of gravitational collapse will be misleading at least to the extent that convective motions are important.}  \\par{\\em Presupernova Models.} Perhaps the most important impact that internal waves have on stellar structure in the late stages of massive star evolution is the degree to which they drive compositional mixing.  In our simulations internal wave modes grow to non-linear amplitudes and mix material at convective boundaries on a hydrodynamical timescale. Given the strong dependence of presupernova structure on the rate at which mixing occurs at convective boundaries we see the incorporation of internal wave physics into stellar evolution codes as a neccesary refinement.  \\par{\\em Symmetry Breaking.} Spherical symmetry in presupernova models is broken by (1) the density perturbations induced by turbulence within the convection zone, (2) the wave interactions between burning shells, and (3) rotationally induced distortions. The perturbations by waves which are trapped between the oxygen and carbon burning shells are correlated on large angular scales, as is rotation, while the turbulent perturbations have both a smaller scale and amplitude. Our restricted simulation domain filters out wave modes with $l<4$, so it is likely that even larger scale perturbations exist in real stars. Symmetry breaking will seed instabilities in an outward propagating supernova shock \\citep{kuranz2005}, and in the collapsing core. The converging case has been intensely studied for inertial confinement fusion \\citep{lindl98}. The diverging case has implications for the problem of $\\rm^{56}Ni$ and$ \\rm^{56}Co$ decay in SN1987A \\citep{herant1991,kifonidis2003}.  The conclusion that internal wave modes do not grow to large amplitudes during core collapse through nuclear driven overstability \\citep{murphy2004}, is based upon an analysis that ignores (1) the dynamics of convective motion and (2) the shell-shell interactions, both of which are expected to become more violent as collapse is approached. Asymmetries in core collapse have implications for pulsar birth kicks, explosion mechanisms, and for gravitational wave generation. \\citet{char05} have shown that internal gravity waves can transport angular momentum at a rate sufficient to be important in the evolution of solar mass stars; we suggest that they are important for evolution of more massive stars to core collapse, and for plausible prediction of the angular momentum distribution in that collapse."
    },
    "0601/astro-ph0601662_arXiv.txt": {
        "abstract": "{In the present paper  some consequences of  the hypothesis  that the supermassive compact object in the Galaxy centre relates to a class of objects without  event horizon are examined. The possibility of the existence of such objects  was substantiated by the author earlier. It is shown that  accretion of a surrounding gas can cause   nuclear combustion in the surface layer which, as a result of comptonization of the superincumbent hotter layer, may  give a contribution to the observed Sgr A* radiation in the range $10^{15} \\div 10^{20}\\, Hz$. It is found a contribution of the possible proper magnetic moment of the object to the observed synchrotron radiation on the basis of   Boltzmann's equation  for photons  which takes into account the influence of gravity to their motion and frequency. We arrive at the conclusion that the hypothesis of the existence in the Galaxy centre  of the object with such extraordinary gravitational properties  at least does not contradict observations. ",
        "introduction": "An analysis of stars motion in the Galaxy centre gives strong evidence for the existence here of a compact object with a mass of ($3\\div4)\\,10^{6}M_{\\odot}$  associated with the radio source Sgr A* (\\cite{gensel1}; \\cite{ gensel2}; \\cite{ghez2};  \\cite{ ghez1}). There are three kind of an explanation of  observed peculiarities of the object emission:\\\\ --    gas accretion onto the central object - a supermassive black hole (SMBH) ( \\cite{melia};  \\cite{ narayan}) ,\\\\ --     ejection of  magnetised plasma from an environment of the Schwarzschild  radius of the SMBH (\\cite{falcke}; \\cite{melia1}),\\\\ -- explanations, based on  hypotheses of other nature of the central object ( clusters of dark objects (\\cite{maoz})  ,  a fermion ball  (\\cite{viollier}), boson stars (\\cite{torres}; \\cite{yuan1}).\\\\  In the present paper   some consequences of the assumption that   emission of Sgr A*  caused by the existence in the Galaxy  centre   of a supermassive compact object without event horizon  are examined. The existence of such  stable configurations of the degenerated Fermi-gas of  masses $10^{2}\\div10^{10}$ $M_{\\odot}$, radiuses of which are  smaller than the Schwarzschild radius $r_\\mathrm{g}$,  is one of  consequences (\\cite{verozub95}) of   the metric-field equations of gravitation  (\\cite{verozub91}; \\cite{verozub01}). According to the  physical principles founding the equations,    gravity of a compact object  manifests itself as a field  in     Minkowski space-time for a remote observer  in an inertial frame of reference, but become apparent as a space-time curvature   for the observer in  comoving  reference   frames associated with particles freely moving   in this field . If the distance from a point  attractive mass is many larger than $  r_\\mathrm{g}$ , the physical consequences, resulting from these equations, are very close to  General Relativity results. However, they are principally  other at  short distances from the central mass . The spherically- symmetric solution of these equations in  Minkowski space-time   have no  event horizon  and  physical singularity in the centre.  Since these gravitational equations were successfully tested by post-Newtonian effects in  the solar system and by  the binary pulsar PSR 1913+16 (\\cite{verkoch00}), and the stability of the supermassive configurations predicted by them was sufficiently strictly  substantiated  (\\cite{verkoch01}), it is of interest to investigate the possibility of the existence of such  a  kind of an object in the  Galaxy centre as  an alternative to the hypothesis of  a supermassive black hole.  The gravitational force of a point mass $M$ affecting a free- falling particle  of mass $m$ is given by (\\cite{verozub91}). \\begin{equation} F=-m\\left[c^{2} C^{\\prime}/2A +(A^{\\prime}/2A  - 2 C^{\\prime}/2C)  \\overset{\\cdot}{r}^{2}\\right]  , \\label{gravaccel1}% \\end{equation} where \\begin{equation} A=f^{\\prime2}/C,\\ C=1-r_\\mathrm{g}/f,\\ f=(r_\\mathrm{g}^{3}+r^{3})^{1/3}. \\label{ABC} \\end{equation} In this equation $r$ is the radial distance from the centre, $r_\\mathrm{g} =2GM/c^{2}$,  $G$ is the gravitational constant, $c$ is the  speed of light at infinity, the prime denotes the derivative with respect to $r$.  For   particles at rest ($\\overset{\\cdot}{r} =0$ ) \\begin{equation} F=-\\frac{GmM}{r^{2}} \\left[ 1-\\frac{r_\\mathrm{g}}{(r^{3}+r_\\mathrm{g}^{3})^{1/3}} \\right] \\label{ForceStat}% \\end{equation}  Fig. \\ref{fig: GRForce} shows the force $F$ affecting    particles at rest  and  particles, free falling from infinity with zero initial speed,   as the function of the distance $\\overline{r}=r/r_\\mathrm{g}$ from the centre.   \\begin{figure}[h] \\resizebox{\\hsize}{!} {\\includegraphics[]{f.eps}} \\caption{The gravitational force (arbitrary units) affecting  free-falling  particles  (curve 1) and   particles  at rest (curve 2) near a point  attractive mass $M$.} \\label{fig: GRForce} \\end{figure}  It follows from  Fig. \\ref{fig: GRForce}  that the gravitational force affecting free falling particles changes its sign at  $r \\approx 2  r_\\mathrm{g}$ . Although we still never observed particles motion  at  distances of the order of $r_\\mathrm{g}$,   we can verify this result for very remote objects in the Universe -- at  large cosmological redshifts . Indeed, consider a simple model -  homogeneous  selfgravitating  dust-balls of  different-size  radiuses $r_\\mathrm{b}$. The force  acting on  particles at the  surface of such a ball     is given by Eq. (\\ref{gravaccel1})  where, in this case,  $r_\\mathrm{g}=(8/3) \\pi c^{-2} G \\varrho_{b}  r_\\mathrm{b}^{3}$ is the  Schwarzschild radius of  the ball and $\\varrho_{b}$ is  its density. If  the ball density  is of the order of the observed density of  matter in the Universe ($\\sim 10^{-29}  \\mathrm{g} \\, \\mathrm{cm}^{-3}$), and the its radius   $r_\\mathrm{b}$   is not less than the radius of the  observed matter  ( $\\sim  10^{28} \\mathrm{cm}$    ), then $r_\\mathrm{b} \\leq r_\\mathrm{g}$. It follows from       Fig. \\ref{fig: GRForce}  that under the sircumstances the radial accelerations   $F/m$   of  speckle particles  on the ball surface are positive.  The repulsive force give rise a deceleration of  the  expansion of such a selfgravitating   dust-like mass (  \\cite{verozubA02} ). More accurate calculations at the redshift  $0<z \\leq 1$       (\\cite{verkoch3}) give a good accordance with  observation data  (\\cite{riess})     .  We assume here for calculations that the mass $M$ of the central object is about  $3\\,10^{6} M_{ \\odot}$. The  radius  $R$  of the object can be found   from the distribution of the density $\\rho_\\mathrm{m}$ of  matter inside the object. It can  be obtained by   solving  the  equations of the  hydrodynamical equilibrium: \\begin{eqnarray} \\label{EquilibrHydEq} dp_\\mathrm{m}/dr=\\rho_\\mathrm{m} g_\\mathrm{m},  \\,\\, d M_\\mathrm{r} /dr=4 \\pi r^2 \\rho _\\mathrm{m},\\, \\,   \\\\ p_\\mathrm{m}=p_\\mathrm{m}(\\rho_\\mathrm{m}).  \\nonumber \\end{eqnarray} In these equations  $M_\\mathrm{r}$ is the  mass  of matter  inside  the sphere $\\mathcal{O}_\\mathrm{r}$  of the radius $r$, $g_\\mathrm{m}=F/m$  is the force   (\\ref{ForceStat}) per unit of mass,  $r_\\mathrm{g}=2 G \\, M_\\mathrm{r}  /c^2$ is the Schwarzschild radius of  the sphere $\\mathcal{O}_\\mathrm{r}$  and $p=p(\\rho_\\mathrm{m})$ is the equation of state of   matter  inside  the object. The equation of state,  which is valid for  the  density range  of $(8 \\div 10^{16} ) \\,   \\mathrm{g}   \\,   \\mathrm{cm}^{-3}$, is given by (\\cite{harrison}) \\begin{equation} p_\\mathrm{m}= \\left(   \\frac{n_\\mathrm{m}}{\\partial n_\\mathrm{m} /  \\partial \\rho_\\mathrm{m} }    - \\rho_\\mathrm{m}  \\right)  \\, c^{2}, \\end{equation} where in CGS units \\begin{equation} n_\\mathrm{m} = Q_{1} \\rho_\\mathrm{m} \\left(    1+Q_{2} \\rho_\\mathrm{m}^{9/16}\\right) ^{-4/9}, \\end{equation} $Q_{1}=6.0228 \\,10^{23}$  and  $Q_{2}=7.7483 \\,10^{-10}$ . For the  gravitational force under consideration there are two types of the solutions of  Eqs. ( \\ref{EquilibrHydEq} ).  Besides the solutions describing white dwarfs and newtron stars, there are     solutions  with very large masses. For the mass of $ 3\\,10^{6} M_{\\odot}$ the object radius resulting from Eqs.  (\\ref{EquilibrHydEq} )   is about $0.4 R_{ \\odot}=0.034\\,  r_\\mathrm{g}=3 \\,10^{10} \\mathrm{cm}   $ where $r_\\mathrm{g}$ is the Schwarzschild radius  of the object.  It seems  at first glance   that  accretion onto the object with a solid surface must lead to too  large  energy release  which contradicts a comparatively low bolometric luminosity ( $\\sim 10^{36}$ $\\mathrm{erg}\\ \\mathrm{s}^{-1}$ ) of Sgr A*. However, it must be taken into account that the radial velocity of test particles free falling from infinity   is given by the equation (\\cite{verozub91}) \\begin{equation} v_\\mathrm{ff}=c\\left[ \\dfrac{C}{A} (1-C)    \\right]^{1/2}. \\label{vff} \\end{equation} The velocity   fast decreases at $r < r_\\mathrm{g}$. As a result, the velocity  at the surface is only about $3.3 \\,10^{8}\\,\\mathrm{cm}\\, \\mathrm{s}^{-1}$ . Consequently, at the accretion rate $\\overset{\\cdot}{M}=10^{-7}$ $M_{\\odot}\\ \\mathrm{yr}^{-1}$ the amount of the released energy $\\overset{\\cdot}{M}v^{2}/2$  is  equal to $ 3.4 \\,10^{35}\\,   \\mathrm{erg}\\, \\mathrm{s}^{-1}$. ",
        "conclusions": "The clarification  of  nature of compact objects in the galactic centres is one of actual problems of fundamental physics and astrophysics. The results obtained above, of course, yet do not allow to draw the well-defined conclusions. However they show that the possibility, investigated here, at last,does not contradict the existing observations, and require  further study."
    },
    "0601/astro-ph0601454_arXiv.txt": {
        "abstract": "We re-assess the question of a systematic time delay between the formation of the progenitor and its explosion in a type Ia supernova (SN Ia) using the \\emph{Hubble} Higher-z Supernova Search sample \\citep{str04}. While the previous analysis indicated a significant time delay, with a most likely value of 3.4~Gyr, effectively ruling out all previously proposed progenitor models, our analysis shows that the time-delay estimate is dominated by systematic errors, in particular due to uncertainties in the star-formation history. We find that none of the popular progenitor models under consideration can be ruled out with any significant degree of confidence. The inferred time delay is mainly determined by the peak in the assumed star-formation history.  We show that, even with a much larger supernova sample, the time delay distribution cannot be reliably reconstructed without better constraints on the star-formation history. ",
        "introduction": "Type Ia supernovae have been used extensively as standard distance indicators and have provided the best evidence to date for an acceleration of the Universe \\citep{rie98, per99, rie04}. Future missions, e.g. \\emph{GAIA} and \\emph{SNAP}, will greatly increase the number of detected SNe Ia and significantly reduce the statistical errors in the determination of cosmological parameters. However, the nature of the progenitors of type Ia supernovae is still unknown and the empirically calibrated \\emph{Phillips relation} \\citep{phi93} is not fully understood physically.  Several progenitor scenarios are under discussion, but there is no consensus due to uncertainties in the evolutionary processes \\citep{HKN96, HKN99, LV97, lan00, HP04} and the explosion mechanism \\citep{HN00, roe05, gam05}. One of the signatures of the various scenarios is the distribution of time delays between the formation of the progenitor systems and their explosion, which could give rise to a significant difference between the redshift dependence of the supernova rate (SNR) and the star-formation history (SFH).  \\citet{str04}, hereafter S04, aimed to detect this difference by studying the distribution of 25 high-z SNe Ia in the \\emph{Hubble} Higher-z Supernova Search sample \\citep{rie04}. Their approach was to infer the mean time delay of the distribution using a Bayesian analysis, which assumed different parametrized time-delay distributions and adopted the star-formation history (SFH) from \\citet{gia04}, hereafter G04. They concluded that mean time delays shorter than $\\sim 2$~Gyr ought to be excluded with a 95 per cent confidence level, ruling out essentially all progenitor scenarios currently under discussion. In a recent re-assessment of the constraints, \\citet{str04erratum} obtained a 95 per cent lower limit ranging from 0.2 to 1.6~Gyr for different time-delay distributions. In the corrected best-fitting model, the 95 per cent confidence interval ranged from 1 to 4.4~Gyr with the most likely value at 3.4~Gyr.  Unlike core collapse SNe (SNe II, Ib/c) that originate from massive progenitors with relatively short main-sequence (MS) lifetimes ($\\sim 3-20$~Myr), SNe Ia are believed to be thermonuclear explosions of white dwarf stars (WDs) whose progenitors have MS lifetimes ranging from $\\sim 30$~Myr to several billion years. This implies a minimum time delay for SNe Ia of the order of $\\sim 30$~Myr.  Most of the SN Ia progenitor scenarios that have been proposed involve mass transfer on to a CO WD in a binary system, either through the expansion and Roche lobe overflow of an evolved companion (\\emph{single degenerate [SD] scenarios}) or through the slow release of gravitational waves, orbital shrinking, Roche lobe overflow and merging of a compact double WD system (\\emph{double-degenerate [DD] scenarios}). Both scenarios have associated time-delay distributions that have been estimated with binary population synthesis codes (BPS), where the properties of binary systems are followed from their birth up to the explosion stage through the many different evolutionary paths. Independently of the particular treatment of the binary interactions, the resulting time-delay distributions differ considerably since their characteristic time-scales have different origins.  The SD scenario is controlled by the process of mass accretion, which has to occur at just the correct critical rate in order to allow the growth of the mass of the companion WD up to the Chandrasekhar limit \\citep{nom91}. The dominant evolutionary path seems to occur via the accretion of matter on to a CO WD from a slightly evolved MS star, the CO WD + MS -- SD scenario \\citep{vH92, rap94, LV97, lan00, HP04}. In this channel, the accretion rate is determined mainly by the mass of the donor star, which must lie in a narrow range in order to satisfy the required accretion-rate constraints. As a consequence, the distribution of MS lifetimes and the time-delay distribution of the channel are relatively narrow, peaking at $\\sim 670$~Myr and rapidly becoming negligible after $\\sim 1.5$~Gyr.  Although recent simulations have suggested that other evolutionary paths within the SD framework are of minor importance \\citep{HP04}, it is quite possible that their contribution has been underestimated. This is particularly important for the CO WD + RG -- SD scenario, where a red-giant (RG) star accretes matter on to a CO WD star \\citep{HKN96}; in this channel the time-delay distribution extends up to several Gyr.  The DD scenario \\citep{IT84, web84}, in contrast, is controlled by the time that it takes for the binary system to coalesce, which depends roughly on the fourth power of the separation of the double-degenerate system \\citep{sha83}. As a result, the time-delay distribution can be described by a low time-delay cutoff ($\\sim 30-100$~Myr) and an approximately power-law decline up to the age of the Universe. The lower time-delay cutoff can be explained by the time required to form the most massive degenerate systems with the shortest MS lifetimes, whereas the power-law tail can be explained by the power-law relation between coalescence time and separation of the double-degenerate systems.  However, the expected accretion rates in the DD scenario are a problem: they are so high that present calculations suggests that this leads to accretion-induced collapse (AIC) and the formation of a compact object rather than a thermonuclear explosion \\citep{NI85, SN85, SN98, TWT94, nom91}.  Therefore, the currently generally most favoured progenitor scenario is the SD scenario. Because the seemingly dominant evolutionary path of this channel would need to be discarded if the mean time-delay were found to be higher than 2~Gyr, it is important to confirm the significance of the S04 results.  In this work we have studied the SN Ia time-delay distribution using the sample of S04 and the same basic analysis, but introducing alternative SFHs found in the literature, avoiding binning effects as much as possible and using a Goodness of Fit (GoF) test that is generally recommended for small samples. We discuss the data and analysis in Section \\ref{sec:analysis}, show the results and Monte Carlo simulations in Sections \\ref{sec:results} and \\ref{sec:Monte Carlo} and discuss their significance in Sections \\ref{sec:discussion} and \\ref{sec:conclusions}. Throughout his paper, we adopted a value for the Hubble constant of $H_{\\rm 0} = 70$ km s$^{-1}$ Mpc$^{-1}$, present ratios of matter, curvature and dark energy density over the critical density of $\\Omega_M = 0.3$, $\\Omega_K = 0$ and $\\Omega_{\\Lambda} = 0.7$, respectively, and a `dark energy' pressure over density ratio ('equation of state') of $w =-1$. ",
        "conclusions": "\\label{sec:conclusions}  We have found that systematic errors associated with the use of alternative star-formation histories are comparable if not larger than the statistical errors reported in \\citet{str04}.  The position of the peak of the SFH was found to be the crucial parameter for the recovered time delays: Later peaked SFHs result in lower time delays and vice versa. Furthermore, the confidence intervals for the time delays depend on the functional form of the delay distributions assumed in the analysis. The use of wider time-delay distributions, in particular the e-folding model, gives considerably longer upper limits for the time delays.  For the data set under investigation, we found that the KS test is better suited to obtain confidence intervals than a Bayesian analysis. The KS test confidence intervals are unambiguously defined, and we have confirmed their validity using Monte Carlo simulations. In contrast, the skewness of the Bayesian probabilities and the small sample size result in somewhat arbitrary confidence intervals.  A KS test using the shape of the theoretical time-delay distributions shows that the extinction corrected model from G04 is incompatible with the CO WD + MS -- SD scenario, although other SFHs are compatible. The DD scenario cannot be rejected at 95 per cent confidence with any combination of SFH and time-delay distribution. Starting from theoretical time-delay distributions, they consistently favour the SFH from CE01.  If we wish to constrain the time-delay distribution and possibly discard progenitor scenarios from the redshift distribution of supernovae, it is of foremost importance to determine the cosmic star-formation history more accurately. Otherwise, the uncertainty in the SFH will continue to limit the interpretation of SN data sets of any conceivable size.  Secondly, we would need a better understanding of our progenitor models: if we ignore factors affecting efficiencies, their cosmic evolution will make the delay time distributions evolve and produce a situation where it is hard to disentangle these different effects without qualitatively different observations.  Only after having solved these issues would deeper and wider surveys help to constrain the delay times and progenitor models by decreasing statistical noise and reducing the influence of environmental cosmic variance on the supernova samples.  Different observations have already produced some evidence for shorter time delays: Recent work by \\citet{man05} showed that the most efficient host galaxies for SNe Ia production among all galaxies are irregulars, and \\citet{dv05} showed that among elliptical galaxies it is in particular the radio-loud galaxies, which are also believed to be associated with recent star formation. Also, \\citep{gar05} have shown that SNe Ia of normal luminosity occur particularly in those elliptical galaxies with quite substantial star-formation rates. Also, SNe Ia in galaxy clusters seem to indicate time delays that are shorter than 2~Gyr \\citep{MGY04}.  Finally, it is clear that, if the SNe Ia phenomenon is composed of several production channels, all conclusions we have drawn apply to the dominant channel. A generally less common channel with different characteristics could again be the dominant channel in a subset of galaxies with older-age or higher-metallicity stellar populations. It is already established that under-luminous 91bg-type SNe Ia preferably occur in non-star-forming hosts, such as ellipticals. It remains to be explored theoretically, whether the CO WD + RG -- SD scenario could be related specifically to old populations. We should also consider the possibility that we may not even have found the dominant progenitor channel for normal SNe Ia \\citep{ham03, tou05}.  Even if the SFH peaks at redshift $\\sim 1$ and the recovered time delays are consequently low, the associated SNR (see Fig.~\\ref{fig:SNR}, lower panel) tends to be over-estimated in the highest-z bin. This effect could be interpreted as the signature of a long time delay component that does not contribute to the total SNR when the universe is too young. However, because the highest-z bin only contains two SNe, this does not lead to low non-rejection probabilities. A study of bimodal time delay distributions could only be done with this method if the uncertainties in the SFH were significantly reduced and the metallicity cutoff on the efficiencies was properly quantified.  If different channels produce SNe Ia from progenitors of distinctly different age or metallicity, then an increase of the supernova sample could greatly help the identification of the various plausible progenitors.  However, such a data set would most be beneficial if it is complemented with host galaxy characterisation \\citep[see][]{vdb05} and spectra with better signal \\citep[see][]{ben05}."
    },
    "0601/hep-ph0601161_arXiv.txt": {
        "abstract": "Generically, universal extra dimension (UED) extensions of the standard model predict the stability of the lightest Kaluza-Klein (KK) particle and hence provide a dark matter candidate. For UED scenarios with one extra dimension, we model-independently determine the size of the induced dimension-five magnetic dipole moment of the KK-neutrino, $\\nu^{(1)}$. We show that current observational bounds on the interactions of dipole dark matter place constraints on UED models with KK-neutrino dark matter. ",
        "introduction": "While the standard model proves an excellent framework for fundamental interactions at energy scales up to at least the sub TeV range, it nevertheless leaves a number of fundamental problems. Theoretically, one of the most outstanding puzzles centers on the origin and mechanism of electroweak symmetry breaking, and the quantum mechanical stability of the hierarchy generated between the electroweak scale and the Planck scale. In addition, recent astrophysical observations \\cite{DarkMatter} concord with $0.094 < \\Omega_{CDM}h^2 < 0.129$, indicating the presence of cold non-baryonic dark matter as the principle form of matter in the Universe, of which the standard model provides no explanation. The most popular candidate for dark matter assumes a non-standard model, stable, electrically neutral, and weakly interacting particle -- the WIMP hypothesis. Clearly, from both a theoretical and phenomenological perspective, the standard model requires extension. In this letter we wish to explore the consequences of the universal extra dimension (UED) \\cite{Appelquist:2000nn} extension of the standard model.  In the UED scenario, all standard model particles can freely propagate in the bulk of one or more extra dimensions and thus each standard model particle is associated with a Kaluza-Klein (KK) tower of states. Each state in the KK tower has the same spin as its standard model counterpart. An important consequence of UED models concerns the existence of a conserved discrete symmetry, KK-parity, which guarantees the stability of the lightest KK particle (LKP) and thus provides a dark matter candidate. Suitable thermal relic dark matter candidates that have been studied extensively \\cite{Servant:2002aq} include the first KK-excitations of the hypercharge boson, the photon, and the neutrino \\ie $B^{(1)}$, $\\gamma^{(1)}$, or $\\nu^{(1)}$.  The tree-level mass spectrum of the KK-excitations of UED models reveals a nearly degenerate spectrum. As an example, a UED model with one extra-dimension compactified on an $S_1/Z_2$ orbifold of radius $R$, leads to the tree level mass relation \\be\\label{mass} m^{(n)} = \\sqrt{(n/R)^2 + (m^{(0)})^2} \\ee for the n-th KK mode, where $m^{(0)}$ constitutes the zero-mode mass (\\ie the standard model particle value). Quantum corrections typically dominate over zero mode level contributions and therefore the resulting mass spectrum depends crucially on radiative effects. In general, a moderately split UED mass spectrum \\cite{Cheng:2002ej} develops. In the minimal UED model (MUED) \\cite{Cheng:2002ej}, one-loop calculations suggest that the LKP is well approximated by the KK hypercharge boson, $B^{(1)}$, and numerous studies have examined the thermal production and prospects of direct and indirect detection of $B^{(1)}$ and $\\gamma^{(1)}$ LKP dark matter \\cite{Servant:2002aq,KKDMcollection}.  However, the non-renormalizability of UED models imply the existence of an ultraviolet cut-off, typically of the order of a few tens of TeV, at which point the model requires UV completion. As such, UED models must be regarded as an effective theory. The presence of incalculable boundary terms arising from the UV complete theory can potentially change the mass spectrum, resulting in different LKP candidates. Recent studies of the relic density of $B^{(1)}$ LKP dark matter with the full MUED spectrum \\cite{coannmat,coannkribs} reveal substantial observational tension with constraints from electroweak precision data \\cite{Flacke:2005hb}. Thus, non-minimal models with brane-localized terms appear as a likely alternative if UED models are to provide a successful phenomenology. Furthermore, model independent studies \\cite{Servant:2002aq} show that $B^{(1)}$, $\\gamma^{(1)}$, and $\\nu^{(1)}$ can all be thermally produced with abundances sufficient to provide the dark matter. Constraints on minimal UED models from limits on weak neutral current nucleon-$\\nu^{(1)}$ elastic scattering in direct searches, together with thermal dark matter production mechanisms, disfavour $\\nu^{(1)}$ dark matter \\cite{Servant:2002hb}. However, given the need for possible new non-minimal interactions, we consider further consequences of KK-neutrino dark matter model-independently.  While there exists compelling evidence for dark matter in the form of WIMPS, there also exists strong constraints on possible electromagnetic interactions of dark matter, even in the limit of complete charge neutrality. A neutral Dirac fermion can posses both a permanent magnetic dipole moment, $\\mu$, and a permanent electric dipole moment, $d$, arising from the dimension-five operator, \\be \\mathcal{L}_{D} = \\frac{i}{2} \\bar f \\sigma_{\\mu\\nu}\\left(\\mu +\\gamma_5 d \\right)f F^{\\mu\\nu}. \\ee While the presence of the magnetic dipole moment does not violate any discrete symmetries, the electric dipole moment requires the violation of parity and CP. Severe constraints exist on the dipole moments of $\\sim 1$ TeV dark matter WIMPS \\cite{Sigurdson:2004zp}. (We are aware that the authors of \\cite{Sigurdson:2004zp} are currently revising their estimates and, as a result, the strength of the constraints in \\cite{Sigurdson:2004zp} may change substantially \\cite{private}.) Thus, if a model predicts Dirac fermionic dark matter, it is important to determine the strength of the induced dipole moment.  Since the KK-neutrino of UED models is a Dirac fermion from the 4-dimensional perspective, UED models that assume KK-neutrino dark matter are constrained by the strength of the induced dipole operator.  In this letter we derive model independent bounds on KK-neutrino dark matter by examining the induced dipole moment and comparing the predictions with the current observational bound. Section II provides a brief review on the properties of KK-fermions in UED models along with a discussion on dipole moments relevant to the calculation presented in section III. Finally, in section IV we present our conclusions. ",
        "conclusions": "UED models have attracted attention as a possible extension to the standard model. A particular appealing feature of the model class centers on the existence of plausible dark matter candidates as the result of KK-parity conservation. While the minimal UED model suggests $B^{(1)}$ dark matter \\cite{Cheng:2002ej}, detailed studies of the relic abundance in the minimal UED model \\cite{coannmat,coannkribs, Matsumoto:2005uh} in combination with electroweak precision constraints \\cite{Flacke:2005hb} show strong observational tension, and thus provide motivation for new possible UED model building avenues. The need for non-minimality has been reported \\cite{Hewett:2004py}. Extensions of the MUED scenario by incalculable boundary terms arising from the UV completion of the model can lead to a different LKP and therefore different possible dark matter candidates. We have taken a model independent approach, following \\cite{Servant:2002aq}, and examined the consequences of UED KK-neutrino dark matter.  As the KK-neutrino is a Dirac fermion, UED models predict an induced KK-neutrino dipole moment. We find that the induced dipole moment, typically $\\mu \\lesssim 10^{-7} \\mu_B$, strongly conflicts -- by over five orders of magnitude -- with the current observational bounds stated in \\cite{Sigurdson:2004zp} for TeV scale dipole dark matter. We reiterate that the constraints provided by \\cite{Sigurdson:2004zp} are currently under revision, and the strength of the stated bounds are expected to change \\cite{private}. The constraints on magnetic dipole moments given in \\cite{Sigurdson:2004zp} would provide the strongest limits on KK-neutrino dark matter. Even in the absence of the strong limits provided by \\cite{Sigurdson:2004zp}, the bounds on dipole moments remain an important constraint on future model building. Not only will new, non-minimal models that predict KK-neutrino LKP need to circumvent the constaints provided by \\cite{Servant:2002hb}, but will also have to evade the constraints provided by radiatively induced magnetic dipole moments, which are generically at least as large as the current experimental bounds.  While we have restricted our discussion to five-dimensional models compactified on $S_1/Z_2$, we expect that the qualitative features carry over to UED models with multiple extra dimensions. We have also assumed the absence of any fine-tuning between the radiatively induced magnetic dipole moment and possible non-renormalizable dimension-five dipole operators arising from the UV complete theory or boundary terms.  Our results indicate that observational limits on dipole dark matter can place significant constraints on UED scenarios where the KK-neutrino is the LKP dark matter candidate."
    },
    "0601/astro-ph0601332_arXiv.txt": {
        "abstract": "Recent progress on the nature of short duration $\\gamma$-ray bursts has shown that a fraction of them originate in the local universe. These systems may well be the result of giant flares from soft gamma-repeaters (highly magnetized neutron stars commonly known as magnetars). However, if these neutron stars are formed via the core collapse of massive stars then it would be expected that the bursts should originate from predominantly young stellar populations, while correlating the positions of BATSE short bursts with structure in the local universe reveals a correlation with all galaxy types, including those with little or no ongoing star formation. This is a natural outcome if, in addition to magnetars formed via the core collapse of massive stars they also form via Accretion Induced Collapse following the merger of two white dwarfs, one of which is magnetic. We investigate this possibility and find that the rate of magnetar production via WD-WD mergers in the Milky Way is comparable to the rate of production via core-collapse. However, while the rate of production of magnetars by core collapse is proportional to the star formation rate, the rate of production via WD-WD mergers (which have long lifetimes) is proportional to the stellar mass density, which is concentrated in early-type systems. Therefore magnetars produced via WD-WD mergers may produce SGR giant flares which can be identified with early type galaxies. We also comment on the possibility that this mechanism could produce a fraction of the observed short duration burst population at higher redshift. ",
        "introduction": "Soft $\\gamma$-repeaters (SGRs) are thought to be highly magnetized neutron stars with rotational periods of seconds and strong ($>10^{14}G$) magnetic fields (see e.g. Kouveliotou et al. 1998). They undergo frequent outbursts and occasional giant flares (Hurley et al. 1999;  Palmer et al. 2005). These giant flares are sufficiently bright to be observed in external galaxies out to large distances ($>50$ Mpc), and would be observable as a short duration $\\gamma$-ray burst (Hurley et al. 2005). However, the paucity of observed giant flares in our own galaxy has so far precluded observationally based determinations of either their luminosity function or rate.   Observations of recently localised short duration GRBs have shown them to be associated with a variety of host galaxy types at typical redshifts of $z \\sim 0.2$ (e.g. Gehrels et al. 2005; Fox et al. 2005; Berger et al. 2005). Their energetics, and location in a variety of host galaxies, including those with little or no ongoing star formation point to an origin in the merger of two compact objects, either NS-NS, or NS-BH (Bloom et al. 2005; Hjorth et al. 2005).  However, if at least some fraction of short duration GRBs are due to SGRs in external galaxies then it may be expected that some correlation between the positions of bursts and the locations of galaxies in the local universe would be seen. Tanvir et al. (2005 - hereafter T05) have recently reported such a correlation, indicating that up to $\\sim 25$\\% of short duration GRBs originate in the local universe (within 100 Mpc). Intriguingly this analysis shows a correlation with all galaxy types, and is strongest when restricted to relatively early-type galaxies (Sbc and earlier). SGRs are thought to be formed in the core collapse of massive stars and due to relatively short lifetimes ($\\sim 10^4$ years; e.g. Kouveliotou 1999) would naturally be located in predominantly star forming galaxies, while essentially none should be seen in ellipticals. Two possibilities present themselves: the first is that rather than representing SGRs the correlations reported by T05 are due to a different progenitor class, such as low-luminosity neutron star mergers. The second possibility is that SGRs are not created exclusively by core collapse events and can be found in older stellar populations. This second option is the subject of this paper.  Here we consider an alternative model for the creation of SGRs, and thus potentially GRBs: namely SGRs which are created via the Accretion Induced Collapse (AIC) of white dwarfs to neutron stars (e.g.  Nomoto \\& Kondo 1991). Such an AIC may occur in WD-WD mergers (King, Pringle \\& Wickramasinghe 2001), or rarely  in binaries with more massive main sequence companions although the precise outcome of WD-WD mergers remains uncertain, and may produce either a SN Ia or an AIC. Once the magnetar has been created its evolution will proceed in an analogous manner to those created via core collapse. ",
        "conclusions": "We have presented a mechanism for the production of magnetars via Accretion Induced Collapse following the merger of two white dwarfs in a tight binary. Considering the observed mass and magnetic field distributions for white dwarfs it appears that the rate of formation of SGRs via this process in the local universe is comparable to the production rate via the core collapse channel (although possibly slightly lower than the core collapse channel in the Milky Way). In early-type galaxies the WD-WD mechanism is the dominant producer of SGRs, while in late type galaxies most SGRs are produced via core collapse. Thus, it would be expected that SGR flares appearing as short duration GRBs may be found from galaxies of all types and this naturally explains the correlation seen between the locations of BATSE short duration bursts and early-type galaxies in the local universe."
    },
    "0601/astro-ph0601618_arXiv.txt": {
        "abstract": "We report the detection of crystalline water ice on the surface of 2003 $\\rm EL_{61}$.  Reflectance spectra were collected from Gemini North telescope from 1.0 to 2.4 \\micron\\/ wavelength range, and from the Keck telescope across the 1.4 to 2.4 \\micron\\/ wavelength range. The signature of crystalline water ice is obvious in all data collected.  Like the surfaces of many outer solar system bodies, the surface of 2003 $\\rm EL_{61}$ is rich in crystalline water ice, which is energetically less favored than amorphous water ice at low temperatures, suggesting that resurfacing processes may be taking place.  The near infrared color of the object is much bluer than a pure water ice model.  Adding a near infrared blue component such as hydrogen cyanide or phyllosilicate clays improves the fit considerably, with hydrogen cyanide providing the greatest improvement.  The addition of hydrated tholins and bitumens also improves the fit but is inconsistent with the neutral $V-J$ reflectance of 2003 $\\rm EL_{61}$.  A small decrease in reflectance beyond 2.3 \\micron\\/ may be attributable to cyanide salts.  Overall, the reflected light from 2003 $\\rm EL_{61}$ is best fit by a model of 2/3 to 4/5 pure crystalline water ice and 1/3 to 1/5 near infrared blue component such as hydrogen cyanide or kaolinite.  The surface of 2003 $\\rm EL_{61}$ is unlikely to be covered by significant amounts of dark material such as carbon black, as our pure ice models reproduce published albedo estimates derived from the spin state of 2003 $\\rm EL_{61}$. ",
        "introduction": "\\label{intro}  The Kuiper belt objects (KBOs) have been known for a decade to have the most diverse surfaces of any minor planet population in terms of global color at visible wavelengths \\citep{1996AJ....112.2310L}.  The carrier of this color has been tentatively identified as a red tholin-like compound \\citep{2001AJ....122.2099J}.  The situation is much different in the near-infrared, where signatures of the most primitive solar system volatiles can be seen.  Although many attempts have been made to detect ices on KBOs in the near-infrared, very few to date have been successful.  Barring the decades of study of Pluto, Charon and Triton (each of which could share origins with the KBOs), detections of simple organic ices on KBOs have been relatively few.  Indisputable methane ice detections have been presented for 2003 $\\rm UB_{313}$ and 2005 $\\rm FY_{9}$ \\citep{2006btr,2005DPS....37.5211B,2006A&A...445L..35L}.  The KBO Sedna has also been reported to have a Triton-like spectrum, largely dominated by methane ice \\citep{2005A&A...439L...1B}.  The following KBOs have all shown signs of water ice: (19308) 1996 $\\rm TO_{66}$ \\citep{1999ApJ...519L.101B}, (26375) 1999 $\\rm DE_{9}$ \\citep{2001AJ....122.2099J} and (90482) Orcus \\citep{2005ApJ...627.1057T}.  The object (50000) Quaoar has been the only KBO so far to show signs of crystalline water ice \\citep{2004Natur.432..731J}.  This finding is significant in that crystalline water ice is unstable on 10 Myr timescales.  However, crystalline water ice has been seen on many other outer solar system bodies such as Charon and the giant planet moons, which argues for a common resurfacing process throughout the solar system.  In this work, we report the discovery of large amounts of crystalline water ice on 2003 $\\rm EL_{61}$ from observations obtained at the Gemini North telescope and Keck Observatory.  The object 2003 $\\rm EL_{61}$ is a unique object in dynamical terms.  It is the only known ternary or higher system in the Kuiper belt \\citep[excepting Pluto,][]{2005IAUC.8625....1W}, with the brighter moon's orbit reported in \\cite{2005ApJ...632L..45B} and the fainter moon reported in \\cite{2006betal}.  Not only is 2003 $\\rm EL_{61}$ a ternary system, but it is the most rapid large ($> 100$ km) rotator in the solar system \\citep{2005DPS....37.5612R}.  The extreme rotational state of 2003 $\\rm EL_{61}$ combined with the measured orbit of its brighter moon yields a surprisingly large amount of information about the physical parameters of the system including limits on its shape, density and albedo \\citep{2006retal}.  With the discovery of crystalline water ice on the surface of 2003 $\\rm EL_{61}$, we can now place constraints on the surface fraction of water ice on the body as well as other components. ",
        "conclusions": "\\label{discussion}  The basic result of our observations is that the spectrum of 2003 $\\rm EL_{61}$ is consistent with a model with roughly 2/3 to 4/5 pure crystalline water ice combined with 1/3 to 1/5 pure blue material. The best candidates for the blue material are pure hydrogen cyanide or possibly phyllosilicate clays, although this identification is not unique.  Likely there are many possible materials that could fit the near-infrared color of 2003 $\\rm EL_{61}$, of which only a fraction have laboratory spectra available.  No neutral absorbers are needed in appreciable amounts to model the relative reflectance of the body. The only feature noted in the near infrared spectrum besides that of crystalline water ice and the basic blue color is the drop in reflectance beyond 2.3 \\micron\\/ which is consistent with copper potassium cyanide.  Laboratory data on cyanide compounds are sparse, so it is likely that other cyanide or nitrile compounds could produce such behavior.  The presence of crystalline water ice has been widely reported for many outer solar system bodies such as the satellites of the major planets, Charon and most recently Quaoar.  As noted by \\cite{2004Natur.432..731J} for (50000) Quaoar, the presence of crystalline water ice is considered problematic in the Kuiper belt because, in general, lower energy amorphous ice is favored in cold environments.  Crystalline water ice may have been prevalent in the early solar nebula and may have been incorporated into bodies at low temperatures ($< 100$ K) if the deposition flux of water molecules was low \\citep{1994A&A...290.1009K}.  However, the disruptive action of cosmic ray and solar flux bombardment on the outer surfaces of icy bodies is expected to rearrange crystalline water ice to amorphous in the outer solar system \\citep{2004JGRE..10901012H}.  Both solar flux and cosmic ray bombardment can cause extensive radiation damage to the outer surface layers of KBOs.  Estimating the time scale of such radiation damage is very difficult due to the sparse amount of data available about energetic particles in the solar system, which comes primarily from spacecraft.  In addition, assumptions about the timescale for density changes in the local interstellar medium must be made which directly affects all estimates of the extent of the heliosphere, which is a key component of radiation damage models.  Our near-infrared reflectance spectra probe only the first $\\sim 1$ mm of the KBO surface when surface density and multiple scattering effects are taken into account. \\cite{2003EM&P...92..261C} suggest that disruptive amounts of energy (of order $\\sim 100$ eV per 16 amu) can be delivered to $\\sim 1$ mm depths by Galactic Cosmic Ray bombardment throughout the heliosphere. Timescales several orders of magnitude shorter may be plausible both closer to and farther from the sun than the $\\sim 40$ AU region of the Kuiper Belt due to proton bombardment from a variety of sources.  As many of these KBOs have been at their present location for $10^9$ years even in extreme dynamical models \\citep{2005Natur.435..466G}, it appears that a solar system wide model for surface renewal is likely needed to explain the presence of crystalline water ice.  Hydrogen cyanide, one of the species that can explain the blue near infrared color of 2003 $\\rm EL_{61}$, is also photolysed rapidly in the outer solar system environment.  Whether alone or frozen in solution with water ice, hydrogen cyanide is destroyed after deposition of about 100 eV per molecule by both irradiation and photolysis \\citep{2005ApJ...620.1140G}.  Thus if HCN is to be considered as a source of the blue material on 2003 $\\rm EL_{61}$, it has the same stability problems as crystalline water ice.  The phyllosilicate clays, another possible infrared blue material, do not share this sensitivity to radiation damage."
    },
    "0601/astro-ph0601104_arXiv.txt": {
        "abstract": "The nature of the X-ray source RX J1914+24 has been the subject of much debate. It shows a prominent period of 569 sec in X-rays and the optical/infra-red: in most models this has been interpreted as the binary orbital period. We present our analysis of new {\\xmm} and {\\chan} data. We find a longer term trend in the {\\xmm} data and power at 556 and 585 sec in 5 sets of data. It is not clear if they are produced as a result of a beat between a longer intrinsic period and the 569 sec modulation or if they are due to secular variations. We obtain a good fit to the {\\xmm} spectrum with a low temperature thermal plasma model with an edge at 0.83keV. This model implies an unabsorbed bolometric X-ray luminosity of $1\\times10^{33}$ \\ergss (for a distance of 1kpc) - this is 2 orders of magnitude lower than our previous estimate (derived using a different model). If the distance is much less, as the absorption derived from the X-ray fits suggest, then it is even lower at $\\sim3\\times10^{31}$ \\ergss. ",
        "introduction": "The nature of the X-ray source RX J1914+24 has been the subject of much debate in recent years. Only one modulation period (569 sec) has been seen in both its X-ray and optical light curves (Ramsay et al 2000). For this, and other reasons, the period of 569 sec has been taken to be the binary orbital period, which would make it the second most compact binary system known.  The models which have been put forward to explain the observational characteristics of RX J1914+24 fall into two broad categories: accretion and non-accretion driven models. The fact that the optical spectrum of RX J1914+24 (also known as V407 Vul) shows no emission lines, puts a question mark over the accretion driven models. The non-accreting model is the unipolar-inductor (UI) model of Wu et al (2002). While this model has been able to explain many of the observed characteristics of this source, it appears that its X-ray luminosity implies an induction torque which is much greater than the amount generated by gravitational radiation. There is some evidence that the observed spin-up rate of RX J1914+24 (Strohmayer 2002, 2004b, Ramsay et al 2005a) is therefore much smaller than that predicted by the UI model (Marsh \\& Nelemans 2005). This conclusion is based on the X-ray luminosity determined using {\\ros} (Cropper et al 1998) and {\\xmm} observations (Ramsay et al 2005a). The {\\ros} data were well fitted using an absorbed blackbody model while the higher signal to noise data of {\\xmm} found a poor fit using this model. Ramsay et al (2005a) found the best fit to the data could be achieved using an absorbed blackbody plus Gaussian model, although they did not expect that this model was physically realistic.  Since those first {\\xmm} observations were made in Nov 2003, a further two sets of observations made in Sept/Oct 2004 have become publically available. These observations were longer than the initial {\\xmm} observations and the particle background was lower. In addition, a further set of X-ray observations of RX J1914+24 made using {\\chan} has become publically available.  We have therefore extracted these data from the archive and use them to make a more detailed investigation of its X-ray spectrum, to determine if the 569 sec is continuing to spin-up at the same rate and to search for evidence for periods other than the 569 sec period. ",
        "conclusions": "\\subsection{The periods in RX J1914+24} \\label{discussion}  Until now there has been no convincing evidence for any other periods in RX J1914+24 apart from the 569 sec period. In the case of 5 X-ray observations we have found evidence for power at 552 and 584 sec. This changes our perception of RX J1914+24. The data are clear - the X-ray flux clearly increases in the {\\xmm} data taken in both orbit 0880 and in orbit 0882. This could be considered to be the result of an underlying $\\sim$6 hour period, or it could be the result of secular variability known already in this system at longer timescales (Ramsay et al 2000, Strohmayer 2004b). Both could produce the observed amplitude spectrum.  To explore this question, we have obtained the mean X-ray flux in each of the {\\chan} and {\\xmm} observations. Since the {\\chan} ACIS-I detector has a low effective area at energies less than $\\sim$0.4keV, we determined the observed mean flux in the 0.4-1.0keV energy band. We show how the flux in this band varies over time in Figure \\ref{long_term}. We also show those epochs at which we detected the 552 sec and 584 sec periods which could be signatures of a longer period trend in the observation. Although not conclusive, it appears that at the higher flux levels there is no (or weak) evidence of these other periods, so that shorter ($<$1 day) timescale variations may take place only when the system is emitting at intermediate or low levels. The detection of power at 552 and 584 sec does not appear to be related to the observation length since the longest observation (the {\\chan} observation of Feb 2003 which lasted 35 ksec) did not show obvious power at these periods.  At this stage, we cannot distinguish between the X-ray long term behaviour being due to random changes in either the accretion rate (in the case of the accreting models) or the electrical current dissipation (in the case of the unipolar-inductor model). We urge the search for radial velocity variations on a timescale of less than 1 day, or dedicated X-ray and/or optical observations of this system to search for variations on this period.  {\\it If} the longer term variation in the X-ray flux is periodic, then we showed earlier that its period would approximately be 6 hrs.  Smith \\& Dhillon (1998) show that for a binary system with this orbital period then its spectral type would be expected to be a mid K-type, although the range is up to 5 sub-classes.  We note with interest the discovery of Steeghs et al (2006) of features in the optical spectrum of RX J1914+24 which appear to resemble a G9/K0 spectral type. Therefore, an orbital period of 6 hrs would be marginally consistent with this.  \\begin{figure} \\begin{center} \\setlength{\\unitlength}{1cm} \\begin{picture}(6,6) \\put(-2,-0.3){\\special{psfile=long_term_flux.ps hscale=40 vscale=40}} \\end{picture} \\end{center} \\caption{The observed flux in the 0.4-1keV energy band (units of $10^{-12}$ \\ergscm) determined in {\\chan} and {\\xmm} observations. `XP' denotes the presence of periods at 552 and 583 sec: in the case of the most recent {\\chan} observation the significance is marginal - `WXP' weak X-ray periods.} \\label{long_term} \\end{figure}  \\subsection{The X-ray spectrum of RX J1914+24}  We have shown that the derived luminosity of RX J1914+24 is highly model dependent. Further, the distance to RX J1914+24 is still rather uncertain, but Steeghs et al (2006) suggest a distance of 1kpc based on its optical spectrum. For a distance of 1kpc the mean bolometric luminosity from the {\\xmm} orbits 0880 and 0882 data is therefore 1.4$\\times10^{33}$ \\ergss - an order of magnitude lower than the previous luminosity of 3.8$\\times10^{34}$ \\ergss. Ramsay et al (2005a) went further and applied a correction to account for projection affects giving a luminosity of $\\sim1\\times10^{35}$ \\ergss. Since the new model is an optically thin plasma, a correction effect for projection is now no longer applicable. Moreover, absorption derived from the X-ray fits is $\\sim$15\\% of that derived from the optical data of Steeghs et al (2006) based on a G9 spectral type. If the G star is not the secondary, then we can reduce the distance by a similar factor and the luminosity is therefore 3$\\times10^{31}$ \\ergss. In this case, the X-ray luminosity of RX J1914+24 is more than $\\sim$3 orders of magnitude less than that found by Ramsay et al (2005a).  In the unipolar-inductor model of Wu et al (2002) the amount of power which is emitted is determined by various factors, with the degree of asynchronisity being one of the key determinants (along with component masses and magnetic field strength). However, to produce the observed luminosity, the asynchronisity would have to be 0.001 or less (Wu et al 2002). This reduces the objections to the UI model on energetic grounds raised by Marsh \\& Nelemans (2005) to a significant degree.  The abundances are clearly non-solar, which is consistent for an evolved secondary as is seen in the fits to spectra of AM CVn systems (eg Marsh, Horne \\& Rosen 1991, Strohmayer 2004a, Ramsay et al 2005b). If this model reflects the properties of the material leaving the secondary star then it rules out the G9/K0 star seen in the optical spectrum being the secondary star."
    },
    "0601/astro-ph0601274_arXiv.txt": {
        "abstract": "\\vskip 0.5cm \\noindent In this article, we examine a model which proposes a common explanation for the presence of additional attractive gravitational effects -- generally considered to be due to dark matter -- in galaxies and in clusters, and for the presence of a repulsive effect at cosmological scales -- generally taken as an indication of the presence of dark energy. We therefore consider the behavior of a so-called dark fluid based on a complex scalar field with a conserved $U(1)$-charge and associated to a specific potential, and show that it can at the same time account for dark matter in galaxies and in clusters, and agree with the cosmological observations and constraints on dark energy and dark matter. \\vspace*{2.cm} ",
        "introduction": "\\noindent After many years of considering that the Universe is filled with a dark matter made of Weakly Interacting Massive Particles (WIMPs) and with a dark energy causing an acceleration of the expansion of the Universe, we have still no direct evidence of their existences. Moreover, observations of supernov\\ae~of type Ia tend to cause trouble to usual dark energy models \\cite{SNIa}, and open the way to new kind of models or analyses to explain the observed repulsing effects \\cite{phantom,coley}.\\\\ Alternative models exist to explain the cosmological observations, and in particular some of them try to solve the dark energy and dark matter problems by unifying both components into a single ``dark fluid''. We can for example note that the Generalized Chaplygin Gas model \\cite{chaplygin} follows this idea and is presently under scrutiny. We have shown in \\cite{arbey_darkfluid} that building such a dark fluid model is very difficult, as the model has to be in agreement with many observational constraints, and especially has to explain at the same time the cosmological repulsive effects and the local binding gravitational effects in the recent Universe. As scalar field-based models for dark energy and dark matter exist in the literature \\cite{quintessence,quintessence2,hessence,matter_field,fuzzy}, it seems interesting to study unifying dark fluid models based on scalar fields. This idea has been proposed in \\cite{padmanabhan}. The crucial question however concerns the form of the potential of the scalar field.\\\\ In this article, we will consider a complex scalar field and propose an adequate form for its potential. We will show that this scalar field can potentially explain correctly the observations of galaxy rotation curves of spiral galaxies and the presence of strong binding gravitational effects in clusters. We will also consider the cosmological behavior of the dark fluid and show that it agrees with the cosmological constraints and observations. We will finally conclude by suggesting some further directions of investigation beyond this study. ",
        "conclusions": "\\noindent In this work, we have shown that it is possible to elaborate a dark fluid model replacing the standard dark energy / dark matter models, using a complex scalar field. We have seen that a potential containing a quadratic term and a Gaussian term is suitable to explain the gravitation observations at galaxy, cluster and cosmological scales. Based on \\cite{arbey_quadratic}, we also know that this model can account for a large variety of rotation curve shapes. Thus, we can conclude that this model could be an interesting alternative to the dark matter / dark energy models.\\\\ Of course, many studies are still needed to test this model. In particular, it would be interesting to consider the formation of large scale structures and to study the influence of the dark fluid on the anisotropies of the Cosmic Microwave Background. Another question concerns the form of the potential. Indeed, this potential has interesting properties, but is not directly derived from high energy theories. We can consider it as a toy-potential rather than a definitive answer. It would then be interesting to perform the study of different kinds of potentials, and perhaps to find an adequate potential derived from a high energy theory.\\\\ In any case, this simple model looks promissing, and reinforces the suggestion that dark fluid models should be studied deeply."
    },
    "0601/astro-ph0601042_arXiv.txt": {
        "abstract": "{There is observed a trend that a lower mass galaxy forms stars at a later epoch. This downsizing of star-forming galaxies has been attributed to hydrodynamical or radiative feedback processes that regulate star formation.} {We explain the downsizing by gravitational processes alone, in the bottom-up scenario where galaxies evolve from subgalactic-scale objects.} {Using the theory of Press and Schechter, as a function of galaxy mass, we study the peak epoch of formation of subgalactic-scale objects, i.e., gravitational collapse of subgalactic-scale density fluctuation.} {The subgalactic-scale density fluctuation is found to collapse at a later epoch for a lower mass galaxy. The epoch is close to the peak epoch of star formation derived from observations.} {The downsizing is inherent in gravitational processes of the bottom-up scenario. Within a region of the initial density field that is to evolve into a lower mass galaxy, subgalactic-scale fluctuation is of a smaller amplitude. Gravitational collapse of the subgalactic-scale density fluctuation and the subsequent onset of star formation occur at a later epoch for a lower mass galaxy. } ",
        "introduction": "\\label{s1}  The current paradigm for the formation and evolution of galaxies is the bottom-up scenario based on the cold dark matter model. The cold dark matter retains significant small-scale fluctuation in the initial density field. Since the first objects are formed within dark matter halos that originate in gravitational collapse of the initial density fluctuation, these objects are not massive. They merge via gravitational clustering of the dark matter halos, evolve into low-mass galaxies, and then evolve into high-mass galaxies.   Observationally, there is a trend that is seemingly inconsistent with the bottom-up scenario, i.e., the ``downsizing'' of star-forming galaxies (Cowie et al. \\cite{Cowie}; Heavens et al. \\cite{Heavens}; Thomas et al. \\cite{Thomas}; Treu et al. \\cite{Treu}). Regardless of environment, a lower mass galaxy forms stars at a later epoch. In particular, a lower mass galaxy undergoes the peak of star formation at a later epoch.  The existing explanations for the downsizing invoke hydrodynamical or radiative feedback processes that regulate star formation (De Lucia et al. \\cite{Lucia}; see also the above references). For example, if the galaxy is not massive, supernova explosions could expel gas from the galaxy and thereby inhibit star formation until the fallback of the gas. However, here we demonstrate that the downsizing is in fact explainable by gravitational processes alone, using the theory of Press \\& Schechter (\\cite{Press}) for the bottom-up formation of dark matter halos. ",
        "conclusions": "\\label{s5}  The downsizing of star-forming galaxies, with a lower mass galaxy forming stars at a later epoch, is inherent in gravitational processes of the bottom-up scenario. Within a region of the initial density field that is to evolve into a lower mass galaxy, subgalactic-scale fluctuation is of a smaller amplitude (Fig. \\ref{f1}). Gravitational collapse of the subgalactic-scale fluctuation and the subsequent onset of star formation occur later for a lower mass galaxy. As a function of galaxy mass, we have calculated the peak epoch of gravitational collapse of subgalactic-scale fluctuation (Fig. \\ref{f2}). The peak epoch is consistent with the peak epoch of star formation derived from observations.   Our explanation for the downsizing differs from the existing explanation for the biasing, where gravitational collapse of small-scale density fluctuation occurs earlier if the large-scale overdensity in that region has a larger height (Kaiser \\cite{Kaiser}; Bardeen et al. \\cite{Bardeen}). We have demonstrated that, if the large-scale overdensity has a fixed height, its scale determines the epoch of gravitational collapse of small-scale density fluctuation (Fig. \\ref{f1}). For the biased galaxy formation, the small and large scales correspond to galactic and supergalactic scales. For the downsizing of star-forming galaxies, the small and large scales correspond to subgalactic and galactic scales. The downsizing of star-forming galaxies is thereby independent of the biased galaxy formation.   Thus far, the downsizing has been explained by hydrodynamical or radiative feedback processes that regulate star formation. Although these processes are at work, they would not be essential. To the formation and evolution of galaxies, hydrodynamical and radiative processes are less essential than gravitational processes that solely explain the downsizing. Also, an observation suggests that the downsizing continuously extends from galaxies to clusters of galaxies where the star formation rate per unit mass at the present epoch is lowest (Feulner et al. \\cite{Feulner}). Over a wide mass range from galaxies to clusters of galaxies, no processes are the same. The exception is gravitational processes, which apply to clusters of galaxies only if we consider subgalactic-scale density fluctuation superposed on cluster-scale overdensities in the initial field.   There is observed another trend that a lower mass galaxy has a longer history of star formation around its peak (Heavens et al. \\cite{Heavens}; Thomas et al. \\cite{Thomas}; see also De Lucia et al. \\cite{Lucia}). This trend, obtained by plotting the star formation rate against epoch $t$, is explainable by the accelerated expansion of the universe and the resultant suppression of gravitational clustering of dark matter halos. The trend is not significant if the star formation rate is plotted against the linear growth factor $D(t)$.   The downsizing is also expected for clusters of galaxies undergoing galaxy formation. Within a region of the initial density field that is to evolve into a lower mass cluster, galactic-scale fluctuation is of a smaller amplitude. The galaxy formation, i.e., gravitational collapse of the galactic-scale fluctuation, occurs at a later epoch. Hence, a lower mass cluster is expected to contain younger galaxies. Further studies would be important to understand more about gravitational processes in the bottom-up scenario."
    },
    "0601/astro-ph0601568_arXiv.txt": {
        "abstract": "The Time Projection Chamber (TPC) has been recognized as a potentially powerful detector for the search of WIMPs by measuring the directions of nuclear recoils, in which the most convincing signature of WIMPs, caused by the Earth's motion around the Galaxy, appears.  We report on the first results of a performance study of the neutron exposure of our prototype micro TPC with Ar-C$_2$H$_6$ (90:10) and CF$_4$ gas at 150 Torr. ",
        "introduction": "\\label{intro} It is considered by many that the galactic halo is composed of weakly interacting massive particles (WIMPs) as dark matter\\cite{jung}. These particles could be detected directly by measuring the nuclear recoils produced by their elastic scattering off nuclei in detectors. The most convincing signature of WIMPs appears in the directions of nuclear recoils. It is provided by the Earth's large velocity through the isothermal galactic halo ($\\sim$230 km/s). Hence, detectors sensitive to the direction of the recoil nucleus would have a great potential to identify WIMPs\\cite{dir}.  Time Projection Chambers (TPCs) with fine spacial resolutions are among such devices, and we are developing a micro TPC, which can detect three-dimensional fine tracks of charged particles\\cite{miuchi}. Since the energy deposits of WIMPs to nuclei are only a few tens of keV and the range of nuclei is limited, the micro-TPC should be operated at low pressures.  We also focused on the detection of WIMPs via spin-dependent(SD) interactions and are interested in operating the micro-TPC with CF$_4$\\cite{NA}, because $^{19}$F has a special sensitivity to SD interactions for its unique spin structure\\cite{collar}.  In the present work, in order to examine the response of the micro-TPC to nuclear recoils at low pressures as a first step, we irradiated a 150 Torr Ar-C$_2$H$_6$ (90:10 mixture) gas (one of the standard gases for TPCs) and CF$_4$ with neutrons from $^{252}$Cf. The track lengths and deposited energies of Ar, C, and F recoils were investigated. ",
        "conclusions": "We successfully showed the nuclear recoils in 150 Torr of Ar-C$_2$H$_6$ (90:10) and CF$_4$ gases according to their $dE/dx$ by changing the detector threshold. The energy loss of protons became lower as the energy increased as opposed to the other nuclei\\cite{srim}. Consequently, the proton band in Fig. \\ref{fig:ArTE}(b) is truncated at the threshold set in the measurements, which corresponds to about 5 keV/400$\\mu$m. In terms of $dE/dx$, the tracks in the micro TPC were much easier to detect for C, F or Ar recoils.  Ultimately, our concern is the recoil direction of such nuclei below 100 keV to allow us to observe the signals of WIMPs. In order to obtain longer tracks and clear Bragg curves, higher gas gain operations at lower pressures with low-energy neutron beams are needed. The measurement of the incident neutron energy with Time-Of-Flight may also be useful to examine the quenching factor of nuclear ionization in the micro-TPC."
    },
    "0601/hep-ph0601080_arXiv.txt": {
        "abstract": " ",
        "introduction": "The  inflationary   paradigm  constitutes  an   ingenious  theoretical framework,  in which  many  of the  outstanding  problems in  standard cosmology  can be  successfully  addressed~\\cite{review}.  The  recent WMAP    data~\\cite{WMAP},    compiled    with    other    astronomical observations~\\cite{MT,Lyman}, improved  upon the precision  of about a dozen of  cosmological parameters.   These include the  power spectrum $P^{1/2}_{{\\cal  R}}$ of curvature  perturbations, the  spectral index $n_s$,  the baryon-to-photon  ratio of  number densities  $\\eta_B$ and others.   The  values of  these  cosmological  observables put  severe constraints on  the model-building  of successful models  of inflation and  their theoretical  parameters.  For  instance, one  of  the basic requirements for  slow-roll inflation  is that the  so-called inflaton potential be flat.  In this respect, supersymmetry (SUSY) emerges as a compelling ingredient in model-building for protecting the flatness of the inflaton potential against quantum corrections.  In addition to the aforementioned element of naturalness, inflationary models would have a greater value if they were predictive and testable as well.  One  such predictive and perhaps most  appealing scenario is the well-celebrated  model of hybrid  inflation~\\cite{Linde}.  In this model, the inflaton  field $\\phi$ can start its  slow-roll from values well  below the  Planck scale  $m_{\\rm Pl}  =  2.4\\times 10^{18}$~GeV. This renders the model very  predictive, in the sense that an infinite set of possible higher-dimensional non-renormalizable operators, being suppressed by  inverse powers of $1/m_{\\rm Pl}$,  will not generically contribute   significantly  to   cosmological  observables,   such  as $P^{1/2}_{{\\cal R}}$  and $n_s$.  In the hybrid  model, inflation ends through  the  so-called waterfall  mechanism,  once  the field  $\\phi$ passes below  a critical value  $\\phi_c$.  When this  happens, another field  $X$  different  from  $\\phi$,  with  vanishing  initial  value, develops  a tachyonic  instability and  rolls  fast down  to its  true vacuum    expectation    value~(VEV).    Super-\\linebreak    symmetric realizations of  hybrid inflation  from $F$-terms were  first analyzed in~\\cite{CLLSW,DSS}, whereas hybrid  inflation triggered by a dominant Fayet--Iliopoulos~(FI) $D$-term~\\cite{FI}  was subsequently considered in~\\cite{Halyo}.  In  this  paper  we  study   $F$-term  hybrid  inflation  in  a  novel supersymmetric extension  of the  Standard Model~(SM) that  includes a subdominant  FI  $D$-term. We  call  this  scenario  for brevity,  the $F_D$-term hybrid model.  To account for the low-energy neutrino data, we introduce 3 singlet neutrino superfields $\\widehat{N}_{1,2,3}$ that contain   3   right-handed   neutrinos  $\\nu_{1,2,3\\,R}$   and   their supersymmetric  scalar   counterparts  $\\widetilde{N}_{1,2,3}$.   Most importantly,  the  model  ties  the  $\\mu$-parameter  of  the  Minimal Supersymmetric  Standard Model~(MSSM) to  an SO(3)  symmetric Majorana mass    $m_N$,   through    the    VEV   of    the   inflaton    field $\\phi$~\\cite{PU2,Francesca}.   Hence,   the  $F_D$-term  hybrid  model naturally predicts lepton-number  violation at the TeV or  even at the electroweak scale. In this  scenario, the non-zero baryon asymmetry in the  Universe (BAU),  $\\eta_B  \\approx 6.1  \\times  10^{-10}$, can  be explained  by  leptogenesis~\\cite{FY,BAUpapers}  and  specifically  by thermal electroweak-scale resonant leptogenesis~\\cite{APRD,PU2}.  In this paper  we also study the constraints on  the parameters of the $F_D$-term hybrid model that  result from a reheat temperature $T_{\\rm reh} \\stackrel{<}{{}_\\sim} 10^9$~GeV, which  is necessary to avoid the well-known gravitino overproduction  problem.  This consideration puts severe  limits on  the size  of  the superpotential  couplings of  the theory, forcing  them {\\em all}  to acquire rather  suppressed values, namely to be smaller than about $10^{-5}$~\\cite{SS}.  To overcome this problem of unnaturalness, the presence of a subdominant FI $D$-term in the  $F_D$-term  hybrid model  is  very  crucial  and provides  a  new mechanism  of  relaxing  dramatically  the above  upper  limit.   More explicitly, the size  of the $D$-term controls the  decay rates of the ultraheavy fermions  and bosons associated with the  extra gauge group U(1)$_X$.    In  the   absence   of  the   $D$-term   and  any   other non-renormalizable   interaction,    these   ultraheavy   gauge-sector particles are  absolutely stable. On  the other hand,  these particles are abundantly  produced during the preheating  epoch, thus dominating the energy  density of  the Universe shortly  after the period  of the first   reheating  caused   by  the   perturbative   inflaton  decays. Therefore, their late decays induced by a non-vanishing $D$-term could give rise  to a second reheating  phase in the evolution  of the early Universe. Depending on the actual size of the FI $D$-term, this second reheat temperature may be as low  as 1~TeV, giving rise to an enormous entropy  release that can  dilute the  gravitinos produced  during the first reheating to an unobservable level.  The paper is organized  as follows: Section~\\ref{FDmodel} presents the model-building   aspects   of  the   $F_D$-term   hybrid  model   with electroweak-scale  lepton-number violation. Technical  details related to the  radiatively-induced FI  $D$-term are relegated  to Appendix~A. In Section~\\ref{reheat},  we estimate the reheat  temperature from the perturbative  inflaton  decays  and  derive  the  resulting  gravitino constraint  on  the  theoretical  parameters.   We  then  discuss  the non-perturbative production of the supermassive fields associated with the U(1)$_X$  gauge group  during the preheating  epoch and  how their late decays can help to lower the reheat temperature even up to~1~TeV. Section~\\ref{inflation} is devoted  to inflation.  Here we investigate two regimes: (i) cold  hybrid inflation, where dissipative effects can be ignored, and (ii)  warm hybrid inflation, where dissipative effects dominate   over   the   expansion    rate   of   the   Universe.    In Section~\\ref{BAU} we  illustrate how  the $F_D$-term hybrid  model can realize  thermal  electroweak-scale  resonant leptogenesis.   Finally, Section~\\ref{conclusions}  summarizes  our  conclusions,  including  a brief discussion  of further  possible implications of  the $F_D$-term hybrid model for particle physics and cosmology.     \\setcounter{equation}{0} ",
        "conclusions": "We have  studied $F$-term hybrid  inflation in a  novel supersymmetric extension of the SM, to which  a subdominant FI $D$-term was added. We called this particular form  of inflation $F_D$-term hybrid inflation. The $F_D$-term hybrid model we  have been analyzing in this paper ties the $\\mu$-parameter  of the MSSM  to an SO(3) symmetric  Majorana mass $m_N$, through the  VEV of the inflaton field.   As a consequence, the model  predicts   {\\em  naturally}  lepton-number   violation  at  the electroweak scale.  In order to obtain predictions for the observables $P_{\\cal R}^{1/2}$, $n_s$    and   $\\eta_B$    compatible    with   global    cosmological analyses~\\cite{Lyman},   as  well   as   interesting  particle-physics phenomenology  that could  be  tested in  laboratory experiments,  one needs to make  certain assumptions for the model  of $F_D$-term hybrid inflation: \\begin{itemize}  \\item[ (i)] Successful hybrid  inflation relies on the assumption that the inflaton field is displaced from its minimum in the beginning of inflation,  whereas all  other non-inflaton  fields have  zero VEVs, according to (\\ref{initial}).  \\item[ (ii)] The present $F_D$-term hybrid scenario utilizes a minimal K\\\"ahler  potential,  where  terms  of  order  $H^2  |S|^2$  in  the potential  are  set to  zero  or  assumed  to be  negligible.   This consideration introduces  some tuning  in general SUGRA  models with non-minimal K\\\"ahler potentials.  \\item[(iii)] In order to get  a red-tilted spectrum with negative $n_s - 1$, one has to assume  that the radiative corrections dominate the slope of  the inflationary  potential.  This possibility  arises for superpotential  couplings:  $10^{-4} \\stackrel{<}{{}_\\sim}  \\kappa,\\ \\lambda,\\ \\rho \\stackrel{<}{{}_\\sim} 10^{-2}$.  \\item[ (iv)] Even  though a bare $D$-tadpole may be  present as a bare parameter  in the  tree-level Lagrangian,  we have  considered here, however, the  possibility that such a term  is generated radiatively after  heavy degrees  of freedom  have been  integrated  out.  These heavy  degrees  of freedom  are  assumed  to  be Planck-mass  chiral superfields  which are  oppositely  charged under  the U(1)$_X$  and which break explicitly the  discrete charge symmetry discussed after (\\ref{VFD}) and in Appendix~A.  \\item[ (v)] We  have assumed that the coupling  $\\rho$ of the inflaton to  neutrino superfields  is SO(3)  symmetric or  very close  to it. After  the inflaton  receives  a VEV,  one  ends up  with 3  nearly degenerate heavy  Majorana neutrinos with masses  at the electroweak scale. This enables  one to successfully address the  BAU within the thermal electroweak-scale  resonant leptogenesis framework  (see our discussion  in Section~\\ref{BAU}).   As has  also been  discussed in Section~\\ref{BAU}, if one assumes that the neutrino-Yukawa couplings $h^\\nu_{ij}$ have a certain hierarchical structure controlled by the approximate  breaking of  global flavour  symmetries, the  model can have  further  testable  implications  for  $e^+e^-$  colliders  and low-energy experiments of lepton flavour and/or number violation.   \\end{itemize}  The requirement for a sufficiently low reheat temperature $T_{\\rm reh} \\stackrel{<}{{}_\\sim} 10^9$~GeV, which does not lead to overproduction of  gravitinos,   provides  an  important  constraint   on  the  basic theoretical  parameters  $\\kappa$, $\\lambda$  and  $\\rho$.  The  naive limits on  these couplings derived from reheating  due to perturbative inflaton  decay   are  very  strict,   i.e.~$\\kappa,\\  \\lambda,\\  \\rho \\stackrel{<}{{}_\\sim} 10^{-5}$. These limits may be completely avoided by considering the late  decays of the U(1)$_X$ gauge-sector particles which  are induced  by  a non-vanishing  FI  $D$-term $m^2_{\\rm  FI}$. Their  decay rates  depend  crucially on~$m^2_{\\rm  FI}$. As  menioned above  in  point~(iv)  and  in  Appendix~A, the  generation  of  a  FI $D$-tadpole  and its  size may  be engineered  by  adding Planck-scale heavy degrees  of freedom to the  theory and by  subjecting these into extended $R$  symmetries.  In  this way, a  phase of  second reheating takes place in the evolution of  the early Universe, which can lead to a significant lowering of the reheat temperature even up to 1~TeV.  The  $F_D$-term  hybrid  model  with electroweak-scale  lepton  number violation can easily be embedded  within a minimal SUGRA theory, where all  soft  SUSY-breaking  parameters  are  constrained  at  the  gauge coupling unification point $M_X$ which  can be chosen to be $M \\approx 10^{16}$~GeV.  Instead,  electroweak baryogenesis  is not viable  in a minimal  SUGRA  scenario  of  the  MSSM.  Moreover,  the  CP-odd  soft SUSY-breaking phases required  for successful electroweak baryogenesis face severe  constraints from the non-observation of  the electron and neutron  electric  dipole  moments,   even  though  the  latter  arise diagrammatically at the 2-loop level~\\cite{CKP}.  The $F_D$-term hybrid model under discussion conserves $R$-parity. The reason is that all  superpotential couplings either conserve the $B-L$ number  or  break it  by  even  number  of units.   Specifically,  the operator   $\\widehat{S}\\widehat{N}_i\\widehat{N}_i$  breaks  explicitly $L$,  as  well  as  $B-L$,  by 2~units.   Consequently,  the  lightest supersymmetric  particle  (LSP) of  the  spectrum  is  expected to  be stable,  thus providing a  viable candidate  to address  the so-called Cold Dark Matter  (CDM) problem.  The new aspect of  our model is that right-handed   sneutrinos  could   be   the  LSPs,   opening  up   new possibilities in the phenomenology of CDM and its detection.  From the  particle-physics point of  view and in the  low-energy limit where  the  waterfall sector  has  decoupled  and the  $\\rho$-coupling neglected  for   simplicity,  the  $F_D$-term   hybrid  model  becomes identical to the  so-called Minimal Nonminimal Supersymmetric Standard Model (MNSSM)  in the decoupling  limit of a  large tadpole~\\cite{PP}. In   particular,    in   the    framework   of   WHI    discussed   in Section~\\ref{inflation},  the  coupling  $\\lambda$  can  be  sizeable, i.e.~$\\lambda \\sim 0.6$. In this  case, the Higgs phenomenology of the MSSM will  modify drastically, despite  the decoupling of  the singlet Higgs  states.  One  striking possibility  in  the MNSSM  is that  the charged  Higgs boson  $H^+$ could  be lighter  than the  SM-like Higgs boson~\\cite{PP2},    thus     pointing    to    particular    collider phenomenologies~\\cite{DP}.   However,   even  within  the  traditional scenario   of  CHI,   where  $\\kappa,   \\lambda  \\stackrel{<}{{}_\\sim} 10^{-2}$, the $F_D$-term hybrid model will favour particular benchmark scenarios  of the  MSSM.  For  example, if  $\\lambda \\gg  \\kappa$, the $F_D$-term hybrid model may account  for a possible large value of the $\\mu$-parameter.   Specifically, if  $\\lambda  = 4  \\kappa$, one  gets from~(\\ref{Send}) the hierarchy $\\mu \\approx 4 M_{\\rm SUSY}$, which is the so-called CPX  benchmark scenario~\\cite{CPX} describing maximal CP violation in the MSSM Higgs sector at low and moderate values of $\\tan \\beta$.   A  possible  natural  solution  to the  famous  cosmological  constant problem  is  expected to  provide  further  constraints  on the  model building of cosmologically viable models in future.  Nevertheless, the $F_D$-term hybrid  model presented in  this paper constitutes  a first attempt  towards the  formulation  of a  minimal Particle-Physics  and Cosmology  Standard  Model, whose  validity  could,  in principle,  be tested   in   laboratory  experiments   and   further  vindicated   by astronomical observations.   \\bigskip\\bigskip  \\subsection*{Acknowledgements}  We  thank Arjun  Berera, Rudnei  Ramos and  Antonio Riotto  for useful discussions. We also thank Constantine Pallis for collaboration in the early stages of this project.  AP~dedicates this work to the memory of Darwin  Chang, an invaluable  friend and  collaborator.  This  work is supported in part by the PPARC research grants: PPA/G/O/2002/00471 and PP/C504286/1.    \\newpage  \\def\\theequation{\\Alph{section}.\\arabic{equation}} \\begin{appendix}  \\setcounter{equation}{0}"
    },
    "0601/astro-ph0601056_arXiv.txt": {
        "abstract": "According to current models, gamma-ray bursts (GRBs) are produced when the energy carried by a relativistic outflow is dissipated and converted into radiation. The efficiency of this process, $\\epsilon_\\gamma$, is one of the critical factors in any GRB model. The X-ray afterglow light curves of {\\it Swift}\\/ GRBs show an early stage of flattish decay. This has been interpreted as reflecting energy injection. When combined with previous estimates, which have concluded that the kinetic energy of the late ($\\gtrsim 10\\;$hr) afterglow is comparable to the energy emitted in $\\gamma$-rays, this interpretation implies very high values of $\\epsilon_\\gamma$, corresponding to $\\gtrsim 90\\%$ of the initial energy being converted into $\\gamma$-rays. Such a high efficiency is hard to reconcile with most models, including in particular the popular internal-shocks model. We re-analyze the derivation of the kinetic energy from the afterglow X-ray flux and re-examine the resulting estimates of the efficiency. We confirm that, if the flattish decay arises from energy injection and the pre-{\\it Swift}\\/ broad-band estimates of the kinetic energy are correct, then $\\epsilon_\\gamma\\gtrsim 0.9$. We discuss various issues related to this result, including an alternative interpretation of the light curve in terms of a two-component outflow model, which we apply to the X-ray observations of GRB 050315.  We point out, however, that another interpretation of the flattish decay --- a variable X-ray afterglow efficiency (e.g., due to a time dependence of afterglow shock microphysical parameters) --- is possible. We also show that direct estimates of the kinetic energy from the late X-ray afterglow flux are sensitive to the assumed values of the shock microphysical parameters and suggest that broad-band afterglow fits might have underestimated the kinetic energy (e.g., by overestimating the fraction of electrons that are accelerated to relativistic energies).  Either one of these possibilities implies a lower $\\gamma$-ray efficiency, and their joint effect could conceivably reduce the estimate of the typical $\\epsilon_\\gamma$ to a value in the range $\\sim 0.1-0.5$. ",
        "introduction": "\\label{sec:introduction}  Recent observations by the {\\it Swift}\\/ X-ray telescope have provided new information on the early behavior of the X-ray light curve of long-duration ($\\gtrsim 2\\; {\\rm s}$) gamma-ray burst (GRB) sources. Specifically, it was found \\citep{Nousek05} that the light curves of these sources have a generic shape consisting of three distinct power-law segments $\\propto t^{-\\alpha}$: an initial (at $t<t_{\\rm break,1}$, with $300\\; {\\rm s}\\lesssim t_{\\rm break,1} \\lesssim 500\\; {\\rm s}$) very steep decline with time $t$ (with a power-law index $\\alpha_1$ in the range $3\\lesssim \\alpha_1 \\lesssim 5$; see also \\citealt{Tagliaferri05} and \\citealt{Barthelmy05}); a subsequent (at $t_{\\rm break,1}<t<t_{\\rm break,2}$, with $10^3\\; {\\rm s}\\lesssim t_{\\rm break,2} \\lesssim 10^4\\; {\\rm s}$) very shallow decay ($0.2\\lesssim \\alpha_2 \\lesssim 0.8$); and a final steepening (at $t>t_{\\rm break,2}$) to the canonical power-law behavior ($1\\lesssim \\alpha_3 \\lesssim 1.5$) that was known from pre-{\\it Swift}\\/ observations.  \\citet{Nousek05} already recognized that these results have direct consequences to the question of the $\\gamma$-ray emission efficiency in GRB sources. This question is important to our understanding of the basic prompt-emission mechanism. In the currently accepted interpretation \\citep[e.g.,][]{P99,P04}, the $\\gamma$-rays originate in a relativistic jet that is launched from the vicinity of a newly born neutron star or stellar-mass black hole. In the simplest picture, a fraction $\\epsilon_\\gamma$ of the energy injected at the source is given to the prompt radiation, with the remaining fraction ($1-\\epsilon_\\gamma$) ending up as kinetic energy of ambient gas that is swept up by a forward shock and then mostly radiated as early afterglow emission. One attractive mechanism for explaining the prompt emission characteristics invokes internal shocks that are generated when the ejecta have a nonuniform distribution of Lorentz factors, which results in outflowing ``shells'' colliding with each other at large distances from the source.  Pre-{\\it Swift}\\/ observations, based on measurements of the $\\gamma$-ray fluence and of the ``late'' ($\\gtrsim 10\\; {\\rm hr}$) afterglow emission, have implied (when interpreted in the context of the basic jet model) comparable (and narrowly clustered) values for the radiated $\\gamma$-ray energy and the kinetic energy feeding the afterglow emission, i.e., $\\epsilon_\\gamma \\sim 0.5$ \\citep[e.g.,][]{F01,PK01a,PK01b,PK02,BKF03,BFK03,Y03}. This result is seemingly problematic for the internal-shocks model, for which an order-of-magnitude lower value for $\\epsilon_\\gamma$ is a more natural expectation \\citep{KPS97,DM98,K99,GSW01}. This apparent difficulty could in principle be overcome if the ejected shells have a highly nonuniform distribution of Lorentz factors \\citep[e.g.,][]{B00,KS01}. However, to fit the data the shells must also satisfy a number of other restrictive conditions, which reduces the attractiveness of this interpretation (see \\S~\\ref{sec:conclusion}). An alternative proposal was made by \\citet{PKG05}, who argued that if the jet consists of an ultra-relativistic narrow component (from which the prompt emission originates) and a moderately relativistic wide component, with the latter having a higher kinetic energy and the former a higher kinetic energy per unit solid angle, then the wide component would dominate the late afterglow emission and the $\\gamma$-ray radiative efficiency of the narrow component could be significantly lower than the value of $\\epsilon_\\gamma$ inferred under the assumption of a single-component jet. As explained in \\citet{PKG05}, this proposal was motivated by observational indications of the presence of two components in the late afterglow light curves of several GRB sources and by the predictions of certain GRB source models.  The earliest (steepest) segment of the afterglow light curve is most naturally explained as radiation at large angles to our line of sight corresponding either to the prompt emission \\citep{KP00a} or to emission from the reverse shock that is driven into the ejecta \\citep{Kobayashi05}. This implies that the early afterglow emission is much weaker than what would be expected on the basis of an extrapolation from the late-afterglow data.  This behavior was interpreted by \\citet{Nousek05} as an indication of an even higher $\\gamma$-ray emission efficiency, typically $\\epsilon_\\gamma \\sim 0.9$. Such a high efficiency could render the internal-shocks model untenable.  Our primary goal is to evaluate the prompt-emission efficiency as accurately and systematically as possible on the basis of current data. For this purpose we re-derive in \\S~\\ref{sec:expressions} expressions that explicitly relate $\\epsilon_\\gamma$ to observable quantities. In particular, we express $\\epsilon_\\gamma$ in terms of the product $\\kappa f$ of two parameters, one ($\\kappa$) encapsulating information that could be obtained by pre-{\\it Swift}\\/ observations, and the other ($f$) representing early-time data obtained in {\\it Swift} measurements. In \\S~\\ref{sec:ek_est} we re-examine the estimates of the kinetic energy during the afterglow phase as inferred from the X-ray flux and present a new general formulation, correcting errors that have propagated in the literature and have generally led to an underestimate of the kinetic energy. We then analyze both pre-{\\it Swift}\\/ (\\S~\\ref{sec:pre_swift}) and {\\it Swift} (\\S~\\ref{sec:swift}) data in a uniform manner in the context of this formalism.  We argue that the kinetic energy estimates remain subject to considerable uncertainties.  In particular, while the simple analysis of a large number of bursts (using ``typical'' values for the microphysical parameters) yields rather high values for the kinetic energy and hence a low inferred $\\gamma$-ray efficiency, a multiwavelength analysis (which determines the microphysical parameters from the fit to the data) of a small subset suggests that the kinetic energy is lower and hence the inferred value of $\\epsilon_\\gamma$ is higher. Some caveats to this analysis are considered in \\S~\\ref{sec:assess}.   The suggestion that the {\\it Swift}\\/ observations imply a higher value of $\\epsilon_\\gamma$ in comparison with the pre-{\\it Swift} results is based on the interpretation of the flattish segment of the X-ray light curve as reflecting an increase in the kinetic energy of the forward shock during the early stages of the afterglow. In this picture, the kinetic energy just after the prompt-emission phase was significantly lower than the kinetic energy estimated from the later stages of the afterglow (the pre-{\\it Swift}\\/ results). One conceivable way of avoiding the need to increase the estimate of $\\epsilon_\\gamma$ in light of the new {\\it Swift}\\/ observations is through a time evolution (specifically, an increase with $t$) of the X-ray afterglow emission efficiency $\\epsilon_{\\rm X}$. Such a behavior could in principle also account for the flattish segment of the light curve and eliminate the need to invoke an increase in the shock kinetic energy. If, in addition, the value of the afterglow kinetic energy at late times were underestimated by the broad-band fits to pre-{\\it Swift}\\/ GRB afterglows, which could be the case if only a fraction $\\xi_e<1$ of the electrons were accelerated to relativistic energies in the afterglow shock \\citep[see][]{EW05}, then the typical afterglow efficiency would be further reduced (to a value as low as $\\epsilon_\\gamma\\sim 0.1$ if $\\xi_e \\sim 0.1$), which might reconcile the new {\\it Swift}\\/ data with the comparatively low efficiencies expected in the internal-shocks model. We discuss these issues in \\S~\\ref{sec:epsilon_X}.  A two-component jet model with the characteristics required for reducing the inferred $\\gamma$-ray emission efficiency is evidently disfavored by the {\\it Swift} data, but a model of this type with different parameters could provide an alternative interpretation of the flattish shape of the light curve between $t_{\\rm break,1}$ and $t_{\\rm break,2}$. We elaborate on these matters in \\S~\\ref{sec:2comp}, where we also present a tentative fit to the X-ray light curve of the {\\it Swift}\\/ source GRB~050315 in the context of this scenario.  Our conclusions on the physical implications of the early X-ray light-curve observations of GRB sources are presented in \\S~\\ref{sec:conclusion}. ",
        "conclusions": "\\label{sec:conclusion}  We have shown that the $\\gamma$-ray efficiency implied by the {\\it Swift}\\/ observations is model-dependent and can vary over a wide range (from typical values of $\\epsilon_\\gamma\\sim 0.9$ or higher to $\\epsilon_\\gamma\\sim 0.1$ or lower) depending on the adopted model assumptions. The $\\gamma$-ray efficiency has been expressed in terms of observable quantities (see eqs. [\\ref{obs_ratio}] and [\\ref{eps_g_fk1}]), namely $\\kappa$ and $f$, where $\\kappa$ relates the $\\gamma$-ray emission to the late-time afterglow emission (and was therefore available in the pre-{\\it Swift}\\/ era) and $f$ relates the early- and late-time afterglow energies (and therefore became available only with the launching {\\it Swift}). We have shown that there is a large uncertainty in the values of both $\\kappa$ and $f$, which translates into a corresponding uncertainty in the value of the $\\gamma$-ray efficiency, $\\epsilon_\\gamma$.  In the following discussion we make the conservative assumption that we observe most of the emitted energy in $\\gamma$-rays ($\\epsilon_{\\rm obs,GRB}\\approx 1$); if a significant fraction of the total radiated energy is emitted outside of the observed photon energy range (e.g., at higher energies) then this would increase the required value of $\\epsilon_\\gamma$ (see equation [\\ref{eps_g_fk1}]).  The kinetic energy of the afterglow shock has been estimated in the pre-{\\it Swift}\\/ era using broad-band afterglow fits for a small number of GRB sources that had the best available broad-band afterglow data, yielding a typical value of $\\kappa\\sim 1$ with a large scatter of almost an order of magnitude (see \\S~\\ref{sec:pre_swift}). Substituting the values of the microphysical parameters ($\\epsilon_e$, $\\epsilon_B$ and $p$) that were derived from these fits into our equations for $E_{\\rm k,iso}$ generally yielded somewhat lower values of $\\kappa$. Finally, using our equations with fiducial values of the microphysical parameters ($\\epsilon_e=0.1$, $\\epsilon_B=0.01$, and $p=2.2$) gives a typical value of $\\kappa\\sim 0.1-0.2$, both for the pre-{\\it Swift}\\/ and the {\\it Swift}\\/ GRB samples that we use, with a moderate scatter. Specifically, for our pre-{\\it Swift}\\/ ({\\it Swift}) sample, $\\langle\\log_{10}\\kappa\\rangle = -0.75$ ($-0.82$) corresponding to $\\kappa = 0.18$ ($0.15$) and $\\sigma_{\\log_{10}\\kappa}= 0.60$ ($0.63$). Obviously, the choice of fiducial values for the microphysical parameters is somewhat arbitrary and affects the resulting value of $\\kappa$. Higher values of the microphysical parameters $\\epsilon_e$ and $\\epsilon_B$ (e.g., $\\epsilon_e \\approx 0.3$, $\\epsilon_B \\approx 0.08$) are required in order to obtain an average value of $\\log\\kappa$, using our equations, similar to that derived from pre-{\\it Swift}\\/ broad-band afterglow fits.  Altogether, there is almost an order-of-magnitude uncertainty in the typical value of $\\kappa$ (which ranges from $\\sim 0.1-0.2$ to $\\sim 1$).  Even if we adopt the high typical values of $\\kappa$ ($\\sim 1$) inferred from pre-{\\it Swift}\\/ afterglow fits, it is important to keep in mind that these values have been estimated on the basis of the standard assumptions of afterglow theory. Changing these assumptions would modify the inferred value of $\\kappa$ for the same fits. For example, as pointed out by \\citet{EW05}, if only a fraction $\\xi_e<1$ of the electrons are accelerated to relativistic energies, then there is a degeneracy where the same observable quantities are obtained for $\\epsilon_e\\to\\xi_e\\epsilon_e$, $\\epsilon_B\\to\\xi_e\\epsilon_B$, $n_{\\rm ext}\\to\\xi_e^{-1}n_{\\rm ext}$, and $E_{\\rm k,iso}\\to\\xi_e^{-1}E_{\\rm k,iso}$. Since this increases the inferred value of $E_{\\rm k,iso}$ by a factor of $\\xi_e^{-1}$, $\\kappa$ is reduced by the same factor in comparison with the estimate from the standard theory (which uses $\\xi_e=1$).  An alternative way of reducing the inferred value of $\\epsilon_\\gamma$ was proposed by \\citet{PKG05} in the context of the pre-{\\it Swift}\\/ observations. Specifically, they considered a two-componet outflow model with parameters that effectively corresponded to the parameter $f$ having a value $< 1$. This parameter choice appears to be inconsistent with {\\it Swift}'s detection of an early flattish decay phase in the X-ray light curve, which, when interpreted in the context of the standard afterglow theory as arising from a gradual increase with time of $E_{\\rm k,iso}$, typically implies $f\\sim 10$ (and, in some cases, values of $f$ that are as high as $\\sim 10^2-10^3$). It is worth noting in this connection that the existence of a two-component GRB jet model can be plausibly expected on various theoretical grounds and has been suggested independently on the basis of fits to several pre-{\\it Swift}\\/ afterglows (see discussion in \\citealt{PKG05}). Furthermore, the {\\it Swift}\\/ observations by no means rule out this model, although they can be used to constrain its parameters. We have illustrated this fact through the fit to the X-ray light curve of GRB 050315 that we performed in \\S~\\ref{sec:2comp} within the framework of this model. This fit has yielded plausible ranges for the kinetic energies and opening angles of the two components as well for as the Lorentz factor of the dominant (wide) component. A key conclusion from this fit is that the kinetic energy of the wide component is much larger than that of the narrow one ($E_{\\rm k,w}/E_{\\rm k,n} \\sim 10^2$). Physically, the narrow and wide components could conceivably correspond to a baryon-poor, black-hole--driven outflow \\citep{LE93,LE03} and a baryon-rich, disk-driven outflow \\citep{VK03a,VK03b}, respectively, although this remains to be demonstrated. We also note that the two-component jet parameters derived in \\S~\\ref{sec:2comp} were based on standard assumptions; they could change if the underlying assumptions (involving, for example, the values and time constancy of the microphysical parameters) were altered.  In conjunction with the $\\kappa\\sim 1$ estimates of the standard broad-band afterglow fits, values of $f \\gtrsim 10$ imply $\\gamma$-ray radiative efficiencies $\\epsilon_\\gamma\\gtrsim 0.9$.  Such high efficiencies would be essentially impossible to achieve in any scheme, such as the internal-shocks model, that tapped the kinetic energy of the outflow for radiation. An alternative possibility that has been discussed in the literature is the direct transfer of Poynting flux \\citep[which evidently is also a major contributor to the flow acceleration --- e.g.,][]{DS02,VK03a} into nonthermal radiation \\citep[e.g.,][]{U94,T94}. It is at present unclear how to assess the efficiency of this process. There are two generic possibilities: dissipative fronts and magnetic reconnection sites. The first option corresponds to overtaking collisions of magnetically dominated relativistic streams and is not expected to result in high radiative efficiencies \\citep[e.g.,][]{RL97,LvP97}. The second case would require magnetic field orientation reversals and would most naturally arise in a pulsar-type outflow from a rapidly rotating neutron star \\citep[e.g.,][]{SDD01}. In this case it is, however, still unclear whether radiative efficiencies $\\gtrsim 0.9$ could be attained even under the most favorable assumptions about the field reconnection rate \\citep{DS02}, and it has in fact been suggested that the reconnection rate might be self-limiting \\citep{LK01}.  An ``intermediate'' situation could prevail if $f \\sim 10$ but $\\kappa\\sim 0.1$, reflecting the possibility that $\\kappa$ was overestimated by the pre-{\\it Swift}\\/ afterglow fits (perhaps because some of the assumptions of the standard theory do not hold --- e.g., $\\xi_e\\sim 0.1$ rather than $\\xi_e=1$). Alternatively, $\\kappa$ could be $\\sim 1$ but $f=1$, corresponding to the early flattish decay phase reflecting an increase with time of the X-ray afterglow efficiency $\\epsilon_{\\rm X}$ (due, e.g., to $p$ being $<2$ or to an increase with time of $\\epsilon_e$ or $\\epsilon_B$; see \\S~\\ref{sec:epsilon_X}) rather than an early increase in $E_{\\rm k,iso}$. In either one of these cases the inferred $\\gamma$-ray radiative efficiency would be reduced to $\\epsilon_\\gamma\\sim 0.5$. Although this value is less extreme than the estimate discussed in the preceding paragraph, it is worth noting that it is still fairly restrictive for the internal-shocks model, in which it could potentially be attained only if {\\it all}\\/ of the following conditions (already summarized in \\citealt{PKG05}) are satisfied: (1) the ratio between the maximum and minimum initial Lorentz factors of the ejected shells is large enough ($\\gtrsim 10$); (2) the distribution of initial Lorentz factors is sufficiently nonuniform; (3) the shells are approximately of equal mass and their number is large enough ($\\gtrsim 30$), and (4) the fraction of the dissipated energy that is deposited in electrons and then radiated away is sufficiently high ($\\epsilon_{e,\\rm GRB}\\epsilon_{\\rm rad,GRB}\\gtrsim 0.5$), with a similar constraint applying to the fraction $\\epsilon_{\\rm obs,GRB}$ of the radiated energy that is emitted as the observed $\\gamma$-rays (see \\citealt{B00} and \\citealt{KS01}).  Only if {\\it both}\\/ $f\\sim 1$ and $\\kappa\\sim 0.1$ were satisfied (which could occur, for example, if $\\epsilon_e$ or $\\epsilon_B$ initially increased with time and $\\xi_e$ were $\\sim 0.1$) would the inferred value of $\\epsilon_\\gamma$ drop to $\\sim 0.1$ and be compatible with the values that are expected to arise under less constrained circumstances in the internal-shocks model."
    },
    "0601/astro-ph0601260_arXiv.txt": {
        "abstract": "Methanol masers which are traditionally divided into two classes provide possibility to study important parts of the star forming regions: Class~II masers trace vicinities of the massive YSOs while class~I masers are likely to trace more distant parts of the outflows where newer stars can form.  There are many methanol transitions which produce observed masers. This allows to use pumping analysis for estimation of the physical parameters in the maser formation regions and its environment, for the study of their evolution. Extensive surveys in different masing transitions allow to conclude on the values of the temperatures, densities, dust properties, etc. in the bulk of masing regions. Variability of the brightest masers is monitored during several years. In some cases it is probably caused by the changes of the dust temperature which follow variations in the brightness of the central YSO reflecting the character of the accretion process.  A unified catalogue of the class II methanol masers consisting of more than 500 objects is compiled. Analysis of the data shows that: physical conditions within the usual maser source vary considerably; maser brightness is determined by parameters of some distinguished part of the object - maser formation region; class II methanol masers are formed not within the outflows but in the regions affected by their propagation.  It is shown that the \"near\" solutions for the kinematic distances to the sources can be used for statistical analysis. The luminosity function of the 6.7~GHz methanol masers is constructed. It is shown that improvement of the sensitivity of surveys can increase number of detected maser sources considerably.  The distribution of class II methanol masers in the Galaxy is constructed on the basis of estimated kinematic distances. It is shown that most of the sources are located in the Molecular Ring and that the dependence of the number of sources on the distance from the Galactic Center has significant peaks at the positions corresponding to the spiral arms.  A survey of CS(2-1) line emission tracing dense gas is performed at Mopra toward the positions of the brightest class II methanol masers. Velocity correlations between the maser and CS lines are analyzed. It is shown that the sources with l from 320 to 350 deg in which the masers are relatively blue-shifted, form a group which is located in the region of the Scutum-Centaurus spiral arm. This can reflect existence of a grand design, i.e., grouping of the sources with similar peculiarity of morphology or evolutionary stage of the massive star forming regions. ",
        "introduction": "Methanol masers are traditionally divided into two classes. Although masers of both classes often co-exist in the same star-forming region, they are usually seen apart from each other. Class~II masers are found in the vicinity of the massive YSOs while class~I masers are believed to trace distant parts of the outflows from these YSOs.  Class II methanol masers display strong emission in transitions at 6.6 and 12.1 GHz (\\cite{men91}). More than 2 dozens of weaker class II maser transitions are detected in several sources (see, e.g., \\cite{sob02}). At present there exist 3 unified catalogues of class II methanol masers containing more than 500 objects (\\cite{mal03}, \\cite{xu03} and \\cite{pest05}) and references to the initial data. Distinguishing feature of the \\cite{mal03} catalogue is that it contains data on several maser transitions and cross-references to the data on molecular shock tracers. Most of the class II maser sources were found toward positions of the IRAS sources while the blind surveys brought some detections in the places without other evidences of star formation. At present class II masers were found only in the high mass star forming regions (\\cite{min03}).  Class I methanol masers are less strong and widespread. Most of them were found in vicinities of class II maser sources (see, e.g., \\cite{sly94} and \\cite{ell05}). There were extensive searches (e.g., \\cite{sly94} and \\cite{val00}) for class I sources but no unified catalogues of these objects do exist and no blind surveys were performed yet. ",
        "conclusions": "\\label{sec:concl}  We have shown that the methanol masers represent a powerful diagnostic tool for the studies of the nature of the massive star formation regions and their distribution in the Galaxy."
    },
    "0601/astro-ph0601110_arXiv.txt": {
        "abstract": "Stellar models which incorporate simple diffusion or shear induced mixing are used to describe canonical extra mixing in low mass red giants of low and solar metallicity.  These models are able to simultaneously explain the observed Li and CN abundance changes along upper red giant branch (RGB) in field low-metallicity stars and match photometry, rotation and \\cc\\ ratios for stars in the old open cluster M67.  The shear mixing model requires that main sequence (MS) progenitors of upper RGB stars possessed rapidly rotating radiative cores and that specific angular momentum was conserved in each of their mass shells during their evolution.  We surmise that solar-type stars will not experience canonical extra mixing on the RGB because their more efficient MS spin-down resulted in solid-body rotation, as revealed by helioseismological data for the Sun.  Thus, RGB stars in the old, high metallicity cluster NGC\\,6791 should show no evidence for mixing in their \\cc\\ ratios.  We develop the idea that canonical extra mixing in a giant component of a binary system may be switched to its enhanced mode with much faster and somewhat deeper mixing as a result of the giant's tidal spin-up.  This scenario can explain photometric and composition peculiarities of RS CVn binaries.  The tidally enforced enhanced extra mixing might contribute to the star-to-star abundance variations of O, Na and Al in globular clusters. This idea may be tested with observations of \\cc\\ ratios and CN abundances in RS CVn binaries. ",
        "introduction": "\\label{sec:intro}  Standard stellar evolution models assume that mixing only occurs in convective regions of stars.  However, there is convincing evidence that a majority of low-mass ($0.8\\la M/\\msol\\la 2$) stars are experiencing nonconvective {\\it extra mixing} on the upper red giant branch (RGB) (e.g., \\citealt{sm79,lea86,sea86,vs88,st92,ch94,ch95,chea98,chdn98,grea00,bea01,gea02,dv03a,sm03,sh03}). These red giants have a helium electron-degenerate core, a hydrogen burning shell (HBS) atop the core, and a thin radiative zone between the HBS and the bottom of the convective envelope (BCE) (Fig.~\\ref{fig:f1}).  The most convincing observational evidence of mixing in red giants is the evolution of their surface carbon isotopic ratio \\cc\\ and abundances $\\eli$, [C/Fe] and [N/Fe]\\footnote{We use the standard spectroscopic notations: [A/B]\\,=\\,$\\log\\,(N({\\rm A})/N({\\rm B}))-\\log\\,(N({\\rm A})/N({\\rm B}))_\\odot$, where $N({\\rm A})$ and $N({\\rm B})$ are number densities of the nuclides A and B; $\\eli = \\log (N(^7\\mbox{Li})/N(\\mbox{H})) + 12$.}  with increasing luminosity $L$ along the RGB ({\\it circles} in Fig.~\\ref{fig:f2}). The first abundance changes that occur on the subgiant and lower red giant branch ($0.6\\la\\log L/L_\\odot\\la 1.0$) are produced by the deepening BCE during the standard first dredge-up episode (\\citealt{i67}). Here, the convective envelope engulfs the stellar layers in which the CN-cycle had altered the relative abundances of $^{12}$C, $^{13}$C and N while on the main sequence (MS). Convection also dilutes the envelope abundance of Li that has survived only in a thin subsurface layer where the temperature was less than $\\sim$\\,$2.5\\times 10^6$\\,K. At the end of the first dredge-up, the BCE stops to deepen and begins to retreat. It leaves a discontinuity of the chemical composition at the level of its deepest penetration.  The evolutionary changes of the surface element abundances resume precisely at the moment when the HBS, advancing in mass inside red giants, erases the composition discontinuity left behind by the BCE. The commonly accepted interpretation of this is based on the following assumptions: (1) some extra mixing is present in the red giants' radiative zones; (2) it connects the BCE with the layers adjacent to the HBS (Fig.~\\ref{fig:f1}), where the CN-cycle is at work and Li is strongly depleted, thus dredging up stellar material enriched in N and deficient in C and Li; (3) at lower luminosities, extra mixing is shielded from reaching the vicinity of the HBS by a large gradient of the mean molecular weight (a $\\mu$-gradient barrier) built up by the composition discontinuity (\\citealt{chea98,dv03a}).  When the HBS meets and erases the $\\mu$-gradient barrier, the evolution of red giants slows down for a while. This results in prominent bumps in the differential luminosity functions of globular clusters, which have really been observed (e.g., \\citealt{zea99,mea02,cea02,sea02,rea03}).  Therefore, the luminosity at which extra mixing starts to manifest itself (at $\\log L/L_\\odot\\approx 1.8$\\,--\\,2.0 in Fig.~\\ref{fig:f2}) is called ``the bump luminosity''. It divides the RGB into its lower and upper part. In stars with $M\\ga 2\\,\\msol$, the HBS fails to cross the $\\mu$-gradient barrier before the core helium ignition is triggered. In concordance with this, no manifestations of extra mixing have been reported in intermediate-mass red giants (\\citealt{g89}).  In the absence of a reliable theory of extra mixing in upper RGB stars, semi-empirical models are worth applying. In particular, a simple diffusion model has shown that in the majority of upper RGB stars the depth and rate of extra mixing do not seem to vary greatly from star to star.  According to \\cite{dv03a}, these quantities can be parameterized by any pair of correlated values within the close limits specified by $\\dlt\\approx 0.19$ and $\\dm\\approx 4\\times 10^8$\\cs, to $\\dlt\\approx 0.22$ and $\\dm\\approx 8\\times 10^8$\\cs.  Here, $\\dlt$ is a difference between the logarithms of temperature at the base of the HBS and at a specified maximum depth of extra mixing $\\mstar_{\\rm mix}$ (Fig.~\\ref{fig:f1}), and $\\dm$ is a constant diffusion coefficient.  Given that the majority of metal-poor upper RGB stars, both in the field and in globular clusters, experience extra mixing with almost the same depth and rate, \\cite{dv03a} have proposed to call this universal process ``canonical extra mixing''. Although its physical mechanism is not identified yet, the phenomenon of Li-rich K-giants and its interpretation are likely to support a rotational mechanism. Indeed, most of the Li-rich giants are located above the bump luminosity (\\citealt{chb00}). Their proportion among rapid rotators ($v\\sin\\,i\\geq 8$\\,\\kms) is $\\sim$\\,50\\%. Compare this to $\\sim$\\,2\\% of Li-rich stars among the much more common slowly rotating ($v\\sin\\,i\\la 1$\\,\\kms) K-giants (\\citealt{dea02}). Because for extra mixing driven by kinetic energy of rotation $\\dm\\propto v^2$, a tenfold increase of $v$ would enhance $\\dm$ from its canonical value of $\\sim 10^9$\\,\\cs\\ to $\\sim 10^{11}$\\,\\cs. \\cite{dh04} have shown that precisely such enhancement of $\\dm$ is required to produce Li in K-giants via the $^7$Be-mechanism (\\citealt{cf71}). Therefore, following \\cite{dv03a}, we hypothesize that sometimes, when an upper RGB star gets spun up by an external source of angular momentum, the presumably rotational extra mixing in its radiative zone may switch from its canonical mode to {\\it an enhanced mode} with much faster and somewhat deeper mixing.  \\cite{dw00} have proposed that single Li-rich giants might have been spun up by engulfed massive planets.  Alternatively, a red giant can be spun up by a close stellar companion.  In this case, its rotation is accelerated by a drag force originating from viscous dissipation of kinetic energy of a tidal hump raised in its convective envelope by a companion star. As a result of this interaction, spin and orbital rotation of the red giant get synchronized, and its initially elliptical orbit can become circular. Fig.~\\ref{fig:f3} illustrates tidal synchronization in real stars. Here, we have plotted observational data of \\cite{dmea02} on projected rotational velocities $v\\sin i$ and orbital periods $P$ of G and K giant components of field binaries with available orbital parameters.  The values of $v\\sin i$ were multiplied by the factor $4/\\pi\\approx 1.27$ that takes into account a random orientation of the stars' axes of rotation.  {\\it Filled circles} are binary systems with a circular or nearly circular orbit (those with eccentricities $e\\leq 0.10$). Systems with eccentric orbits ($e > 0.10$) are represented by {\\it open circles}.  A set of theoretical curves is constructed by us under the assumption that $\\Omega = \\omega$, where $\\Omega$ and $\\omega$ are the spin and orbital angular velocity of a red giant, for 4 values of the red giant's radius: $R/R_\\odot =$\\,5, 10, 20 and 40 ({\\it solid curves} from left to right). The tidal circularization time is longer than the synchronization time (e.g., \\citealt{hjea02}), therefore the giants represented by filled circles are most likely to have synchronized their rotation.  The goal of this paper is twofold. First, in \\sect{sec:single} we demonstrate that canonical extra mixing in solar-metallicity red giants in the open cluster M67 can be modeled using essentially the same parameters as those adjusted for metal-poor stars. After that, we attempt to model enhanced extra mixing in  red giant components of tidally locked (corotating) binaries that are spun up as a result of tidal synchronization between their spin and orbital rotation (\\sect{sec:binary}). In order to get a self-consistent picture of switching from canonical to enhanced extra mixing, in \\sect{sec:binary} we make assumptions and use $\\Omega$--profiles and parameters of extra mixing compatible with those discussed in \\sect{sec:single}. Finally, \\sect{sec:concl} contains a summary of our main conclusions along with a proposal of observational tests that could confirm or reject our models. ",
        "conclusions": "\\label{sec:concl}  In this paper, we have elaborated upon the ideas proposed by \\cite{dv03a} about canonical extra mixing in single upper RGB stars and enhanced extra mixing in low-mass red giants spun-up as a result of their tidal synchronization in close binary systems. In order to put as many as possible observational constraints on properties of extra mixing, we have supplemented the old data on the Li and CN abundance changes along upper RGBs in field low-metallicity stars with the new data on photometry, chemical peculiarities and rotational periods/velocities of stars from the solar-metallicity open cluster M67, globular clusters 47\\,Tuc, NGC\\,6397, NGC\\,6752, and RS CVn binaries.  We have confirmed the conclusions made in the earlier paper that the secular shear instability driven by differential rotation of the red giants' radiative zones can be considered as a promising physical mechanism for both modes of extra mixing while the turbulent diffusion coefficient (\\ref{eq:dv}) derived by \\cite{mm96} can be used to model them appropriately, provided that: {\\it (i)} unlike the Sun, low-mass MS progenitors of those red giants already possessed differentially rotating radiative cores; {\\it (ii)} the specific angular momentum was conserved in each mass shell, including convective regions, inside those stars during their entire evolution from the MS through the RGB tip; and {\\it (iii)} the diffusion coefficient (\\ref{eq:dv}) is taken with the enhancement factor $f_{\\rm v}\\approx 20$. For a binary red giant, we have additionally assumed that the tidal force brings to corotation only an upper part of its convective envelope. For the orbital and stellar parameters typical for the RS CVn binaries, this assumption results in spinning up of their radiative zones by a factor of $\\ga$\\,10.  Although we present some arguments in support of the made assumptions, it still seems unlikely that even our tidally enforced enhanced extra mixing with $\\dm\\sim 10^{11}$\\,\\cs\\ is not accompanied by a fast transport of angular momentum that would work toward flattening the $\\Omega$-profile in the radiative zone, thus reducing the mixing efficiency. From this point of view, our approach to modeling extra mixing in upper RGB stars is similar to the ``Maximal Mixing Approach'' recently used by \\cite{chea05}. The question ``how far is our approach from reality?'' can be answered in the usual way --- by comparing the theoretical predictions made by us with future observations. Here are some possible observational tests that can support or reject our models.  First, we propose to determine \\cc\\ ratios in upper RGB stars in the solar-metallicity open cluster NGC\\,6791 whose MS turn-off mass is $\\sim\\,1.1\\,\\msol$. If the Sun is not an exceptional case then, like the Sun, MS progenitors of the NGC\\,6791 red giants must have rotated as solid bodies. Therefore, we expect that canonical extra mixing in upper RGB stars in NGC\\,6791 should be inefficient. Hence, these stars should keep their post-first-dredge-up values of \\cc\\,$\\approx 25$ unchanged instead of having the values of \\cc\\,$\\approx 10$ similar to those measured in the M67 evolved stars that have definitely experienced canonical extra mixing.  Second, although we do believe that the depth of extra mixing is constrained by equation (\\ref{eq:dvmmix}), when computing its rate, we still cannot choose between the constant diffusion coefficient ($\\dm\\approx 10^9$\\,\\cs\\ for canonical and $\\dm\\approx 10^{11}$\\,\\cs\\ for enhanced extra mixing) and the turbulent diffusion coefficient (\\ref{eq:dvmix}) that is proportional to the radiative diffusivity $K$. The first coefficient better reproduces the Na--Li data for RS CVn binaries and, unlike the second one, it actually results in a steeper decline of the \\cc\\ ratio immediately after the bump luminosity as appears to be dictated by observations (e.g., compare the $\\log$\\,\\cc\\ panels in our Fig.~\\ref{fig:f2} and in Fig.~4 from \\cite{dv03a}). Fortunately, one probably can discriminate between the diffusion coefficients by measuring the Li abundance in RGB tip and clump stars. As we have predicted, in the case of $\\dm\\propto K$ these evolved stars should have values of $\\eli\\ga 1$ comparable with the post-first-dredge-up values or even exceeding them. Tentatively, this theoretical prediction seems to have been already confirmed by \\cite{pea01}. However, we think more observational work should be done in this direction.  Finally, we surmise that red giant components of the RS CVn binaries experience enhanced extra mixing caused by their tidal spin-up. We think that the large Na abundances found in these stars by \\cite{mea04} support this idea. We did not mention before that the same authors had also reported high Al abundances in the RS CVn binaries. It is interesting that a similar Na--Al correlation commonly inheres in the star-to-star abundance variations in globular clusters. One of its possible interpretations also employs a model of tidally enforced enhanced extra mixing in a binary red giant (\\S\\,4.2.2). If it could be possible to corroborate our hypothesis of enhanced extra mixing in the RS CVn binaries observationally that would also prove that a large amount of Al can be produced at the same physical conditions under which Na is synthesized (this actually requires higher rates for the MgAl-cycle reactions responsible for the Al production). We propose to determine \\cc\\ ratios and/or CN abundances in the RS CVn binaries to see whether they are consistent with those resulting from enhanced extra mixing."
    },
    "0601/astro-ph0601326_arXiv.txt": {
        "abstract": "{} {We analyze the impact on the Galactic nitrogen abundances of using a new set of low and intermediate mass star yields. These yields have a significant yield of primary nitrogen from intermediate mass stars.} {We use these yields as an input to a Galactic Chemical Evolution model and study the nitrogen abundances in the halo and in the disc, and compare them with models obtained using other yield sets and with a large amount of observational data.} { We find that, using these new yields, our model adequately reproduce the observed trends. In particular, these yields solve the historical problem of the evolution of nitrogen, giving the right level of relative abundance N/O by the production of a primary component in intermediate mass stars.  Moreover, using different evolutionary rates in each radial region of the Galaxy, we may explain the observed N dispersion.} {} ",
        "introduction": "Most elements are created in the interiors of stars by nucleosynthesis processes \\citep[see][for a review]{wall97}, starting with hydrogen and progressing toward heavy elements. These processes are called {\\sl primary production}.  Some elements, however, can be formed from nuclei heavier than hydrogen originally present in the star.  They are called {\\sl secondary}. This is the case of nitrogen, that can be created during the CNO cycle using seeds of original carbon and/or oxygen.  From a theoretical point of view, it has been considered that massive stars produce secondary nitrogen \\citep{pei87}, while low and intermediate mass (LIM) stars have mechanisms, like the third dredge-up and the Hot Bottom Burning processes, to produce both, primary and secondary nitrogen \\citep{edm78,all79}.  The third dredge-up event is a consequence of the thermal pulses in the star, and transport C and He to the outer layers. The Hot Bottom Burning occurs when the CNO cycle takes place at the base of the convective envelope.  Observationally, there are several open questions about the primary or secondary character of nitrogen that up to now remain unsolved. When N and O data are represented as log(N/O) {\\sl vs} log(O/H), including the galactic stars, H{\\sc ii} regions for the Milky Way Galaxy (MWG), external galaxies \\citep{gar95,gar99,vze98,izo99} and the high redshift data \\citep[][and references therein]{pet02,pro02,cen03}, a clear positive slope appears for abundances larger than $\\rm 12+log(O/H)=7.8-8$ dex which indicates a secondary behavior, but the plot shows a flat slope for low metallicities that can only be explained with a primary component of nitrogen. Taking into account that this flat slope occurs for low abundances, the first idea proposed, shared by some authors \\citep{pag79,dia86,dah95}, is that observations would be reproduced if the nitrogen ejected by massive stars would be primary, while intermediate mass stars might have both primary and secondary components.  Thus, some authors have tried to look for mechanisms that explain how massive stars could produce primary nitrogen. This is the case of \\citet{mey02} that have recently proposed rotation as a possible source of primary nitrogen, since low metallicity stars show a bigger rotation than high metallicity ones. \\cite{chi03-1} have used these yields in their chemical evolution models, concluding that they are only a lower limit for the primary nitrogen production since the Hot Bottom Burning is not considered in their calculation.  In fact, \\cite{chi05} find that an extra-- production of N in low metallicity massive stars by a large factor, between 40 and 200 along the mass range, is necessary to explain the data of very metal-poor halo stars since these yields do not produce a sufficient amount of primary N. Moreover, if the production of primary nitrogen would proceed from massive stars, the left side of the (N/O) {\\sl vs} (O/H) plot should not show any scatter.  Although some authors claim to observe \\citep{izo99,pil03} this no-scatter, recent observations from low metallicity objects \\citep{pet02,pro02,cen03,isr04,spi05} do show a clear dispersion.  \\citet{ser83} already claimed that a secondary production by intermediate mass stars must exist and suggested that the zero slope may be explained by two factors: 1) a delay in the ejection of N to the ISM due to the different mean-lifetimes of stars and 2) the gas infall effects. The advantage of taking a delay into account is that the great data scatter can be explained by considering different evolutionary states for each galaxy and, therefore, this possibility has been supported by a large number of authors \\citep{vila93,pil92,pil93,vze98b,hen00}. These last ones also include gas flows --infall or/and outflow--, and low efficiency for the star formation rate (the equivalent mechanism to produce a delay) in the low evolved regions, in order to reproduce the flat slope in the (N/O) vs (O/H) plot.  They conclude that the secondary production of nitrogen should dominate in high metallicity environments while the primary one should act at low metallicities.  Some new yields for LIM stars have been given in \\citet[][hereinafter Paper I]{gav05}, where they were adequately evaluated and calibrated by using them in a Galaxy chemical evolution model. It was shown that the results about C and O abundances adequately reproduce the Galactic and Solar Neighborhood data. The purpose of this work is to analyze the impact of these stellar yields on the nitrogen abundances. In particular, we check, using the same prescriptions of Paper I, if the contribution to N given by these yields for LIM stars is sufficient to justify the amount of primary nitrogen the observations point out.  We describe the yields in Section 2, analyzing in particular the primary and secondary components of the nitrogen production. In section 3 we describe briefly the chemical evolution model. Section 4 is devoted to the results, and the conclusions are presented in section 5. ",
        "conclusions": "Our conclusion can be summarized as follows:  \\begin{enumerate} \\item The primary component of nitrogen, necessary to explain the trend of N/O with O/H, may be mostly produced by LIM stars, and adequately fits all the data, including the observed dispersion. In this way the integrated yield produced by LIM stars must be directly proportional to Z, while the ratio $N_{P}/N_{tot}$ must increase for Z decreasing. A primary component, larger for lowest metallicities, has an important effect on explaining low abundance range data. \\item The dependence of the N yield on stellar mass would have a maximum around 5-6 $M_{\\odot}$, while the primary component shows other around 3.5-5 $M_{\\odot}$. This constraints the time where these contributions have important effects on the evolution. \\item The high dispersion on N/O data for low and high metallicity galactic regions may be explained with these yields as we have demonstrated with the radial regions of the disc models which have different star formation efficiencies. Our findings go in the same address than \\cite{hen00} and \\cite{pra03} using VG yields, but we claim that it is easier to reproduce the whole data set when BU yields are used. Our model for MWG with small efficiencies to form stars is consistent with data. \\item These results seem suggest that models with differences in the star formation histories for different types of galaxies, such as those calculated in \\cite{mol05}, might produce final abundances with high dispersion, in agreement with the observed one when dwarfs galaxies or DLA galaxies are included in a plot N/O-O/H, such as we will show in \\cite{mol05b}. \\item As \\cite{chi03-2}, we also support that the halo and the disc have different evolutions.  The set of stellar data [C/Fe], [N/Fe] and [N/C] may be divided into two trends. The first one is well reproduced by our disc models, while the second one is well fitted by our halo results. \\item Due to the primary N component of BU yields, and since the intermediate stars have short lifetimes, it is possible to produce high [N/Fe] abundances even at low metallicities which are in perfect agreement with the recent halo stars data obtained by \\cite{isr04,cay04,spi05}. \\item We claim that the fit of the whole set of data with only one model is not an easy task.  We may reproduce the observed trend with BU yields combined with yields from WW. In summary, our model BU reproduce reasonably well the whole CNO data set. \\end{enumerate}"
    },
    "0601/astro-ph0601489_arXiv.txt": {
        "abstract": "The non-Keplerian galactic rotational curves and the gravitational lensing data~\\cite{Persic:1995ru,Wittman:2000tc,Hoekstra:2003pn} strongly indicate a significant dark matter component in the universe. Moreover, these data can be combined to deduce the equation of state of dark matter \\cite{Faber:2005xc}. Yet, the existence of dark matter has been challenged following the tradition of critical scientific spirit. In the process, the theory of general relativity itself has been questioned and various modified theories of gravitation have been proposed~\\cite{Milgrom:1983pn,Sanders:2002pf,Bekenstein:2004ne}. Within the framework of the Einsteinian general relativity, here I make the observation that if the universe is described by a spatially flat Friedmann-Robertson-Walker (FRW) cosmology with Einsteinian cosmological constant then the resulting cosmology predicts a significant dark matter component in the universe. The phenomenologically motivated existence proof refrains from invoking the data on galactic rotational curves and gravitational lensing, but uses as input the age of the universe as deciphered from studies on globular clusters. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601440_arXiv.txt": {
        "abstract": "We present VLT/NACO SDI images of the very nearby star SCR 1845-6357 (hereafter SCR 1845).  SCR 1845 is a recently discovered \\citep{ham04} M8.5 star just 3.85 pc from the Sun \\citep{hen06}.  Using the capabilities of the unique SDI device, we discovered a substellar companion to SCR 1845 at a separation of 4.5 AU (1.170''$\\pm$0.003'' on the sky) and fainter by 3.57$\\pm$0.057 mag in the 1.575 $\\mu$m SDI filter. This substellar companion has an H magnitude of 13.16$^{+0.31}_{-0.26}$ (absolute H magnitude of 15.30$^{+0.31}_{-0.26}$), making it likely the brightest mid-T dwarf known. The unique Simultaneous Differential Imager (SDI) consists of 3 narrowband filters placed around the 1.6 $\\mu$m methane absorption feature characteristic of T-dwarfs (T$_{eff}$ $<$ 1200 K). The flux of the substellar companion drops by a factor of 2.7$\\pm$0.1 between the SDI F1(1.575 $\\mu$m) filter and the SDI F3(1.625 $\\mu$m) filter, consistent with strong methane absorption in a substellar companion. We estimate a spectral type of T5.5$\\pm$1 for the companion based on the strength of this methane break. The chances that this object is a background T dwarf are vanishing small -- and there is no isolated background T-dwarf in this part of the sky according to 2MASS. Thus, it is a bound companion, hereafter SCR 1845-6357B. For an age range of 100 Myr - 10 Gyr and spectral type range of T4.5-T6.5, we find a mass range of 9 - 65 M$_{Jup}$ for SCR 1845B from the \\citet{bar03} COND models. SCR 1845AB is the 24th closest stellar system to the Sun (at 3.85 pc); the only brown dwarf system closer to the Sun is the binary brown dwarf Eps Indi Ba-Bb (at 3.626 pc).  In addition, this is the first T-dwarf companion discovered around a low mass star. ",
        "introduction": "After decades of little change in the number of known stellar systems within 5 pc, numerous previously unknown low mass stars have recently been discovered in the solar neightborhood \\citep[e.g.~][]{sch03,hen04,ham04}. Because they are extremely nearby and intrinsically low luminosity, these objects are ideal targets to search for low mass companions, since even a close companion will appear reasonably separated on the sky. For example during commissioning of the Simultaneous Differential Imager (SDI) at the VLT \\citep{clo05a,len04}, one such object, $\\epsilon$ Indi B was resolved into a binary T dwarf \\citep{mcc03}. $\\epsilon$ Indi Ba,Bb are the closest brown dwarfs to the Earth (3.626$\\pm$0.009 pc). At a distance of 3.85$\\pm$0.02 pc \\citep{hen06}, the recently discovered M8.5 star SCR 1845-6357 is just slightly further away than $\\epsilon$ Indi Ba,Bb \\citep{ham04} and is the 24th closest stellar system from the Earth.  We report the discovery of SCR 1845-6357B (hereafter SCR1845B), a methane rich substellar companion to this star.  This companion object is the third closest brown dwarf to the Sun.  It is also the only example of a T dwarf companion to a low mass star and one of the tightest known brown dwarf companions to a star. ",
        "conclusions": "SCR 1845B is the brightest mid-T dwarf yet discovered. In addition, it is the first T dwarf companion found around a low mass star. At only $\\sim$4.5 AU from its primary, it is one of the tightest known brown dwarf companions to a star and is a further piece of evidence that the brown dwarf desert does not exist for companions to very low mass stars \\citep[][]{clo03,giz03}. Both the primary and secondary mass can be accurately measured within a decade."
    },
    "0601/astro-ph0601395_arXiv.txt": {
        "abstract": "{We have performed a search for variable stars in the dwarf irregular galaxy NGC\\,6822 using wide-field multi-epoch VI photometry down to a limiting magnitude $V$ $\\sim$ 22. Apart from the Cepheid variables in this galaxy already reported in an earlier paper by Pietrzynski et al. (2004), we have found 1019 \"non-periodic\" variable stars, 50 periodically variable stars with periods ranging from 0.12 to 66 days and 146 probably periodic variables. Twelve of these stars are eclipsing binaries and fifteen are likely new, low-amplitude Cepheids. Interestingly, seven of these Cepheid candidates have periods longer than 100 days, have very low amplitudes (less than  0.2 mag in $I$), and are very red. They could be young, massive Cepheids still embedded in dusty envelopes. The other objects span a huge range in colours and represent a mixture of different types of luminous variables. Many of the variables classified as non-periodic in the present study may turn out to be {\\it periodic} variables once a much longer time baseline will be available to study them.  We provide the catalogue of photometric parameters and show the atlas of light curves for the new variable stars. Our present catalogue is complementary to the one of Baldacci et al. (2005) which has focussed on very short-period and fainter variables in a subfield in NGC 6822. ",
        "introduction": "NGC\\,6822 is a relatively well studied dwarf irregular galaxy of the Local Group. In recent years its content of variable stars has been investigated by several authors, mainly with emphasis on distance indicators, like Cepheids and RR Lyrae stars, and short period variables. Antonello et al. (2002a, 2002b) discovered 130 variable stars in NGC\\,6822, 21 of them Cepheids, 18 other periodic variables and 91 irregular or semiregular variables. Their study was based on a 3-year monitoring of a small (3.8 x 3.8 arcminute) field of NGC\\,6822. In a series of papers, Clementini et al. (2003), Baldacci et al. (2004a, 2004b), and Baldacci et al. (2005), using deep  VLT images of a field of 6.8 x 6.8 arcminute, reported 390 candidate variables. They used data taken on 3 half-nights distributed over 5 days in August 2001. They classified the variables in Main Sequence Variables (MSV, 36 stars), Classical Instability Strip Variables (CISV, 160 stars) and Red Giant Branch Variables (RGBV, 66 stars). They also found 6 eclipsing binaries. As they noted, their data distribution was optimized for searching periods in the range of 0.2-0.7 days. On the other hand, Pietrzy{\\'n}ski et al. (2004, hereafter P04) monitored the whole galaxy in the $V$ and $I$ bands during 77 nights, spread over nearly one year, with the aid of mosaic images  (35 x 35 squared arcminutes) taken with the Warsaw 1.3 m telescope at Las Campanas Observatory, discovering 116 Cepheids with periods ranging from 1.7 to 124 days. These authors determined, using the reddening-free Wesenheit index, a distance modulus for NGC\\,6822 of 23.34 ($\\pm$ 0.04 statistical and $\\pm$ 0.05 systematic) mag. In the present paper, we use the photometric database provided by P04 to search for variable stars in NGC\\,6822. The investigation presented in this paper complements the previous ones since it is the first simultaneous monitoring of the whole galaxy during a time interval of almost one year, allowing us to detect and characterize the population of relatively long-period, bright variables from these images.  In Section 2, we give a review of the observations and explain our searching tools and selection methods. In Section 3, we show the results of our investigation, and discuss them in Section 4. Conclusions  and future prospects are outlined in Section 5. ",
        "conclusions": "In this paper we have continued our search for extragalactic variable stars making use of the image databases provided by the Araucaria Project (for variable stars detected in NGC\\,300 in the framework of this project, see Mennickent et al.\\,2004; for details on the Araucaria Project, see for instance Pietrzynski et al., 2002, and more recently Gieren et al., 2005). With this study, we have significantly contributed to the census and understanding of the population of bright variable stars towards the galaxy NGC\\,6822. We have presented data for more than 1200 variable stars, far exceeding the number of previously known objects of this kind  toward this galaxy.  Unfortunately, for many objects we cannot unambiguously determine their membership to NGC 6822, since contamination by foreground Milky Way stars is a strong problem in such surveys in Local Group galaxies. Our present survey is not an exception, especially due to the low galactic latitude of NGC\\,6822 ($b^{II}$ = -18$\\degr$). According to the Bahcall \\& Soneira galaxy model for star densities, we should expect  16.2 foreground stars per square arcminute with magnitudes in the range of 13 $ < V < $ 23 towards NGC\\,6822 (Ratnatunga \\& Bahcall 1985). This yields approximately 18\\% of contamination for the total number of stars detected in our survey. If we translate this figure to the detected variable stars, we would expect that 219 stars of our  1215 objects could be galactic foreground stars. True membership to NGC\\,6822 can only be established with spectroscopic follow-up surveys and, in some cases, with multicolor photometry (Massey 1998). In any case, our work has been exploratory; the detection of variable star candidates constitutes the first step towards a complete understanding of the variable star population in this galaxy. Our search was limited to $M_{V} < -2$, excluding pulsators like RR-Lyrae and $\\delta$ Scuti stars, but we should be able to detect $\\beta$ Cephei, Mira and semiregular variables, as well as RV Tau stars. Many of our targets could finally be classified as these kind of variables once follow-up observations become available. Further spectroscopic work and  continuous photometric monitoring are therefore encouraged to understand the nature of these variables, and to compare their properties with those of the corresponding populations in other galaxies.  An interesting byproduct of our investigation has been the discovery of a small population of stars which could be strongly reddened, small-amplitude Cepheids of extremely long periods. They seem to follow the Cepheid $PL$ relationship, but turn to be underluminous in the $V$ band. We suspect that these 7 objects might be indeed very young Cepheids which are still embedded in dusty envelopes. Further work will be necessary to corroborate this hypothesis."
    },
    "0601/physics0601011_arXiv.txt": {
        "abstract": "The irregular polarity reversals of the Earth's magnetic field have attracted much interest during the last decades. Despite the fact that recent numerical simulations of the geodynamo have shown nice polarity transitions, the very reason and the basic mechanism of reversals are far from being understood. Using a paradigmatic mean-field dynamo model with a spherically symmetric helical turbulence parameter $\\alpha$ we attribute the essential features of reversals to the magnetic field dynamics in the vicinity of an {\\it exceptional point} of the spectrum of the non-selfadjoint dynamo operator. At such exceptional (branch) points of square root type two real eigenvalues coalesce and continue as a complex conjugated pair of eigenvalues. Special focus is laid on the comparison of numerically computed time series with paleomagnetic observations. It is shown that the considered dynamo model with high supercriticality can explain the observed time scale and the asymmetric shape of reversals with a slow decay and a fast field recovery. ",
        "introduction": "We consider a simple mean-field dynamo model of $\\alpha^2$ type with a supposed spherically symmetric, isotropic helical turbulence parameter $\\alpha$ \\cite{KRRA}. The induction equation for the magnetic field $\\bf B$ reads \\begin{eqnarray} {\\dot{ \\bf{B}}} ={\\bf \\nabla} \\times (\\alpha {\\bf{B}}) + (\\mu_0 \\sigma)^{-1} \\Delta {\\bf{B}} \\; , \\end{eqnarray} with magnetic permeability $\\mu_0$ and electrical conductivity $\\sigma$. For the Earth's  core we will assume the diffusion time   $\\tau_{diff}=\\mu_0 \\sigma R^2$ to be $\\sim 200$ kyr.  \\begin{figure}[t] \\begin{center} \\unitlength=\\textwidth \\includegraphics[width=\\textwidth]{fig2.eps} \\caption{Explanation of the field dynamics for $C=7.237$ and $D=0$. (a) Half an anharmonic oscillation with nine selected instants $\\tau_i$, $i=1...9$. (b) Instantaneous growth rates resulting from the instantaneous $\\alpha(r,\\tau_i)$ profiles. (c) Instantaneous frequencies. (d) Instantaneous poloidal field $s(r,\\tau_i)$. (e) Profiles $\\alpha(r,\\tau_i)$, compared to the unquenched $\\alpha(r)$ (K). Note that the deformation  of $\\alpha(r)$ during a reversal is not very strong. (f) Details of (e) close to the origin. At the instants 6 and 7, the $\\alpha(r)$ profile comes close to the unquenched  $\\alpha(r)$ (K).} \\end{center} \\end{figure}  The divergence-free magnetic field  ${\\bf{B}}$ can be decomposed into a poloidal and a toroidal components, according to ${\\bf{B}}=-\\nabla \\times ({\\bf{r}} \\times \\nabla S)-{\\bf{r}} \\times \\nabla T $. The defining   scalars $S$ and $T$ are expanded in spherical harmonics of degree $l$ and order $m$ with expansion coefficients $s_{l,m}(r,\\tau)$ and $t_{l,m}(r,\\tau)$. For the envisioned spherically symmetric and isotropic $\\alpha^2$ dynamo problem, the induction equation decouples for each degree $l$ and order $m$ into the following pair of equations: \\begin{eqnarray} \\frac{\\partial s_l}{\\partial \\tau}&=& \\frac{1}{r}\\frac{\\partial^2}{\\partial r^2}(r s_l)-\\frac{l(l+1)}{r^2} s_l +\\alpha(r,\\tau) t_l \\; ,\\\\ \\frac{\\partial t_l}{\\partial \\tau}&=& \\frac{1}{r}\\frac{\\partial}{\\partial r}\\left( \\frac{\\partial}{\\partial r}(r t_l)-\\alpha(r,\\tau) \\frac{\\partial}{\\partial r}(r s_l) \\right) -\\frac{l(l+1)}{r^2} [t_l-\\alpha(r,\\tau) s_l] \\; . \\end{eqnarray} Since these equations are independent of the order $m$, we have skipped $m$ in the index of $s$ and $t$. The boundary conditions are $\\partial s_l/\\partial r |_{r=1}+{(l+1)} s_l(1)=t_l(1)=0$. In the following we consider only the dipole field with $l=1$.  \\begin{figure}[t] \\begin{center} \\unitlength=\\textwidth \\includegraphics[width=\\textwidth]{fig3.eps} \\caption{The role of exceptional points for the understanding of reversals. (a) Instantaneous growth rate curves for the kinematic profile and for the quenched $\\alpha(r,\\tau_i)$ profiles. The $E_i$ indicate, for each of the considered $\\alpha$ profiles, the exceptional point. (b) Growth rates for the profiles $\\alpha(r)= 1.916 \\, C^* \\, C  \\, (1-6 \\; r^2+5 \\; r^4)/(1+E_{mag}(r)/E_{mag,0})$ with $C=$6.78, 7.2, 10, 20, 50, 100, 200. The label ''7.2 (m)'' refers to the maximally quenched $\\alpha$ during the reversal. Note the remarkable move of the exceptional point well above the zero line and back to it (indicated by the thick dashed line).} \\end{center} \\end{figure}  We will focus on a  particular radial profile $\\alpha_{kin}(r)=1.916 \\; C  \\; (1-6 \\; r^2+5 \\; r^4)$, which is characterized by a  sign change along the radius (the factor 1.916 results from normalizing the radial average of $|\\alpha(r)|$ to the corresponding value for constant $\\alpha$). The motivation for this choice is that quite similar $\\alpha(r)$ profiles had been shown to exhibit oscillatory behaviour \\cite{OSZI}.  Saturation is ensured by assuming the kinematic profile $\\alpha_{kin}(r)$ to be algebraically quenched by the  magnetic field energy averaged over the angles which can be expressed in terms of $s_{1}(r,\\tau)$ and $t_1(r,\\tau)$. Note that this averaging over the angles represents a severe simplification, since in reality (even for an assumed spherically symmetric kinematic $\\alpha$) the energy dependent quenching would result in a breaking of the spherical symmetry.  In addition to this quenching, the $\\alpha(r)$ profiles are perturbed  by \"blobs\" of noise which are considered constant within a correlation time $\\tau_{corr}$. Physically, such a  noise term  can be understood as a consequence of changing boundary conditions for the flow in the outer core, but also as a substitute for the omitted influence of higher multipole modes on the dominant axial dipole mode.  In summary, the $\\alpha(r,\\tau)$ profile  entering Eqs. (2) and (3) is written as \\begin{eqnarray} \\alpha(r,\\tau)&=&C \\frac{1.916 \\; (1-6 \\; r^2+5 \\; r^4)}{1+ E^{-1}_{mag,0} \\left[ \\frac{2 s_1^2(r,\\tau)}{r^2}+ \\frac{1}{r^2}\\left( \\frac{\\partial (r s_1(r,\\tau))} {\\partial r} \\right)^2 +t_1^2(r,\\tau) \\right] }     \\nonumber\\\\ && +\\xi_1(\\tau) +\\xi_2(\\tau) \\; r^2 +\\xi_3(\\tau) \\; r^3+\\xi_4(\\tau) \\; r^4 \\; , \\end{eqnarray} where the noise correlation is given by $\\langle \\xi_i(\\tau) \\xi_j(\\tau+\\tau_1) \\rangle = D^2 (1-|\\tau_1|/\\tau_{corr}) \\Theta(1-|\\tau_1|/\\tau_{corr}) \\delta_{ij}$. $C$ is a normalized dynamo number, $D$ is the noise strength, and $E_{mag,0}$ is a constant measuring the mean magnetic field energy.  \\begin{figure}[t] \\begin{center} \\unitlength=\\textwidth \\includegraphics[width=\\textwidth]{fig4.eps} \\caption{Time series for some selected values of $C$ and $D$.     } \\end{center} \\end{figure} ",
        "conclusions": "Based on former results in \\cite{PRL} and \\cite{EPSL}, we have shown that a simple but strongly  supercritical $\\alpha^2$ dynamo model exhibits a number of features which are typical for Earth's magnetic field reversals, in particular an asymmetric shape  and correct timescales for the field decay and field recreation.  As it does not include the necessary  North-South asymmetry of $\\alpha$ we cannot claim that our model is an appropriate model of the geodynamo. However, recent papers by Giesecke et al. have shown that a) even in such realistic models $\\alpha$ may exhibit a sign change along the radius \\cite{GIES} and b) that such models can also give rise to reversals \\cite{GIESAN}. Therefore, it seems worthwhile to identify the indicated reversal scenario in this and in more complicated geodynamo models."
    },
    "0601/astro-ph0601506_arXiv.txt": {
        "abstract": "We report on ESO Very Large Telescope optical spectroscopy of 42 BL Lacertae objects of unknown redshift.  Nuclear emission lines were observed in 12 objects, while for another six we detected absorption features due to their host galaxy. The new high S/N spectra therefore allow us to measure the redshift of 18 sources. Five of the observed objects were reclassified either as stars or quasars, and one is of uncertain nature.  For the remaining 18 the optical spectra  appear without intrinsic features in spite of our ability to measure rather faint (EW $\\sim$0.1 \\AA) spectral lines. For the latter sources a lower limit to the redshift was set exploiting the very fact that the absorption lines of the host galaxy are undetected on the observed spectra. ",
        "introduction": "\\label{sec:intro}  BL Lac objects (hereinafter BL Lacs or BLL) are active galactic nuclei (AGN) characterized by luminous, rapidly variable UV--to--NIR non--thermal continuum emission and polarization, strong compact flat spectrum radio emission and superluminal motion. Similar properties are observed also in flat spectrum radio quasars (FSRQ) and the two types of active nuclei are often grouped together in the blazar class. From the spectroscopical point of view BL Lacs are characterized by quasi featureless optical spectra. In fact their spectra are often dominated by the non--thermal continuum that arises from the nucleus. To this emission it is superimposed a thermal contribution due to the stellar component of the host galaxy. Like in other AGN, emission lines  could be generated by fluorescence in clouds surrounding the central black hole. Moreover, as it happens for high z quasars in some cases absorption lines due to intervening gas in the halo of foreground galaxies can be observed in the spectra of BL Lacs and one can derive a lower limit to the redshift of the object. The detectability of spectral features depends on the brightness of the nuclear source: in fact during low brightness states, intrinsic absorption features can be more easily  revealed, while during high states one can better discover intervening absorption systems. Because of the strong contribution from the continuum the equivalent width (EW) of all these spectral features is often very small and their detection represents a challenging task.  In the past decade a number of projects were carried out to derive the redshift of BL Lac objects. Most of these works were based on optical spectra collected with 4 m class telescopes, and are therefore limited by relatively low signal-to-noise  ratio (S/N), low spectral resolution and limited wavelength range \\citep[e.g.][]{falomo93, stickel93, veron93, bade94,falomo94, falomo96, marcha96, drinkwater97,rgb, landt01, rector01, londish02, carangelo03, hook03}. Recently, however, some observations with 8 m class telescopes were carried out \\citep{heidt04,sowards05}. Despite these efforts, a significant fraction of known BL Lacs (e. g. 50 \\% in \\citet{veron03} catalogue) have still unknown redshift.   In order to improve the knowledge of the redshift of BL Lacs we  carried out a project to obtain optical spectra of sources with still unknown or uncertain redshift using the European Southern Observatory (ESO) 8-m Very Large Telescopes (VLT). This allows one to improve significantly the S/N of the spectra and therefore the capability to detect faint spectral features. A first report on this work, giving the redshift of 12 objects, has been presented by \\citet[][Paper I]{sbarufatti05a}, and here we refer on the results for the full sample of 42 observed sources.  The outline of this paper is the following. In section  \\ref{sec:sample} we give some characterization of the 42 observed objects. The observation and analysis procedures are described in section \\ref{sec:obsred}. In sections \\ref{sec:results} and \\ref{sec:notes} we report the results of our spectroscopic study. Finally in section \\ref{sec:discussion} a summary and conclusions of this study are given. Throughout this paper we adopted the following cosmological parameters: H$_0$= 70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{\\Lambda}$=0.7, $\\Omega_m$=0.3. ",
        "conclusions": "\\label{sec:discussion}  Out of 42 objects observed we confirm the BL Lac classification for 36 sources and for 18 of them we are able to measure/confirm the redshift. This information allows us to derive the luminosity of the objects.  The distribution in the V band luminosity-distance plane is indeed fully consistent with what observed for BL Lacs of known redshift in the combination of the \\citet{PG95} and the SS sample (see Fig. \\ref{fig:distz}). We note that the sources in this combined list are affected by the typical selection effect of incomplete, flux limited samples: the envelope of the objects follows in fact the expected behavior for sources with constant V magnitude, centered around V=18, with a spread of $\\sim$ 3 mag. The objects discussed here (filled circles and lower limits) follow the same distribution, with absolute magnitudes ranging between -21.5 and -27.5, slightly increasing with the redshift.  In 18 cases the optical spectra remain lineless in spite of the high  S/N of the obtained optical spectra. This indicates that if they are hosted by galaxies of standard luminosity they have likely very luminous or extremely beamed nuclei \\citep[see also][]{sbarufatti05b}. In the latter case one may expect to see the most extreme cases of relativistic beaming, making these sources ideal targets for milli-arcsec resolution radio observations. Alternatively, if the host galaxies were under-luminous, these objects could be rare  examples of dwarf galaxies hosting an AGN \\citep[see][]{sbarufatti05b}.  The high S/N of most of the optical spectra obtained at VLT represents a frontier for the determination of the redshift of BL Lacs with current instrumentation and further improvement of the issue will not be easy to get. In particular the four brightest objects (R$<$16: 0048--099, 1553+113, 1722+119, 2136-428) we observed at the VLT have values of \\EWmin \\ smaller than 0.25 \\AA. These objects belong to an interesting sub-population of BL Lac objects with extreme nuclear (and/or host) properties for which it is actually not possible to derive the intrinsic physical parameters. One  possibility is to image  the source when it is particularly faint in order to improve the detection of the host galaxy and to derive  an imaging redshift \\citep{sbarufatti05b}. Deep spectroscopic observations in the near-IR may also prove to be effective in the determination of the redshift considering this region of the spectrum is poorly known."
    },
    "0601/astro-ph0601620_arXiv.txt": {
        "abstract": "The recent detection with \\chandra\\ of two warm--hot intergalactic medium (WHIM) filaments toward Mrk 421 by Nicastro et al.~provides a measurement of the bulk of the ``missing baryons'' in the nearby universe.  Since Mrk 421 is a bright X-ray source, it is also frequently observed by the \\xmmnewton\\ Reflection Grating Spectrometer (RGS) for calibration purposes.  Using all available archived \\xmm\\ observations of this source with small pointing offsets ($<15$\\arcsec), we construct the highest--quality \\xmm\\ grating spectrum of Mrk 421 to date with a net exposure time (excluding periods of high background flux) of 437~ks and $\\sim 15000$ counts per resolution element at 21.6\\,\\AA, more than twice that of the \\chandra\\ spectrum.  Despite the long exposure time neither of the two intervening absorption systems is seen, though the upper limits derived are consistent with the \\chandra\\ equivalent width measurements.  This appears to result from (1) the larger number of narrow instrumental features caused by bad detector columns, (2) the degraded resolution of XMM/RGS as compared to the \\chandra/LETG, and (3) fixed pattern noise at $\\lambda \\ga 29$\\,\\AA. The non--detection of the WHIM absorbers by XMM is thus fully consistent with the \\chandra\\ measurement. ",
        "introduction": "Most of the baryons that comprise 4\\% of the mass--energy budget of the universe are found in the intergalactic medium (IGM), primarily appearing as the Lyman--alpha ``forest'' in high--redshift quasar spectra \\citep{weinberg97}.  At more recent times ($z\\la 2$) the process of structure formation has shock--heated the IGM to temperatures of $\\sim 10^{5-7}$\\,K, thus rendering the hydrogen nearly fully ionized and producing (at most) broad, extremely weak Lyman--alpha absorption \\citep[e.g.][]{sembach04,richter04}.  Known as the warm--hot intergalactic medium (WHIM), this phase is predicted to contain roughly half of the baryonic matter at low redshifts \\citep{cen99,dave01}. Its extremely low density (typically $\\delta\\sim 10-100$) precludes the detection of WHIM thermal or line emission with current facilities, so the only way to directly measure these ``missing'' baryons is through far--ultraviolet and X-ray spectroscopic measurements of absorption lines from highly ionized heavy elements \\citep{perna98,hellsten98,fangb02}.  Several early attempts to detect these intervening WHIM absorbers in X-rays \\citep[e.g.][]{fang02,mathur03,mckernan03} and more recent surveys \\citep{fang05} yielded only tentative detections at best.  Intervening Lyman--alpha \\citep{shull96,sembach04} and \\ion{O}{6} \\citep[e.g.][]{savage02} absorbers had also been seen in \\fuse\\ and HST quasar spectra, but their ionization states and possible galactic halo origins \\citep[e.g.][]{tumlinson05} are quite uncertain.  These uncertainties are largely mitigated with the recent detection by \\citet{nicastro05a,nicastro05b} of two X-ray absorption systems at $z=0.011$ and $z=0.027$ along the line of sight to the blazar Mrk 421.  These filaments account for a critical density fraction of $\\Omega_{\\rm WHIM}=0.032^{+0.042}_{-0.021}$, fully consistent with the mass of the missing baryons in the local universe (albeit with large uncertainties).  While future proposed missions such as \\emph{Constellation--X}, \\emph{XEUS}, or \\emph{Pharos} (Nicastro et al., in preparation) will be able to measure $\\Omega_{\\rm WHIM}$ to far greater precision with detections of numerous weaker X-ray forest lines, the \\citet[][hereafter N05]{nicastro05a} results present a key early confirmation of numerical predictions using observational capabilities that are \\emph{currently} within our grasp --- hence any test of their correctness is of great importance.  Although each of these two absorption systems was detected with high confidence through multiple redshifted X-ray absorption lines, the \\emph{individual} absorption lines were generally quite weakly detected, mostly at the $2-4\\sigma$ level.  Moreover, while they employed high--quality \\chandra\\ and \\fuse\\ data taken during exceptionally bright outbursts of Mrk 421, the many archived \\xmmnewton\\ observations of this source were not included in the analysis.  With roughly twice the effective area of \\chandra/LETG, \\xmm/RGS is in principle superior for X-ray grating spectroscopy between $\\sim 10-40$\\,\\AA; however, its slightly worse resolution (approximately 60\\,m\\AA\\ FWHM, versus 50\\,m\\AA\\ for \\chandra/LETG), higher susceptibility to background flares, and multitude of narrow instrumental features can present serious complications for WHIM searches.  Independent confirmation of the \\chandra\\ results with a separate instrument like \\xmm\\ is thus important.  While some groups have searched for WHIM features in a limited number of \\xmm\\ Mrk 421 spectra \\citep[e.g.,][]{ravasio05}, a complete and systematic analysis has yet to be performed.  Here we present a search for $z>0$ WHIM features employing all ``good'' observations of Mrk 421 available in the \\xmm\\ archive, and a comparison of these results to those presented by N05. ",
        "conclusions": "We have presented the highest signal--to--noise coadded \\xmm\\ grating spectrum of Mrk 421 to date, incorporating all available archival data.  This spectrum serves as an independent check on the recent detection of two $z>0$ WHIM filaments by N05.  While none of the \\chandra--detected absorption lines are seen in the \\xmm\\ spectrum, the upper limits derived from the \\xmm\\ data are consistent with the equivalent widths reported by N05 (even though the \\xmm\\ data contain a larger number of counts), and hence do not jeopardize the validity of the \\chandra\\ measurement.  The non--detections can be attributed primarily to narrow instrumental features in RGS1 and RGS2, as well as the inferior spectral resolution of \\xmm\\ and fixed--pattern noise at longer wavelengths. This underscores the extreme difficulty of detecting the WHIM, illustrates how the aforementioned (apparently small) effects can greatly affect the delicate measurement of weak absorption lines, and re--emphasizes the importance of high resolution and a smooth instrumental response function for current and future WHIM absorption line studies."
    },
    "0601/astro-ph0601599_arXiv.txt": {
        "abstract": "We report the discovery of Andromeda~X, a new dwarf spheroidal satellite of M31, based on stellar photometry from the Sloan Digital Sky Survey (SDSS).  Using follow-up imaging data we have estimated its distance and other physical properties. We find that Andromeda~X has a dereddened central surface brightness of $\\mu_{V,0} \\sim 26.7$ mag arcsec$^{-2}$ and a total apparent magnitude of $V_{\\rm tot} \\sim 16.1$, which at the derived distance modulus, $(m - M)_0 \\sim 24.12 - 24.34$, yields an absolute magnitude of $M_V \\sim -8.1 \\pm 0.5$; these values are quite comparable to those of Andromeda~IX, a previously-discovered low luminosity M31 satellite. The discoveries of Andromeda~IX and Andromeda~X suggest that such extremely faint satellites may be plentiful in the Local Group. ",
        "introduction": "\\label{txt:intro}  In hierarchical cold dark matter (CDM) models, large galaxies like the Milky Way and M31 form from the merger and accretion of smaller systems. Such models, while successful at large scales, predict at least 1 -- 2 orders of magnitude more low-mass dark subhalos at the present epoch than the observed abundance of dwarf galaxies \\citep[e.g.,][]{klyp99,moor99,bens02b}. This discrepancy, the ``missing satellite'' problem, is one of the most serious obstacles for matching CDM theory to observations.  A number of solutions have been proposed to address the problem, at least qualitatively. Star formation in low mass subsystems could be inhibited \\citep[e.g.,][]{some02,bens02a}. It is possible that the observed satellites are in fact embedded in much larger, more massive dark subhalos \\citep{stoe02}, or originally formed in more massive subhalos that were later tidally stripped \\citep{krav04}, and that we are therefore observationally sampling a less populous region of the initial satellite mass function.  All of these solutions aim to resolve the discrepancy between theory and observation by creating models which predict fewer directly observable satellites. A complementary observational approach would be to place more stringent constraints on the faint end of the galaxy luminosity function, but attempts to do so are hampered by the low surface brightnesses expected of galaxies in that regime. The advent of wide-field CCD surveys such as the Sloan Digital Sky Survey (SDSS) has made it possible to detect stellar structures with extremely low surface brightnesses \\citep[see, e.g.,][]{will02}, and in the past two years data from SDSS have been used to find two new Local Group members, Andromeda IX \\citep[\\aix;][]{zuck04b} and Ursa Major \\citep{will05}, satellites of M31 and the Milky Way, respectively; at the time of its discovery, each of these two objects was determined to be the least luminous, lowest surface brightness galaxy found up to that point.  Since the discovery of \\aix, we have been using SDSS photometry of M31 and its surroundings to select fields for deeper observations on other telescopes, in order to search for additional M31 companions. In this letter, we report the discovery, using SDSS and followup photometric data, of a new dwarf spheroidal companion to M31, one which appears comparable in luminosity to \\aix. Keck spectroscopy of stars in the field of this new satellite, and the kinematic and abundance information derived therefrom, are presented in a companion paper \\citep{guha05b}. For this work, we have assumed a distance to M31 of 783\\,kpc \\citep[$(m - M)_0 = 24.47$; e.g.,][]{holl98,stan98}. ",
        "conclusions": "As shown in Table \\ref{tbl:pars}, \\ax\\, is comparable in size, surface brightness, and apparent magnitude to \\aix, and, despite the uncertainty in its distance, \\ax\\, appears to be somewhat lower in luminosity ($M_V \\sim -8.1$). We conclude that \\ax\\, is a new extremely low luminosity dSph companion of M31, a result confirmed by kinematic data \\citep{guha05b}. The earlier discovery of \\aix\\, raised the question of whether such low luminosity galaxies were a rarity in the Local Group, or if \\aix\\, represented the tip of an iceberg, one of a large population of faint  dwarf satellites that have remained undetected owing to their extremely low surface brightnesses and low luminosities. \\ax's discovery, in conjunction with the recent discovery of the Ursa Major dwarf \\citep{will05}, strongly suggests that the latter scenario is closer to the truth -- that M31 and the Milky Way may in fact have a large number of low luminosity satellites.  Further analysis of SDSS data is therefore likely to yield more such dwarfs. To test this scenario, we are following up other stellar overdensities in the SDSS M31 scan with deeper observations on other telescopes.  Intriguingly,  in projection \\aix\\, and \\ax\\, lie relatively close to each other and to the major axis of M31, northeast of M31's center (Fig. \\ref{fig:m31sats}); as more new M31 satellites are found (as appears quite likely), we will be able to place constraints on the spatial distribution of M31 satellites and their (an)isotropy \\citep[e.g.,][]{koch05}, as well as on the satellite luminosity function. Thus the discovery and characterization of extremely low luminosity galaxies like \\aix\\, and \\ax\\, play a critical role in reconciling the predictions of current CDM galaxy formation models with the observable local Universe."
    },
    "0601/astro-ph0601550_arXiv.txt": {
        "abstract": "Recent optical and infrared studies have revealed that the heavily-reddened starburst cluster  Westerlund 1 (Wd 1) contains at least 22 Wolf-Rayet (WR) stars, comprising the richest WR population of any galactic cluster. We present results of a sensitive {\\em Chandra} X-ray observation of Wd 1 which  detected 12 of the 22 known WR stars and the mysterious emission-line star W9.  The fraction of detected WN stars is nearly identical to that of WC stars. The WN stars WR-A and WR-B as well as W9 are exceptionally luminous in X-rays and have similar hard heavily-absorbed X-ray spectra with strong Si XIII  and S XV  emission lines. The luminous high-temperature X-ray emission of these three stars is characteristic of colliding wind binary systems but their binary status remains to be determined. Spectral fits of the X-ray bright sources WR-A and W9 with isothermal plane-parallel shock models require high absorption column densities log N$_{\\rm H}$ = 22.56 (cm$^{-2}$)  and yield characteristic shock temperatures  kT$_{s}$ $\\approx$ 3 keV (T$_{s}$ $\\approx$ 35 MK). ",
        "introduction": "The heavily-reddened  cluster Westerlund 1 (Wd 1; Westerlund 1961, 1987) in Ara has recently been recognized as a rare example of a starburst cluster in our own Galaxy.  The cluster is massive, compact, and young with age estimates of $\\sim$3 - 5 Myr (Brandner et al. 2005 = B05; Clark et al. 2005 = C05). Wd 1  contains a remarkable collection of massive post-main sequence stars including early and late-type supergiants, a luminous blue variable candidate, and the largest known population of  Wolf-Rayet (WR) stars of any galactic cluster (C05; Negueruela \\& Clark 2005 = NC05; Clark \\& Negueruela 2002). A recent {\\em VLT} study has also revealed a faint population of  low-mass pre-main sequence stars and gives a photometric distance d = 4.0 $\\pm$ 0.3 kpc (B05), but spectroscopic studies allow a larger range of distances (C05). The extinction is A$_{\\rm V}$ $\\approx$ 9.5 - 13.6 mag (B05, C05).  Ongoing studies have so far identified 22 WR stars in Wd 1 and the census is likely incomplete. WR stars are highly-evolved  evolutionary  descendants of massive O-type stars that are in advanced nuclear burning stages and undergoing extreme mass-loss from high-velocity winds, rapidly approaching the end of their lives as supernovae. At least one supernova has already occurred in Wd 1 as evidenced by the  discovery of a new X-ray pulsar (Figure 1; Skinner et al. 2005a = S05a; Muno et al. 2006).   There is at present no  comprehensive theory of X-ray emission from WR stars. Previous observations have focused mainly on X-ray bright  WR $+$ OB binaries such as $\\gamma^2$ Vel and WR 140 (Skinner et al. 2001; Zhekov \\& Skinner 2000), whose hard emission  (kT $\\geq$ 2 keV) is thought to originate primarily in colliding wind shocks. Much less is known about the X-ray emission of single WR stars, but by analogy with O-type stars they are expected to emit soft X-rays (kT $<$ 1 keV) from instability-driven shocks formed in their supersonic winds (Gayley \\& Owocki 1995). Despite these expectations, X-ray emission from single WR stars has proven difficult to detect. Recent sensitive observations have shown that single carbon-rich WC stars are exceedingly faint in X-rays or perhaps even X-ray quiet,for reasons that are not yet fully understood (Skinner et al. 2005b = S05b). Additional X-ray observations are needed to quantify X-ray emission properties across the full range of WC and WN spectral subtypes.  The presence of a rich, equidistant, coeval population of WR stars in Wd 1 makes it an opportune target for X-ray observations, which are capable of penetrating the high extinction. We present the results of a sensitive {\\em Chandra} X-ray observation of Wd 1, focusing here on the WR population. This observation yields 12 new WR X-ray detections and provides valuable new information on the X-ray properties of this unique sample of galactic WR stars that can be used to test shock emission models and guide new theoretical development. ",
        "conclusions": "There are good reasons to believe that most of the WR X-ray detections in Wd 1 are binaries. This conclusion is more secure for WC stars than WN stars since there have been no previous X-ray detections of single WC stars, even at better sensitivities than obtained  here. Large differences in L$_{x}$ between WR stars of similar spectral type can be naturally explained if the luminous X-ray sources are colliding wind binaries. Furthermore, {\\em Chandra} preferentially detects harder X-ray sources in Wd 1 because of the high extinction. Plane-parallel shock models give good fits of the brightest X-ray detections and require shock temperatures kT $\\geq$ 2 keV. Such temperatures are higher than predicted for radiative shocks distributed in the winds of single stars, but are consistent with colliding wind emission in binary systems. Even so, more definitive proof of binarity is needed from optical/IR follow-up work. And, interesting questions remain in the X-ray regime. What is the origin of the excess absorption that is inferred from X-ray spectral fits of WR-A  and W9? How does the WR X-ray luminosity function behave at low L$_{x}$: are the undetected WR stars faint sources below our detection limit or are they  X-ray quiet?"
    },
    "0601/astro-ph0601285_arXiv.txt": {
        "abstract": "{We present here new images of relics and halo sources in rich cluster of galaxies and the correlation between the halo radio surface brightness versus the cluster bolometric X-ray luminosity. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/hep-th0601217_arXiv.txt": {
        "abstract": "We consider a Friedmann brane moving in a bulk impregnated by radiation. The setup is strongly asymmetric, with only one black hole in the bulk. The radiation emitted by this bulk black hole can be reflected, absorbed or transmitted through the brane. Radiation pressure accelerates the brane, behaving as dark energy. Absorption however generates a competing effect: the brane becomes heavier and gravitational attraction increases. We analyse the model numerically, assuming a total absorbtion on the brane for $k=1$. We conclude that due to the two competing effects, in this asymmetric scenario the Hawking radiation from the bulk black hole is not able to change the recollapsing fate of this brane-world universe. We show that for \\textit{light} branes and early times the radiation pressure is the dominant effect. In contrast, for \\textit{heavy} branes the self-gravity of the absorbed radiation is a much stronger effect. We find the \\textit{critical} value of the initial energy density for which these two effects roughly cancel each other. ",
        "introduction": "According to brane-world models, our universe is a hypersurface (the brane) embedded in a $5$-dimensional spacetime (the bulk), in which Einstein's gravity holds. The attempts to increase the number of dimensions in physics have a long history, originating with the pioneering works of Kaluza and Klein. In the more recent models presented in \\cite{ADD}, the (flat) extra dimensions are still compactified. The novelty of the brane-world models generalizing the original proposal \\cite{RS2} of Randall and Sundrum (RS) consist in allowing for one additional \\textit{noncompact} dimension. Instead of compactifying, the bulk is warped.  Gravitational dynamics on the brane is given by a modified Einstein equation, which arises from the projection of the 5-dimensional dynamics and the junction conditions across the brane. A major difficulty in this theory is that due to a new source term, arising from the projection of the bulk Weyl tensor to the brane, the system of differential equations describing gravitational evolution is not closed on the brane. This new gravitational theory allows general relativity as the low energy density limit (as compared to the brane tension $\\lambda $).  Brane-world models were initially $Z_{2}$-symmetric, with identical copies of the bulk on the two sides of the brane. Equivalently, the brane could be imagined as a domain boundary. This assumption, although it considerably simplifies the model, seems unnecessarily restrictive in the context of brane-world models with curved branes (generalized RS type 2 models). In this context, our observable universe may be imagined as a Friedmann brane moving either in a Schwarzschild-anti de Sitter bulk (if the bulk is vacuum), or in a 5-dimensional Vaidya-anti de Sitter (VAdS5) spacetime (if the bulk contains radiation). The modified Einstein equation in the $Z_{2}$% -symmetric case was given covariantly in \\cite{SMS}, and generalized for asymmetric embedding in \\cite{Decomp}. The effects of asymmetric embedding were also analyzed in a more generic class of models, containing induced gravity contributions \\cite{Induced}. There it was shown that the late time universe behaves differently under the introduction of a small asymmetry of the embedding: on the two branches of such theories. While on the RS branch asymmetry produces late-time acceleration, on the DGP branch the late-time acceleration is diminished by asymmetry.  Whenever there is radiation in the bulk, it can arise from both brane and bulk sources. The situation when the brane radiates into the bulk was already considered both in a symmetric \\cite{LSR} and in an asymmetric scenario \\cite{RadBrane}, \\cite{VernonJennings} in the framework of generalized RS theories. In both scenarios a radial emission of radiation into the bulk was considered. A more realistic set-up, allowing for the emitted gravitons to follow geodesic paths, was advanced recently in \\cite% {Langlois}.  In this paper we discuss the other possibility, when the radiation is emitted by sources in the bulk. For this, we embed the Friedmann brane asymmetrically into VAdS5 spacetime. The symmetry of the embedding is severely broken by allowing for a radiating black hole\\ on one side and no black hole on the other side of the brane. Due to this asymmetry, there is a radiation pressure acting on one side of the brane. This radiation pressure continuously accelerates the motion of the brane, leading to accelerated cosmic expansion. In principle, this mechanism can be a \\textit{dark energy candidate}.  In brane-world models, standard matter is confined to the brane. Therefore the radiation coming from the bulk should be of non-standard nature. We discuss here the case when this is the Hawking radiation of the single bulk black hole. The expression of the Hawking radiation in a 5-dimensional Schwarzschild-anti de Sitter (SAdS5) spacetime was derived for the curvature index $k=1$ in \\cite{EHM}-\\cite{GCL}. This energy density was employed, but for $k=0$ in \\cite{Jennings} for the study of asymmetrically embedded Friedmann branes into VAdS5.  Some part of the radiation reaching the brane will be absorbed, some part will be reflected and the remaining part will go through. We will study here only the models without reflection in order to have the VAdS5 spacetime in both bulk regions (Fig \\ref{Fig:1}). We disregard the possibility of reflection, because there is no known exact solution with cosmological constant describing a crossflow of radiation streams. Such a spacetime with both incoming and outgoing null dust streams, but without a cosmological constant is known in $4$-dimensions \\cite{GergelyND}, however no generalization including a cosmological constant has been found. Thus we discuss the case when the reflected component can be neglected.  We keep however the absorption, as an essential element of our model, in contrast with the one presented in \\cite{Jennings}, where there is a black hole on each side of the brane, but their radiation is \\textit{completely transmitted} through the brane.  By retaining the absorption on the brane, we obtain novel features. The absorbed radiation appears on the brane as continuously emerging energy. Thus the energy density on the brane increases, strengthening the gravitational self-attraction of our universe. As consequence of absorption, cosmic expansion slows down.  Therefore there are two competing effects in our model: radiation pressure accelerates, while the absorbed radiation decelerates cosmic expansion. We would like to study how their balance affects the Friedmann brane-world. We derive the relevant equations of the model in Section II and we comment on the physical interpretation in Section III.  As our model emphasizes the effects of the absorption (as opposed to \\cite% {Jennings}), we finally choose a \\textit{total absorption} on the brane in Section IV. We also employ $k=1$, as the formula for the energy density of the Hawking radiation was derived for this case \\cite{EHM}-\\cite{GCL}. Then we introduce new dimensionless variables, adapted to our choice of $k=1$.  We provide quantitative results by numerical analysis in Section V. First we discuss cosmological evolution in the case of a non-radiating bulk black hole. Then, by switching on the radiation, we see how the two competing effects modulate the cosmological evolution. We show that a critical behaviour can appear, when these two competing effects roughly cancel each other. We discuss the relevance of these results in a broader context in Section VI.  Throughout the paper a tilde distinguishes the quantities defined on the 5-dimensional spacetime and a hat denotes dimensionless quantities and units with $c=1=\\hbar $ are chosen. ",
        "conclusions": "In this paper we have considered a scenario which is able to deal with the possibility of a Hawking radiation escaping a bulk black hole and we have investigated the consequences of such a radiation on the cosmological evolution of a closed $k=1$ universe.  For this purpose first we have developed a formalism suitable for describing the asymmetric brane-world containing a single radiating bulk black hole on the left bulk region. In the most generic scenario the radiation can be partially absorbed, partially transmitted and partially reflected on the brane. We have considered here the case of negligible reflection, because we wanted to dispose of an exact solution of the Einstein equations in the bulk. By suppressing the reflection, and considering the geometrical optics limit of the radiation, the bulk can be described by the VAdS5 metric.  The part of radiation absorbed on the brane generates quite interesting effects, which were not studied before. These can be related to both dark energy and dark mattter. In order to concentrate on these effects, the absorption was chosen as maximal. By suppressing transmission, the equations could be partially integrated, cf. Eq. (\\ref{mint}). Then we have introduced dimensionless quantities in order to perform a complex numerical analysis.  We have carefully compared the cosmological evolutions of the models both without and with radiation and found that the Hawking radiation of the bulk black hole represents merely a perturbation of the non-radiating model. This holds true during the entire cosmological evolution.  We have studied these perturbations in detail and shown that in the presence of the radiation, there are two competing effects. The absorption on the brane contributes towards the self-gravity of the brane, which then tends to recollapse faster. This phenomenon appears as continuously emerging dark matter on the brane. By contrast, radiation pressure tends to accelerate away the brane from the bulk black hole, producing an equivalent of dark energy. For the numerical data employed, we have found the critical value of the initial brane energy density for which these two competing effects roughly annihilate each other. We have also analysed cosmological evolution for both lighter and heavier branes and found that (at early times) the radiation pressure is dominant for light branes, while for heavy branes the self-gravity dominates during the whole cosmological evolution.  As a main result, we have proved that the asymmetry introduced in the model is not able to change the recollapse of the universe in the $k=1$ case, regardless of whether the radiation is switched off or on. Our numerical analysis has shown that the bulk Hawking radiation cannot change the final fate of the recollapsing universe either."
    },
    "0601/astro-ph0601146_arXiv.txt": {
        "abstract": "We estimate the luminosity evolution and formation rate for over 900 GRBs by using redshift and luminosity data calculated by Band, Norris, $\\&$ Bonnell (2004) via the lag-luminosity correlation. By applying maximum likelihood techniques, we are able to infer the true distribution of the parent GRB population's luminosity function and density distributions in a way that accounts for detector selection effects.  We find that after accounting for data truncation, there still exists a significant correlation between the average luminosity and redshift, indicating that distant GRBs are on average more luminous than nearby counterparts.  This is consistent with previous studies showing strong source evolution and also recent observations of under luminous nearby GRBs.  We find no evidence for beaming angle evolution in the current sample of GRBs with known redshift, suggesting that this increase in luminosity can not be due to an evolution of the collimation of gamma-ray emission. The resulting luminosity function is well fit with a single power law of index $L'^{-1.5}$, which is intermediate between the values predicted by the power-law and Gaussian structured jet models. We also find that the GRB comoving rate density rises steeply with a broad peak between $1<z<2$ followed by a steady decline above $z> 3$.  This rate density qualitatively matches the current estimates of the cosmic star formation rate, favoring a short lived massive star progenitor model, or a binary model with a short delay between the formation of the compact object and the eventual merger. ",
        "introduction": "\\label{sec:Introduction_ch3}  There are currently roughly three dozen gamma-ray burst events (GRBs) for which we have independently measured redshifts.  Most of these redshift determinations come from either identification of absorption lines in the afterglow spectra, attributed to the gas in the host galaxy, or from observations of emission lines from the host galaxy. The combination of these techniques has resulted in a small but growing GRB sample with redshifts ranging from 0.0085 to 4.5 and a distribution peaking around $z \\sim 1$.  From this small sample, it is already abundantly clear that the isotropic equivalent energy $E_{iso}$ released in the prompt GRB phase is not a standard candle. The total radiated energy taken at face value (i.e. when not correcting for a beaming factor $d\\Omega$) clearly spans several orders of magnitude, ranging from $10^{47}$ for the closest event, GRB 980425 at $z=0.0085$ \\citep{Kulkarni98}, to $10^{54}$ for GRB 990123 at $z=1.6004$ (Kulkarni et al. 1999). Recently \\citet{Sazonov04} and \\citet{Soderberg04} have reported on gamma-ray observations of a nearby underluminous GRB occurring at redshift of $z=0.106$. These new findings have added to the speculation that there is either a substantial under luminous population of GRBs which cannot be seen at large distances and/or that nearby events ($z <$ 0.15) are underluminous compared to distant counterparts, pointing to the evolution of the average energy emitted by a GRB with time.  A measure to the extend to which luminosity evolution exists in the GRB population, along with their true luminosity function and density distribution, may yield important clues regarding the nature of gamma-ray bursts and how they're progenitors have evolved with time.  Although the physics of the underlying GRB engine is hidden from direct observation and is yet uncertain, the total GRB energy budget is most likely linked to the mass and/or rotational energy of the GRB progenitor. Understanding of how this energy budget has changed with time may offer constraints on progenitor properties and may ultimately point to the physics leading to their explosions. Since GRB progenitors are most likely linked to compact objects (supermassive rotating star, black hole or neutron star mergers) understanding how the GRB luminosity function evolves with time may give insight to the host environment in the early universe, namely the star formation rate or initial stellar mass functions at high redshifts.  Any attempt at quantifying the evolution of intrinsic source properties of parent populations must account for Malquist type biases.  Detection thresholds prevent events below a certain flux from being observed, resulting in the detection of only bright objects at large distances.  Combined with the fact that bright events are typically rare, it is very easy in astronomy to incorrectly conclude that the distant universe is filled with extremely bright rare objects.  Any attempt at measuring the correlation between luminosity and redshift without properly accounting for selection effects will grossly overestimate the correlation strength between the two variables. Flux limited samples are a classic problem in astronomy, which manifested prominently in early quasar studies. Fortunately, straight forward methods have been devised to account for such effects based on maximum likelihood techniques. These methods allow for the correct estimation of the correlation strength between a truncated data set as well as an estimate on the underlying parent population. The \"catch\" of such techniques is that the overall normalization of the resulting parent distributions cannot be determined, although their functional forms are constructed in such a way to account for the data truncations.  These techniques also have to limitation of requiring a large sample sizes and more importantly, an extremely good understanding of the survey's detection thresholds (i.e. the flux cutoff for magnitude limited samples). The use of the current sample of GRBs with known redshift is limited by both of these restrictions.  The current size of a little over two dozen bursts does not lend itself well to producing statistically robust results, especially in the high and low redshift regimes for which only a handful of events have been detected.  Furthermore, the sample is an accumulation of observations from several different spacecraft, all of varying detector thresholds.  It would seem that these limitations could only be overcome by the accumulation of a larger data set with consistent detector thresholds which is expected to come from the Swift spacecraft and the upcoming GLAST mission.  Fortunately, several authors have announced empirical Cepheid like correlations linking intrinsic burst properties, such as luminosity \\citep{norris00} and the total radiated energy \\citep{amati02,ghirlanda04a} to other GRB observable.  These correlations may allow for the determination of burst redshifts directly from the gamma-ray data, which has the advantage of being relatively insensitive to extinction and observable at far greater distances than afterglow line measurements. The first of these correlations was reported by \\citet{norris00}. Using 6 BATSE detected bursts with known redshift, they found an anti-correlation between the {\\it source} frame lag between the 25-50 keV and 100-300 keV emission and the absolute luminosity of the GRB.  More recently, \\citep{ghirlanda04a} reported an empirical correlation between the collimation correction total energy $E_{\\gamma}$ radiated by the burst and the rest frame energy at which most of the prompt radiation is emitted $E_{pk}$.  Using these relationships, it is now possible to estimate \"pseudo\" redshifts for a much larger number of GRBs detected by the BATSE instrument which perviously lacked any information as to their distance.  More importantly, the BATSE detector threshold is relatively well understood for the entire sample, making the resulting pseudo redshift data excellent for statistical analysis.  In this paper, we examine the issue of luminosity and density evolution by using a sample of over 900 BATSE GRBs for which the luminosity and redshift where recently estimated by \\citet{band04} through the use of the lag-luminosity correlation.  We limit our analysis to the lag-luminosity correlation primarily due to the lack of jet opening angle $\\theta_{j}$ information that is required for the use of the $E_{\\gamma}$-$E_{pk}$ relation.  This relationship requires knowledge of $\\theta_{j}$ in order to determine the collimation factor, which is only known for bursts with measured jet break times and hence cannot be used with the BATSE sample in consideration for this paper.  We found that the more general, and much broader, correlation between intrinsic $E_{pk}$ and $E_{iso}$ reported by \\citet{amati02} did not provide redshift constraints for a majority of the bursts in our sample.  This is consistent with recent observations by \\citet{nakar} and \\citet{bandpreece} who also found large fractions of their BATSE samples to be inconsistent with the Amati correlation.  Therefore we limit the current analysis to distances estimated through the use of the lag-luminosity relationship.  To our sample, we apply statistical techniques developed by \\citet{Lynden-Bell71} and \\citet{efron92} and first applied to GRB analysis by \\citet{lloyd02} to measure the underlying luminosity and density distribution in a way that properly accounts for the detection thresholds of the BATSE instrument.  We find a strong (11.63 $\\sigma$) correlation between luminosity and redshift that can be parameterized as $L(z) = (1+z)^{1.7 \\pm 0.3}$.  The resulting cumulative luminosity function $N(L')$ is well fit by double power law separated by a break energy of about $10^{52}$ ergs s$^{-1}$, with the differential luminosity function $dN/dL'$ exhibiting a power law shape of $L^{-1.5}$ below this luminosity. We show that the GRB comoving rate density increases roughly as $\\rho_{\\gamma}(z) \\propto (1+z)^{2.5}$ to a redshift of $z \\approx 1$ followed by a flattening and eventual decline above $z>3$.  This rate density is in qualitatively agreement with recent photometric estimates of the cosmic star formation rate (SFR), as would be expected from massive short lived progenitors.  In $\\S 2$, we describe the data set that we use in our study.  In $\\S 3$ we discuss the statistical methods applied to this data to estimate the GRB luminosity function and comoving rate density as well as to test for any correlation between luminosity and redshift. In $\\S 4$ we present the resulting demographic distribution functions of this analysis followed in $\\S 5$ by a discussion of the implications of the shape and evolution of the luminosity function and comoving rate density on various jet profile.  We show that there is no evidence for beaming angle evolution in the current sample of GRBs with known redshift, suggesting that the variation of the observed luminosity with redshift can not be due to an evolution of the collimation of gamma-ray emission.  We conclude by examining how the similarity between the SFR and the GRB comoving rate density tentatively favors short lived progenitor models. ",
        "conclusions": "\\label{conclusion:ch3}  In this work we perform demographic studies on a large sample of luminosity and redshift data found through the use of the lag-luminosity correlation.  By applying maximum likelihood techniques, we are able to obtain an estimate of the luminosity evolution, luminosity function, and density distributions in a way that accounts for detector selection effects.  We find that there exists a strong (11.63 $\\sigma$) correlation between luminosity and redshift that can be parameterized as $L(z) = (1+z)^{2.58}$.  The resulting cumulative $\\Phi(L')$ and differential $d\\Phi/dL'$ luminosity functions are well fit by double power laws separated by a break energy of about $10^{52}$ ergs s$^{-1}$, with $d\\Phi/dL'$ exhibiting a power law shape of $L^{-1.5}$ below this luminosity. This value does not immediately discriminate against any proposed structured jet models, but it may indicate that a more complicated jet profile is need to explain the observed luminosity function of GRBs.  The GRB comoving rate density is found to increase as $\\rho_{\\gamma}(z) \\propto (1+z)^{2.5}$ to a redshift of $z \\approx 1$ followed by flattening and eventual decline above $z>6$. Although, the conversion between $\\rho_{\\gamma}(z)$ and an estimate of the SFR cannot be quantitatively made due to the uncertainty regarding the GRB progenitors and their initial stellar mass functions, it can be said that $\\rho_{\\gamma}(z)$ does qualitatively follow recent photometric estimates of the SFR, as would be expected from massive short lived progenitors.  We stress that these conclusions are based on the validity of the lag-luminosity correlation, which still stands to be confirmed as new redshift data becomes available.  A full confirmation, and most probably further calibration, of this distance indicator will have to wait for addition detections with the Swift spacecraft, which should have the collecting area necessary to obtain high signal-to-noise energy dependant light curves from which to measure statistically significant lags. Even with the slew of directly measured luminosity and redshift data expected to come from the Swift mission, empirical distance indicators may still play an important roll in expanding the available GRB data set through the use of archival BATSE and BepooSAX data for statistical studies such as the analysis performed in this work."
    },
    "0601/astro-ph0601370_arXiv.txt": {
        "abstract": "{} {To predict time delays for a sample of gravitationally lensed quasars and to evaluate the accuracy that can be realistically achieved on the value of $H_0$.} {We consider 14 lensed quasars that are candidates for time-delay monitoring and model them in detail using pixelized lens models. For each system, we provide a mass map, arrival-time surface and the distribution of predicted time-delays in a concordance cosmology, assuming $H_0^{-1}=14$~Gyr ($H_0=70$ in local units). Based on the predicted time-delays and on the observational circumstances, we rate each lens as `excellent' or `good' or `unpromising' for time-delay monitoring.  Finally, we analyze simulated time delays for the 11 lens rated excellent or good, and show that $H_0$ can be recovered to a precision of 5\\%.} {In combination with COSMOGRAIL paper~I on the temporal sampling of lensed quasar light curves, the present work will help design monitoring campaigns of lensed quasars.} {} ",
        "introduction": "Gravitational lensing of distant quasars is one of many possible routes to $H_0$. It has unique advantages. First, lensing depends on well-understood physics: gravitation.  Second, time-delay observations require modest resources, hence low-demand telescopes can make a significant contribution.  Third, the galaxy models used to convert time-delays into $H_0$ have made considerable progress in the past decade.  As a result, time-delay measurements have become an increasingly active research topic.\\footnote{According to ADS, the original paper by \\cite{refsdal64} pointing out the connection between gravitational lensing time-delays and $H_0$ was cited on-average once every two years through the 1960s and 70s, whereas nowadays it is cited about once every two weeks.}  So far there are 12 secure time-delays (Table~\\ref{tablobs}), of which 10 yield an estimate of the Hubble constant --- the lens identification in PKS~1830--211 remains controversial (\\cite{cou02}, \\cite{winn02}), and the lensing galaxy HE~0435--122 may be anomalous (Kochanek 2005).  The dominant uncertainty in measuring $H_0$ from lensing is the non-uniqueness of lens mass profiles that can reproduce the observables.  Before the non-uniqueness of mass-models was widely appreciated, researchers would usually fit a single family of mass models to data, leading to over-optimistic error bars.  Experimenting with different kinds of mass model for the same data pointed to much larger uncertainties (Schechter et al.\\ 1997, \\cite{sw97,bernstein99,keeton00}).  More recently, procedures involving sampling an ensemble of models according to some prior are being preferred, in order to derive a more useful picture of the uncertainties (\\cite{ws00,sw04,oguri04a}, Jakobsson et al.\\ 2005).  A fair summary of current $H_0$ results from lensing is that the error-bars are competitive on the young-Universe (i.e., high $H_0$) side but the old-Universe side needs improvement.  Clearly, to reach the 5--10\\% accuracy claimed by some other techniques (e.g., \\cite{freedman01,spergel03}), more time-delays are needed.  But to run monitoring campaigns efficiently, it is important to have preliminary estimates for time-delays --- witness the tenfold range in the known values in Table~\\ref{tablobs} --- as well as to identify the most promising systems to monitor.  This paper supplies such information.  For a sample of 14 lenses we provide predicted time delays with uncertainties, a rating of prospects as `excellent', `good', or `unpromising' based on both models and the observational situation, and finally an estimate of the precision on $H_0$ obtainable from these lenses.  A companion paper by Eigenbrod et al.\\ (2005; COSMOGRAIL I) is devoted to the determination of the optimal strategy to use in order to measure time-delays (temporal sampling of the light curves, object visibility and variability, contamination by microlensing, etc).  Together, these papers help design an observational campaign.  An ideal time-delay lensing system has the following features: {\\bf 1-} bright optical images, {\\bf 2-} large angular image separations ($>$1\\arcsec), {\\bf 3-} light path unperturbed by nearby structure, {\\bf 4-} known or easy-to-measure lens redshift $\\zl$.  Our sample of 14 has been selected using these criteria as a guideline, though not a strict requirement. We have considered only objects for which the time-delay can be measured in the optical.  The main results of this paper are in Sect.~\\ref{indivsec} and \\ref{predacc}, which present ensembles of models for the 14 individual systems and then estimate the precision to which $H_0$ could be recovered from them. But before going into details of the models, it is useful to preview the results and compare them with measured systems.  We do this in Sect.~\\ref{compsec}. ",
        "conclusions": "In this paper we do three things: first, we introduce a simple rough predictor for time delays $\\Tastrom$, second, we make model predictions for 23 time delays covering 14 lenses, and finally we estimate the precision in the Hubble time inferred from simulated data on the 11 best lenses.  The main conclusion is that no single lens can usefully constrain $H_0$, but time delays accurate to $\\simeq1$d on $>10$ lenses can yield $H_0$ accurate to 5\\%.  In the histograms in Figs.~\\ref{plotJ1155}--\\ref{plotJ0924}, typically 90\\% of the area ranges over a factor of two in time delays.  Hence, a monitoring program can have 90\\% confidence in succeeding --- provided the quasar is sufficiently variable --- if the sampling allows for the appropriate 90\\%-range of possible time delays.  There is no single characteristic shape for the histograms, but the pattern of a low-end tail and a high-end cliff is common.  The large uncertainty in the predicted time delays reflects the large variety of mass models models that can reproduce the observed image positions in any given lens. The prior we have for deciding what is an allowable mass model for a galaxy is very conservative, so our models have more variety than reality.  But not very much more --- if all real galaxy lenses belonged to some known parametrization, then error-bars on $H_0$ from fitting such parametric models to observed time-delay systems would overlap, but as is evident from Fig.~12 in \\cite{cou03}, those error-bars do not overlap.  Actually, the basic results about predicted time delays are already present in the summaries Figs.~\\ref{plotobs} and \\ref{plocand}.  To interpret these figures, recall that $\\Tastrom$ makes a preliminary prediction for the time-delay, while the deviation from the oblique line depends on the details of mass distribution and lens morphology. From Fig.~\\ref{plotobs} we see that $\\Tastrom$ gets us to within a factor of two of the observed values.  Now, the error bars ---note that the error bars in tables and figures are 68\\% confidence--- in Fig.~\\ref{plocand} are generally less than a factor of two; thus detailed modelling does provide a better prediction than $\\Tastrom$ alone, but not dramatically better.  We can also see that several of the error bars in Fig.~\\ref{plocand} are shorter on top, thus indicating a low-end tail and a high-end cliff.  The image morphology of a lens is correlated with the uncertainty in the time delays, especially in quadruples.  Core quadruples generically have short time delays and are unlikely to be useful for time delays; the case of B0435-122 is illustrative.  Long- and short-axis quadruples, having three images close together, are likely to have only one measurable time delay.  Inclined quadruples are the most promising, since they usually have two time delays in the measurable range, and sometimes three.  Among doubles, inclined systems tend to be somewhat better constrained than axial systems. Significant asymmetry in the lens is a disadvantage, but in compensation, asymmetry increases the chance of having three measurable time delays.  Thus B1608+656 is an asymmetric inclined quadruple with three measured time delays; J2033--472 may prove to be another, and is among our excellent candidates.  Surprisingly, a large external shear appears to reduce uncertainties in the time delay. This is particularly noticeable in the inclineddoubles J1650+425, B0909+532, and J0903+502, and the short-axis quads B1422+231 and J1131--123. The reason is not clear; it may be that since external shear reduces amount of mass needed in the main lens to produce multiple images, it reduces the available model-space.  Finally, the results from combining several lenses are very encouraging. Assuming time delays accurate to 1\\thinspace d we find that the model-dependent uncertainty in $H_0^{-1}$ reduces to less than 10\\% on combining the 5 best lenses, and about 5\\% on combining the best 11 lenses.  The uncertainties are asymmetric, with the lower limit on the Hubble time being tighter than the upper limit.  More work needs to be done on the model prior before we can truly attain 5\\% accuracy, but meanwhile our results help provide both motivation and observing strategies for accurate time-delay measurements."
    },
    "0601/astro-ph0601693_arXiv.txt": {
        "abstract": "Over the last few years our understanding of local Type Ibc supernovae and their connection to long-duration gamma-ray bursts has been revolutionized.  Recent discoveries have shown that the emerging picture for core-collapse explosions is one of diversity.  Compiling data from our dedicated radio survey of SNe Ibc and our comprehensive HST survey of GRB-SNe together with ground-based follow-up campaigns, I review our current understanding of the GRB-SN connection.  In particular, I compare local SNe Ibc with GRB-SNe based on the following criteria: (1) the distribution of optical peak magnitudes which serve as a proxy for the mass of $^{56}$Ni produced in the explosion, (2) radio luminosity at early time (few days to weeks) which provides a measure of the energy coupled to on-axis relativistic ejecta, and (3) radio luminosity at late time (several years) which constrains the emission from GRB jets initially directed away from our line-of-sight.  By focusing on these three points, I will describe the complex picture of stellar death that is emerging. ",
        "introduction": "Twenty years have passed since the class of Type Ibc supernovae (SNe Ibc) was initially recognized as a distinct population of core-collapse explosions \\cite{emn+85,wl85,f97}.  Their lack of homogeneity and low event rate, ($\\sim 10\\%$ of locally discovered SNe), did not motivate focused observational programs.  In 1998, however, SNe Ibc enjoyed an explosion of new-found interest thanks to the discovery of Type Ic SN\\,1998bw ($d\\approx 36$ Mpc) within the {\\it BeppoSAX} localization error box of gamma-ray burst, GRB\\,980425 \\cite{gvv+98,paa+00}.  While the $\\gamma-$ray energy release, $E_{\\gamma}$, of GRB\\,980425 was a factor of $10^{4}$ below that of GRBs, SN\\,1998bw was (and still remains) the most luminous radio SN ever observed \\cite{kfw+98}.  Two unusual features were noted from the radio data: significant energy ($E_{\\rm radio} \\sim 10^{49}\\,$erg) coupled to mildly relativistic (Lorentz factor, $\\Gamma\\sim 3$) ejecta and evidence for episodic energy injection \\cite{kfw+98,lc99}.  Moreover, the bright optical emission required production of $\\sim 0.5\\,M_\\odot$ $^{56}$Ni, comparable to that inferred for Type Ia supernovae while the broad absorption lines (indicative of photospheric velocities above 30,000 km~s$^{-1}$) implied a total kinetic energy of $3\\times 10^{52}$ erg \\cite{i+98,w+99}.  These observations have been interpreted under the framework of the ``collapsar model'' (e.g. \\cite{m+01}) in which a central engine (accreting black hole) plays a significant role in exploding the star. ",
        "conclusions": "We compare the optical and radio properties for local SNe Ibc and GRB-SNe.  We show that the optical luminosities for GRB-SNe and local SNe Ibc are comparable and therefore cannot be used to distinguish between the two samples.  From our comprehensive radio survey of local SNe, however, we are able to show a clear distinction between GRB-SNe and SNe Ibc: the radio luminosities of local events are typically $10^4$ times fainter than those of typical GRB-SNe and $10^2$ times fainter than SN\\,1998bw.  We conclude that GRB-SN explosions couple the bulk of their energy to highly relativistic ejecta while local SNe couple a relatively tiny fraction.  Finally, we use our late-time radio observations to constrain the fraction of SNe Ibc harboring off-axis GRBs to less than 10\\%.  This holds in particular for the local broad-lined SNe Ibc which have been argued (based on their spectral similarity to GRB-SNe) to be associated with off-axis GRBs.  In conclusion we find that while most GRB explosions have a supernova component, only a small fraction of SNe Ibc are capable of producing the copious relativistic ejecta characteristic of gamma-ray bursts."
    },
    "0601/astro-ph0601687_arXiv.txt": {
        "abstract": "Rate of period change $\\dot{P}$ for a Cepheid is shown to be a parameter that is capable of indicating the instability strip crossing mode for individual objects, and, in conjunction with light amplitude, likely location within the instability strip. Observed rates of period change in over 200 Milky Way Cepheids are demonstrated to be in general agreement with predictions from stellar evolutionary models, although the sample also displays features that are inconsistent with some published models and indicative of the importance of additional factors not fully incorporated in models to date. ",
        "introduction": "Cepheids represent a brief phase in the post-main-sequence evolution of stars that originally had masses in excess of $\\sim 3 \\frac{1}{2} M_{\\sun}$ \\citep{tu96}. As evolved objects they populate the Cepheid instability strip in the HR diagram according to the manner in which they generate energy, depending upon strip crossing. Intermediate mass stars can become unstable to radial pulsation during shell hydrogen burning (first crossing), twice during core helium burning (second and third crossings), and according to some evolutionary models \\citep{ib65,be77,be85,xl04}, twice during shell helium burning (fourth and fifth crossings). A common feature of many recent evolutionary models for intermediate-mass stars, which incorporate the new opacity tables \\citep[e.g.,][]{mm00,mm02,bo00,sa00}, is that they permit only three strip crossings for such stars, since core oxygen ignition occurs prior to a separate shell helium-burning phase.  In all cases the evolution of stars through the Cepheid instability strip should be associated with gradual changes in overall dimensions, and hence periods of pulsation: period increases for evolution towards the cool edge of the instability strip as the stellar radius grows, and period decreases for evolution towards the hot edge as the stellar radius decreases. The observed parabolic trends in Cepheid O--C diagrams (plots of the differences between Observed times of light maximum and those Computed from a linear ephemeris) have been recognized for the past half century as evidence for the evolution of such stars through the instability strip \\citep{pa58,st59,ei82}. As noted by \\citet{st59}, ``It appears that studies of period change are by far the most sensitive test available to the astronomer for detecting minute alterations in the physical characteristics of a star.''  Observations of period changes in Cepheids have been matched with some confidence to evolutionary models of massive stars in various crossings of the instability strip \\citep[e.g.,][]{tu98,tb01,tb04} in order to identify the direction of strip crossing for individual variables. When used for such purposes, the study of Cepheid period changes becomes an important tool for the characterization of individual members of the class.  In principle it should also be possible to use rate of period change for individual Cepheids to establish likely location within the instability strip. Because strip crossings for individual Cepheids occur at different rates and at different luminosities for specific stellar masses, the observed rates of period change must be closely related to strip crossing mode and location within the instability strip. Potential constraints are imposed by variations in chemical composition and pulsation mode, e.g., fundamental mode, first overtone, etc. \\citep{be97,te99}, as well as by our limited ability to establish small rates of period change for O--C data containing sizeable observational uncertainties \\citep{sz83}. In this paper we demonstrate the link in more detail. ",
        "conclusions": "Our intent here is to demonstrate that rate of period change for a Cepheid is a useful parameter that permits one to characterize the variable in terms of specific evolutionary state. Information on $\\dot{P}$ for a Cepheid, in conjunction with its known pulsation period and light amplitude, can be used to identify the strip crossing mode for the object as well as its likely location within the strip, the latter independent of its observed color and reddening. The parameter $\\dot{P}$ may even be useful for establishing if a Cepheid is a fundamental mode pulsator or an overtone pulsator, although we leave that as a future exercise.  If we interpret the results of Figs. 5--7 as a generic indicator of how pulsation amplitude varies across the instability strip, then the width of the strip in $\\dot{P}$ at constant period, which amounts to $\\sim 1.2$ in log $\\dot{P}$, must encompass a range of $\\sim 16$ in $\\dot{P}$. Of that, an intrinsic dispersion in log $\\dot{P}$ values amounting to perhaps 0.4--0.5, a factor of $\\sim 3$, presumably arises from actual internal differences in the Cepheids resulting from different histories for their progenitor stars. Specific stellar evolutionary models presented in Fig. 2 predict a smaller variation in log $\\dot{P}$ than what is observed, which may reflect the simplicity of the models. In that regard, observed rates of period change in Cepheids can play an important role as a check on how closely stellar evolutionary models match real stars. Until now Cepheid period changes have not been used for that purpose."
    },
    "0601/astro-ph0601222_arXiv.txt": {
        "abstract": "The tidal  torque exerted  by a protoplanetary  disk with  power law surface   density  and   temperature  profiles   onto   an  embedded protoplanetary embryo is generally a negative quantity that leads to the  embryo inwards  migration. Here  we investigate  how  the tidal torque balance  is affected at  a disk surface density  radial jump. The jump has two consequences : \\begin{itemize} \\item it affects the  differential Lindblad torque. In particular if the  disk  is merely  empty  on  the  inner side,  the  differential Lindblad torque almost amounts  to the large negative outer Lindblad torque. \\item It  affects the  corotation torque, which  is a  quantity very sensitive to  the local  gradient of the  disk surface  density.  In particular if  the disk is  depleted on the  inside and if  the jump occurs radially  over a  few pressure scale-heights,  the corotation torque  is  a  positive quantity  that  is  much  larger than  in  a power-law disk. \\end{itemize} We show  by means  of customized numerical  simulations of  low mass planets embedded  in protoplanetary  nebulae with a  surface density jump that the second effect is dominant, that is that the corotation torque  largely dominates  the differential  Lindblad torque  on the edge  of a central  depletion, even  a shallow  one. Namely,  a disk surface density  jump of about  $50$~\\% over $3-5$  disk thicknesses suffices to cancel  out the total torque. {As a consequence the type I  migration of  low mass objects  reaching the jump  should be halted, and all these objects  should be trapped there provided some amount  of  dissipation  is  present  in the  disk  to  prevent  the corotation  torque  saturation.    As  dissipation  is  provided  by turbulence, which induces a jitter of the planet semi-major axis, we investigate  under which conditions  the trapping  process overcomes the trend  of turbulence to  induce stochastic migration  across the disk. We  show that  a cavity  with a large  outer to  inner surface density ratio efficiently traps embryos from $1$ to $15$~$M_\\oplus$, at any radius up to $5$~AU  from the central object, in a disk which has same  surface density profile  as the Minimum Mass  Solar Nebula (MMSN).  Shallow surface density  transitions require light disks to efficiently trap embryos. In the case of the MMSN, this could happen in the very  central parts ($r < 0.03$~AU).}  We  discuss where in a protoplanetary  disk   one  can   expect  {a   surface  density jump}. This  effect could  constitute a solution  to the  well known problem that  the build up of  the first protogiant solid  core in a disk takes much longer than its type I migration towards the central object. ",
        "introduction": "} The migration of low  mass protoplanets ($M_p<15\\;M_\\oplus$) under the action of disk tides is long known  to be a fast process in disks with power  law  surface  density  profiles \\citep{w97,tanaka}.   The  fast migration  timescale estimates  of  these objects  even constitutes  a bottleneck for the core accretion scenario, which implies a slow build up  of  a  solid core  until  it  reaches  the mass  threshold  ($\\sim 15\\;M_\\oplus$) above  which rapid  gas accretion begins.   Indeed, the solid core build up  time is $10^{6-7}$~yrs \\citep{petal96}, while the migration  timescale of  a $M_p=1\\;M_\\oplus$  planet  is $O(10^5)$~yrs \\citep{w97,tanaka} and  scales inversely proportionally  to the planet mass. The  existence of gaseous  giant planets at large  distances ($a \\sim  0.1-10$~AU)  from their  central  star  therefore constitutes  a puzzle. Recent work by \\cite{alibert05}  has shown that the core build up time scale  can be lowered by taking  migration into account, which prevents  the  depletion of  the  core  feeding  zone. However,  these authors  find that  the most  up  to date  type~I migration  timescale estimate, which includes three dimensional effects and the co-rotation torque \\citep{tanaka}, still needs to  be lowered by a factor $10-100$ in  order to  allow for  the solid  core survival.   The  total torque exerted by the  disk onto the planet can be split  into two parts: the differential Lindblad  torque, that corresponds  to the torque  of the spiral wake  that the planet excites  in the disk,  and the corotation torque,  exerted  by the  material  located  in  the planet  coorbital region.  The role  of the corotation torque has  often been overlooked in migration rate estimates. The two  main reasons for that is that it is harder to evaluate than  the differential Lindblad torque, and that it saturates (i.e. tends to  zero) in the absence of dissipation.  The corotation torque scales with the radial gradient of $\\Sigma/B$, where $\\Sigma$  is the disk  surface density  and $B$  is the  second Oort's constant,  or half  the disk  flow vorticity  vertical component  in a non-rotating  frame.   This  scaling  makes the  corotation  torque  a quantity  very  sensitive to  local  variations  of  the disk  surface density or rotation  profile. Here we investigate the  behavior of the total (Lindblad + corotation) tidal  torque exerted on a planet in the vicinity of a  surface density radial jump, in  order to investigate a suggestion  by \\citet{m02}  that localized,  positive  surface density jumps  may be  able to  halt migration.   We assume  that  the surface density  transition  occurs on  a  length  scale  $\\lambda$ of  a  few pressure scale heights $H$.  We consider the case in which the surface density is  larger on  the outside  of the transition,  but we  do not limit ourselves to the case where the surface density on the inside is negligible compared  to its value  on the outer  side.  The case  of a virtually  empty  central  cavity  has already  been  contemplated  by \\cite{KL02}  in  the  context,  different  of ours,  of  giant  planet migration.   They conclude that  giant planet  migration is  halted or considerably slowed down  once the planet is inside  the cavity and in $2:1$ resonance  with the  cavity edge, as  beyond this  resonance the disk torque onto a planet on a circular orbit becomes negligible.  In section~\\ref{sec:analytic}  we provide simple  analytical estimates of  the   Lindblad  and  corotation  torques  at   a  surface  density transition.  We show  that the corotation torque, which  is a positive quantity  there, is  likely to  overcome  the Lindblad  torque if  the transition  is localized enough,  i.e.  $\\lambda  < $  a few  $H$.  In section~\\ref{sec:setup} we  describe the numerical setup  that we used to   check  this   prediction   with  a   numerical  hydro-code.    In section~\\ref{sec:num},  we  present   the  results  of  our  numerical simulations  which indeed  exhibit for  a wide  range of  parameters a fixed point at  the transition, i.e. a point  where the corotation and Lindblad  torques  cancel each  other  and  where planetary  migration stops.   {We  also discuss  in  this  section  the issue  of  the saturation  of the  corotation torque  and the  need of  turbulence to prevent it, and  the conditions under which turbulence  is able or not to  unlock  a  planet  from   the  transition.}  We  then  discuss  in section~\\ref{sec:discuss} where  in protoplanetary disks  such surface density transitions  can be  found, and what  are the  consequences of these planet traps on giant planet formation. ",
        "conclusions": "We have shown that the  type I migration of $M_p<\\;15M_\\oplus$ planets can be  halted at radially  localized disk surface density  jumps (the surface density  being larger on the  outside of the  jump). {Some turbulence  is   then  required  to  prevent   the  corotation  torque saturation.  The  latter tends to  induce stochastic migration  of the embryos, and therefore tends to  act against the trap. The lighter the disk, the  easier it is  to keep an  embryo trapped, since  its radial jitter over the timescale needed  to reach torque convergence is small and  can  be  confined  to  the  vicinity of  the  trap  stable  fixed point. Inversely, the lighter the  embryo, the more difficult it is to trap it:  in the zero mass limit,  the embryo has no  tidal torque and only feels  the torque fluctuations  arising from turbulence.  We have shown that a jump with a large surface density ratio efficiently traps embryos above one Earth mass up even in relatively massive disks, such as  the MMSN  at  $r=5$~AU,  while shallow  jumps  (a $40$~\\%  surface density increase over $\\sim 4$ disk scale heights) will retain embryos of a few Earth  masses and more only in light disks,  such as the MMSN at $r=0.1$~AU.}  We have illustrated by means of numerical simulations that  a surface  density jump  traps in-falling  embryos, and  that it drives them  along with it if  it moves radially as  the disk evolves. Such a jump can occur at different locations in a protoplanetary disk: it can be found at  the transition between the standard accretion disk solution (SAD)  and the jet  emitting disk (JED),  thought to be  at a distance $0.3$  to $1-2$~AU  from the central  object. It can  also be found at the inner edge of a dead zone, where one expects a very large ratio of  outer to inner surface  densities.  It may also  be found at the outer  edge of  a protoplanetary gap  if one giant  planet already exists in the  disk. We have stressed that  these planetary traps have two important consequences: (i) they  can hold the type~I migration of protogiant cores  over the timescale  needed to overcome  the critical mass ($\\sim 15\\;M_\\oplus$) over  which rapid gas accretion begins onto these cores, which are then  quickly converted into giant planets that should enter the slow, type~II migration regime; (ii) as solid embryos accumulate  at the  trap, the  build up  of critical  cores  should be faster than elsewhere in the disk. Finally, we mentioned that although temperature radial  jumps in principle  also yield large peaks  of the corotation  torque, they  are unlikely  to be  able to  counteract the Lindblad torque  by themselves, as the  temperature gradients required for that are unrealistically large in a standard accretion disk.  \\appendix"
    },
    "0601/astro-ph0601708_arXiv.txt": {
        "abstract": "Using the Swift data of GRB 050315, we progress in proving the uniqueness of our theoretically predicted Gamma-Ray Burst (GRB) structure as composed by a proper-GRB, emitted at the transparency of an electron-positron plasma with suitable baryon loading, and an afterglow comprising the ``prompt radiation'' as due to external shocks. Detailed light curves for selected energy bands are theoretically fitted in the entire temporal region of the Swift observations ranging over $10^6$ seconds. ",
        "introduction": "GRB 050315 \\citep{va05} has been triggered and located by the BAT instrument \\citep{b04,ba05} on board of the {\\em Swift} satellite \\citep{ga04} at 2005-March-15 20:59:42 UT \\citep{pa05}. The narrow field instrument XRT \\citep{bua04,bua05} began observations $\\sim 80$ s after the BAT trigger, one of the earliest XRT observations yet made, and continued to detect the source for $\\sim 10$ days \\citep{va05}. The spectroscopic redshift has been found to be $z = 1.949$ \\citep{kb05}. We present here the first results of the fit of this source in the framework of our theoretical model and point out the new step toward the uniqueness of the explanation of the overall GRB structure made possible by the Swift data of this source. ",
        "conclusions": "In view of the above results, which clearly fits the overall luminosity in fixed energy bands, we will return in a forthcoming publication to the identification of the P-GRB using our theoretically predicted values for both its intensity and its time lag relative to the afterglow peak. We will also address the spectral analysis which is a most powerful theoretical prediction in order to evidence the continuity between the ``prompt radiation'' and the late phases of the afterglow and so to prove the uniqueness of the overall GRB structure.  \\begin{theacknowledgments} We thank P. Banat, G. Chincarini, A. Moretti and S. Vaughan for their help in the analysis of the observational data. \\end{theacknowledgments}"
    },
    "0601/gr-qc0601053_arXiv.txt": {
        "abstract": "We study gravitational lensing by compact objects in gravity theories that can be written in a Post-Post-Newtonian (PPN) framework: i.e., the metric is static and spherically symmetric, and can be written as a Taylor series in $\\gravr/r$, where $\\gravr$ is the gravitational radius of the compact object. Working invariantly, we compute corrections to standard weak-deflection lensing observables at first and second order in the perturbation parameter $\\vep = \\vthbh/\\vthE$, where $\\vthbh$ is the angular gravitational radius and $\\vthE$ is the angular Einstein ring radius of the lens.  We show that the first-order corrections to the total magnification and centroid position vanish universally for gravity theories that can be written in the PPN framework.  This arises from some surprising, fundamental relations among the lensing observables in PPN gravity models.  We derive these relations for the image positions, magnifications, and time delays. A deep consequence is that any violation of the universal relations would signal the need for a gravity model outside the PPN framework (provided that some basic assumptions hold). In practical terms, the relations will guide observational programs to test general relativity, modified gravity theories, and possibly the Cosmic Censorship conjecture.  We use the new relations to identify lensing observables that are accessible to current or near-future technology, and to find combinations of observables that are most useful for probing the spacetime metric.  We give explicit applications to the Galactic black hole, microlensing, and the binary pulsar J0737$-$3039. ",
        "introduction": "\\label{sec:intro}  Gravitational lensing is now a central field of astronomy with wide-ranging applications that relate to extra-solar planets, dark matter substructures, and cosmological parameters (including dark energy) \\cite{SEF,petters,Saas-Fee}. Theoretical studies have also explored lensing by black holes within the context of general relativity \\cite{darwin}-\\cite{PPN-2}, braneworld gravity \\cite{bhadra}-\\cite{majumdar}, and string theory \\cite{gibbons,paperI}. Most black hole lensing studies have focused on ``relativistic'' images corresponding to light rays that loop around a black hole, which probe gravity in the strong-deflection limit, but are extremely difficult to detect observationally \\cite{virbhadra,petters-SgrA}.  We are proposing and evaluating new possibilities for using gravitational lensing by compact objects (including black holes) to test various theories of gravity using current or near-future technology.  In Paper I of this series \\cite{paperI}, we introduced an analytic framework for studying gravitational lensing by a compact deflector with mass $\\Mbh$, in which the lensing scenario satisfies three basic assumptions: \\begin{itemize}  \\item[\\assump{1}] The gravitational lens is compact, static, and spherically symmetric, with an asymptotically flat spacetime geometry sufficiently far from the lens \\footnote{In lensing, asymptotic flatness can be generalized to a Robertson-Walker background by using angular diameter distances and including appropriate redshift factors.}.  The spacetime is vacuum outside the lens and flat in the absence of the lens.  \\item[\\assump{2}] The observer and source lie in the asymptotically flat regime of the spacetime.  \\item[\\assump{3}] The light ray's distance of closest approach $r_0$ and impact parameter $b$ both lie well outside the gravitational radius $\\gravr = G \\Mbh/c^2$, namely, $\\gravr/r_0 \\ll 1$ and $\\gravr/b \\ll 1$.  The bending angle can then be expressed as a series expansion in $\\gravr/b$, as follows: \\beq \\label{eq:bangle-ser} \\hat{\\alpha}(b) = A_1 \\left(\\frac{\\gravr}{b}\\right) \\ + \\ A_2 \\left(\\frac{\\gravr}{b}\\right)^2 \\ + \\ A_3 \\left(\\frac{\\gravr}{b}\\right)^3 \\ + \\ \\order{\\frac{\\gravr}{b}}{4} . \\eeq The coefficients $A_i$ are independent of $\\gravr/b$, but may include other fixed parameters of the spacetime.  Since $b$ and $\\gravr$ are invariants of the light ray, (\\ref{eq:bangle-ser}) is independent of coordinates.  Note that the subscript of $A_i$ conveniently indicates that the component is affiliated with a term of order $i$ in $\\gravr/b$.  \\end{itemize} We applied the lensing framework to gravity theories that can be written in a Post-Post-Newtonian (PPN) expansion, meaning that the metric can be written as a Taylor series in $\\gravr/r$. Working invariantly, we computed the observable properties of the primary (positive-parity) and secondary (negative-parity) lensed images.  Assuming that \\assump{3} holds, we wrote the lensing observables as series expansions in $\\vep = \\vthbh/\\vthE$, where $\\vthbh$ is the angular gravitational radius and $\\vthE$ is the angular Einstein ring radius.  For gravity theories that agree with general relativity in the weak-deflection limit ($A_1 = 4$ in eqn.~\\ref{eq:bangle-ser}), the zeroth-order terms in the $\\vep$ expansions give the familiar results for lensing by a point mass in the weak-deflection limit of general relativity.  The higher order terms give corrections to the lensing observables, which differ for different gravity theories. We studied general third-order PPN models, which allowed us to compute the weak-deflection lensing results plus the first two correction terms (order $\\vep$ and $\\vep^2$).  In Paper I we found the interesting result that the first-order corrections to two lensing observables --- the total magnification and the magnification-weighted centroid position --- vanish in PPN gravity models that agree with general relativity in the weak-deflection limit (i.e., $A_1 = 4$).  More generally, we found that these corrections depend on $A_1-4$, suggesting that they nearly vanish in gravity theories that agree only approximately with general relativity in the weak-deflection regime (i.e., $A_1 \\approx 4$). Note that existing observations constrain this parameter to $A_1 = 3.99966 \\pm 0.00090$ \\cite{sh}.  Working more generally than in Paper I, we have now discovered the surprising result that the first-order corrections to the total magnification and centroid in fact vanish exactly in all PPN models.  This depends on precise cancellations that are striking because the PPN framework covers quite a broad range of gravity theories.  Understanding why the cancellations occur has led us to identify some new fundamental relations between lensing observables in the PPN framework.  Some of the relations are universal for PPN models, in the sense that they hold for all values of the invariant PPN parameters of the light bending angle.  These new lensing relations will play key roles in planning observing missions to test theories of gravity.  First, the relations allow us to determine which observables are most accessible to current or near-future instrumentation.  We give explicit applications to the Galactic black hole, Galactic microlensing, and the binary pulsar J0737-3039. Second, the lensing relations help us identify combinations of lensing observables that are most useful for probing the spacetime metric by constraining invariant PPN parameters. One of the measurable parameters is connected to the existence of naked singularities in certain gravity models (see Paper I), so we may have an observational test of the Cosmic Censorship conjecture \\cite{CosmicCensorship}. Third, the universal relations provide a powerful means to test the entire PPN framework.  Any violation of relations that are universal in the PPN framework would suggest that a fundamentally different theory of gravity is needed (provided that assumptions \\assump{1}--\\assump{3} hold).  In this paper we present the new lensing relations and use them to assess prospects for using lensing observations to test PPN gravity theories.  Section~\\ref{sec:formalism} reviews the third-order PPN lensing framework.  Section~\\ref{sec:relations} derives the new relations between the image positions, magnifications, and time delays, and discusses their conceptual implications.  Section~\\ref{sec:applications} considers applications of the relations to various astrophysical settings. ",
        "conclusions": "\\label{sec:conclusions}  Using gravitational lensing by compact objects, we have presented new propects for testing theories of gravity within the PPN framework.  In this paper we generalized the PPN lensing formalism from Paper I to include fully general third-order PPN models.  We determined the weak-deflection limits plus first- and second-order corrections in $\\vep = \\vthbh/\\vthE$ for observable properties of lensed images (positions, magnifications, and time delays).  During the PPN analysis, we discovered some surprising new fundamental relations between lensing observables.  Some of the relations are universal for the entire family of PPN gravity models.  A deep conceptual implication is that any observed violation of the universal lensing relations (given that assumptions \\assump{1}--\\assump{3} apply) would indicate that a fundamentally different theory of gravity is at work --- one outside the PPN framework.  The new relations have enabled us to identify combinations of lensing observables that are key to probing the spacetime metric by constraining the invariant PPN parameters of the light bending angle. The parameter $A_2$ is related to the existence of naked singularities in certain gravity models (see Paper I), so constraining $A_2$ also provides a possible observational test of the Cosmic Censorship conjecture.  The new lensing relations will, in other words, play important roles in planning observing missions to test theories of gravity.  In a practical application, we identified lensing observables that are accessible to current or near-future instrumentation, considering three likely lensing scenarios: the Galactic black hole, Galactic microlensing, and the binary pulsar J0737$-$3039.  A noted application of the new lensing relations is the ability to find combinations of observables that will yield a direct method for knowing when observations are probing beyond the standard weak-deflection regime."
    },
    "0601/hep-th0601203_arXiv.txt": {
        "abstract": "Positively-curved, oscillatory universes have recently been shown to have important consequences for the pre-inflationary dynamics of the early universe.  In particular, they may allow a self-interacting scalar field to climb up its potential during a very large number of these cycles. The cycles are naturally broken when the potential reaches a critical value and the universe begins to inflate, thereby providing a `graceful entrance' to early universe inflation. We study the dynamics of this behaviour within the context of braneworld scenarios which exhibit a bounce from a collapsing phase to an expanding one. The dynamics can be understood by studying a general class of braneworld models that are sourced by a scalar field  with a constant potential. Within this context, we determine the conditions a given model must satisfy for a graceful entrance to be possible in principle. We consider the bouncing braneworld model proposed by Shtanov and Sahni and show that it exhibits the features needed to realise a graceful entrance to inflation for a wide region of parameter space. ",
        "introduction": "The inflationary scenario is presently the favoured model for large-scale structure formation in the universe \\cite{inf1,inf2,inf3,inf4,inf5,linde83}. (For reviews, see \\cite{llkcba,lythriotto,lidlyth,btw}). Inflation arises whenever the universe undergoes a phase of accelerated expansion. Within the context of relativistic cosmology, a necessary and sufficient condition for acceleration is a violation of the strong energy condition, i.e., $\\rho +3 p< 0$. This is equivalent to the condition $\\gamma < 2/3$, where $\\gamma \\equiv (\\rho+ p)/\\rho$ defines the equation of state parameter. In the simplest versions of the scenario, such a condition is realised by a scalar `inflaton' field that is rolling sufficiently slowly down its self-interaction potential, $\\dot{\\phi}^2 < V$ \\cite{linde83}.  Given the recent successes of inflation when confronted with observations of the Cosmic Microwave Background (CMB) radiation and high redshift surveys \\cite{wmap}, a key question to address is how the initial conditions for successful inflation were established in the very early universe. This question was recently examined within the context of Loop Quantum Cosmology (LQC), which is the application of Loop Quantum Gravity to symmetric states. (For reviews, see \\cite{loopreview,loop1,bojoreview}). In this scenario, semi-classical modifications to the standard Friedmann equations cause a minimally coupled scalar field to behave as an effective phantom fluid below a critical value of the scale factor, $a_*$, with an effective equation of state, $\\gamma_{\\rm eff} < 0$ \\cite{Bojowald2002,lmnt,mntl,mtle,nunes}. This causes a collapsing, isotropic, spatially closed Friedmann-Robertson-Walker (FRW) universe to undergo a non-singular bounce into an expansionary phase. On the other hand, if the field satisfies the strong energy condition above $a_*$ (where classical dynamics is recovered), such a universe will eventually recollapse. The overall result, therefore, is that the universe oscillates between cycles of contraction and expansion \\cite{lmnt}.  If the field's kinetic energy dominates its potential energy during these cycles, it will effectively behave as a massless field with its value either increasing or decreasing monotonically with time. In principle, therefore, the field may gradually roll {\\em up} its potential (assuming implicitly that it is evolving along a region of the potential that is increasing in magnitude). As a result, the potential energy will gradually increase during each successive cycle, with the net result that it becomes progressively harder to satisfy the strong energy condition on scales above $a_*$. Eventually, therefore, a cycle will be reached when the strong energy condition becomes violated during the expansionary phase of the cycle ($a>a_*$), and this will initiate a phase of inflation. This implies that slow-roll inflation may be possible even if the field is initially located in a region of the potential that would not generate accelerated expansion, such as a minimum. This effect has been termed a `graceful entrance' to inflation \\cite{nunes}.  A key feature of the dynamics described above is that the universe is able to oscillate because the effective equation of state satisfies $\\gamma_{\\rm eff} < 2/3$ for $a<a_*$ and $\\gamma_{\\rm eff} >2/3$ for $a>a_*$. Since these are rather weak conditions, it is important to investigate whether similar effects are possible in other cosmological scenarios, such as the braneworld paradigm. This is the purpose of the present paper. The braneworld scenario is motivated by string/M-theory \\cite{string1,string2,string3} and has attracted considerable attention in recent years. (See \\cite{royreview,braxreview,braxreview1} for reviews). Our observable four-dimensional universe is interpreted as a co-dimension one brane propagating in a five- (or higher-) dimensional `bulk' space. A number of bouncing braneworld models have been developed to date \\cite{shsh,burgess,bounce1,bounce2,bounce3,bounce4}. In the model proposed by Shtanov and Sahni (S-S), for example, the extra bulk dimension is timelike and this results in modifications to the effective four-dimensional Friedmann equation that induce a non-singular bounce \\cite{shsh}.  The paper is organized as follows. In Section II, we outline how a phase space analysis for a field with a constant potential can yield valuable insight into the cosmic dynamics that leads to a graceful entrance for inflation. We then proceed in Section III to determine the necessary and sufficient conditions for graceful entrance when the Friedmann equation has an arbitrary dependence on the energy density of the universe. We consider the S-S braneworld scenario as an explicit example in Sections IV--VI and determine the regions of parameter space where a graceful entrance is possible. We conclude with a discussion in Section VII. ",
        "conclusions": "In this paper, we have examined the criteria that a cosmological model must satisfy in order for it to undergo a `graceful entrance' into a phase of inflation, whereby a scalar field is able to move up its interaction potential whilst the universe undergoes a large number of oscillatory cycles. Our approach has been to consider an idealized model, where the matter is comprised of a scalar field evolving along a constant potential. This system provides a good approximation to more realistic scenarios where the potential is field-dependent if the change in the potential is sufficiently small over a large number of cycles. It therefore provides us with a methodology for identifying whether a particular cosmological scenario will meet the necessary requirements for graceful entrance.  The mechanism we have outlined is very generic. The important ingredients are that the phase space should exhibit both a centre and saddle equilibrium point when the magnitude of the potential $V$ falls below a critical value $V_{\\rm crit}$. These points gradually move towards each other as $V$ increases and eventually merge when $V=V_{\\rm crit}$. Above this scale, the points should disappear from the phase space. From the physical viewpoint, this behaviour arises because the potential now dominates the field's kinetic energy, $V > \\dot{\\phi}^2$, and consequently drives a phase of inflationary expansion. It follows that the critical value $V_{\\rm crit}$ sets the energy scale for the onset of inflation.  We have developed a framework for investigating a general class of braneworld models in this context that are characterized by a Friedmann equation with an arbitrary dependence on the energy density. Such a class of models can be represented as a standard, relativistic Friedmann universe by reinterpreting the effective equation of state of the matter on the brane. As a concrete example, we focused on the Shtanov-Sahni braneworld \\cite{shsh}, where the extra bulk dimension is timelike. We found that the question of whether the braneworld oscillates or expands indefinitely is dependent on the field's kinetic and potential energies, as well as the dark radiation and the brane tension, $\\sigma$. If the dark radiation is negligible, the phase space has a similar structure to that outlined above and a graceful entrance to inflation can therefore be realised, where the energy scale at the onset of inflation is given by $V_{\\rm crit} \\approx \\sigma$. The dark radiation can have a significant effect on the dynamics, however, and the question of whether a graceful entrance is possible when this term is present is sensitive to the initial conditions.  In conclusion, therefore, we have presented a non-singular, oscillating braneworld that can in principle exhibit a graceful entrance to inflation. We anticipate that such behaviour will apply for a wide range of collapsing braneworld scenarios that are able to undergo non-singular bounces (see, e.g., \\cite{burgess}) and, moreover, will not be strongly dependent on the precise form of the inflaton potential. Similar dynamics was recently shown to arise in models inspired by loop quantum cosmology \\cite{lmnt,mntl,mtle} and it is interesting that such behaviour is also possible in braneworld models motivated by string/M-theory.  Finally, our analysis also provides the basis for realising the recently proposed emergent universe scenario in a braneworld context \\cite{em,mtle}. In this scenario, the initial state of the universe is postulated to be either an Einstein static universe or one that is oscillating about such a solution in the infinite past. The potential has an asymptotically flat section, where $V<V_{\\rm crit}$ as $\\phi  \\to -\\infty$, but subsequently increases above $V_{\\rm crit}$ once the value of the field has increased beyond a certain value. If the field is initially located in the plateau region of the potential and moving in the appropriate direction, it will eventually reach the section of the potential which rises above $V_{\\rm crit}$, thus initiating inflation. Within the framework of classical Einstein gravity, the emergent universe scenario suffers from severe fine-tuning, since the static universe corresponds to a saddle point in the phase space. However, a mechanism which allows for a graceful entrance to inflation provides a natural realization of this scenario with less fine tuning and a greater degree of freedom in the choice of inflationary potential."
    },
    "0601/astro-ph0601258_arXiv.txt": {
        "abstract": "We explore the hypothesis that the magnetic fields of neutron stars are of fossil origin. For parametrised models of the distribution of magnetic flux on the Main Sequence and of the birth spin period of the neutron stars, we calculate the expected properties of isolated radio pulsars in the Galaxy using as our starting point the initial mass function and star formation rate as a function of galacto-centric radius.  We then use the 1374 MHz Parkes Multi-Beam Survey of isolated radio pulsars to constrain the parameters in our model and to deduce the required distribution of magnetic fields on the main sequence.  We find agreement with observations for a model with a star formation rate that corresponds to a supernova rate of 2 per century in the Galaxy from stars with masses in the range $8 - 45\\Msun$ and predict 447,000 active pulsars in the Galaxy with luminosities greater than 0.19 mJy~kpc$^2$. The progenitor OB stars have a field distribution which peaks at $\\sim 46$~G with $\\sim 8$ percent of stars having fields in excess of 1,000~G. The higher field progenitors yield a population of 24 neutron stars with fields in excess of $10^{14}$~G, periods ranging from 5 to 12 seconds, and ages of up to 100,000 years, which we identify as the dominant component of the magnetars. We also predict that high field neutron stars ($\\log B>13.5$) originate preferentially from higher mass progenitors and have a mean mass of $1.6\\Msun$, which is significantly above the mean mass of $1.4\\Msun$ calculated for the overall population of radio pulsars. ",
        "introduction": "A striking feature of the white dwarfs and the neutron stars is the wide range of magnetic field strengths seen in each of these groups. In isolated magnetic white dwarfs (about 15 percent of the total white dwarf population), the field strengths are measured directly through the Zeeman effect and range from $\\sim 10^5 - 10^9$~G.  In isolated radio pulsars, the field estimates are in the range $\\sim 10^{11} - 10^{14}$~G and are based on the measured spin down rates and the assumption of dipole radiation. Observations in the X and $\\gamma$-ray spectral regions have revealed the presence of a class of neutron stars with higher inferred fields, namely the Anomalous X-ray Pulsars (AXPs) and the Soft Gamma Repeaters (SGRs). These neutron stars are often referred to as ``magnetars'' and have estimated field strengths in the range $\\sim10^{14} - 10^{15}$~G.  Possible explanations for the magnetic fields in neutron stars are that the fields are either (i) essentially of fossil origin (e.g. Rudermann 1972) or (ii) generated by a convective dynamo (Thompson \\& Duncan 1993). At the present time, this issue remains unresolved (see Mestel, 1999, for a comprehensive review).  In the case of the white dwarfs, a direct link can be made between magnetism in the compact star phase and the Main Sequence via the well documented properties of the chemically peculiar Ap and Bp stars (Wickramasinghe \\& Ferrario 2005) and has been a major factor supporting (i) for the white dwarfs.  Magnetism is not as well documented in the more massive ($\\gsimeq 8 \\Msun$) upper Main Sequence stars that evolve into neutron stars. The only known stars with a directly measured magnetic field are $\\theta$ Orion C, with a dipolar magnetic field $B_p = 1,100$~G (Donati et al. 2002), and HD191612 with a field of $B_p = 1,500$~G (Donati et al. 2005).  Their mass is estimated around $40 \\Msun$, with the latter star slightly more evolved than the former. Interestingly, the magnetic fluxes of both these stars ($1.1\\times 10^{27}$ G~cm$^2$ for $\\theta$ Orion C and $7.5\\times 10^{27}$ G~cm$^2$ for HD191612) are comparable to the flux of the highest field magnetar SGR~1806-20 ($5.7\\times 10^{27}$ G~cm$^2$). There is also indirect evidence from studies of stellar winds from massive stars that magnetism may be widespread among OB stars (e.g. Wade 2001). These facts, taken together with the scaling that exists between the dipolar field strengths of the white dwarfs and the neutron stars (Figure 1), suggest that a prima facie case can be made for a fossil origin of fields also in the neutron stars. The fossil field hypothesis provides a natural explanation for the wide range of observed magnetic field strengths in a given class as being the result of inhomogeneities in the interstellar magnetic field strength in star forming regions.  Calculations of pre-Main Sequence evolution have shown that stars more massive than $\\sim 2\\msun$ are likely to begin their Main Sequence phase with a primordial magnetic flux entrapped in radiative regions supporting the fossil hypothesis (Tout, Livio \\& Bonnell 1999; Moss 2003).  There are no detailed studies on the interplay between fossil fields and subsequent stellar evolution. This problem has been discussed by Tout, Wickramasinghe \\& Ferrario (2004, hereafter TWF). They argued that provided there are regions of the star that are radiative during subsequent evolution, a fossil Main Sequence poloidal magnetic flux could survive through to the compact star phase.  \\begin{figure} \\begin{center} \\hspace{0.1in} \\epsfxsize=0.98\\columnwidth \\epsfbox[30 154 577 512]{NS_WD_distr.ps} \\caption{Black histogram (solid line): Field distribution of neutron stars. Red histogram (dashed line): Field distribution of High Field Magnetic White Dwarfs when their radii are shrunk to that of a typical neutron star radius of $\\sim 10^6$~cm.} \\end{center} \\end{figure}  In this paper, we investigate the consequences of the fossil field hypothesis for the origin of magnetic fields in neutron stars by carrying out population synthesis calculations for different assumptions on the distribution of the magnetic flux of massive ($\\gsimeq 8 \\Msun$) Main Sequence stars and of the field dependence of the initial birth period of neutron stars. We model the observed properties of the population of isolated radio pulsars in the 1374 MHz Parkes Multi-Beam Survey (PMBS, Hobbs et al. 2004; Kramer et al. 2003; Morris et al. 2002; Manchester et al. 2001) to constrain our model parameters, and use these to deduce the required magnetic properties of the progenitor OB stars. ",
        "conclusions": "Our approach has been to focus on constraining the key parameters that we have introduced for describing the magnetic flux distribution on the Main Sequence and the distribution of the initial birth period assuming that the remaining parameters of the model are within the ranges given by previous investigators. We have used as our basic observational constraints the 1-D projections of the data comprising the number distributions in period $P$, magnetic field $B_{NS}$, period derivative $\\dot P$ and radio luminosity $L_{1374}$, noting that these distributions are not all independent.  \\begin{figure*} \\begin{center} \\hspace{0.1in} \\epsfxsize=0.8\\textwidth \\epsfbox[18 439 570 700]{model_fit1.ps} \\caption{Our model (blue) overlapped to the PMBS isolated radio pulsars data from the ATNF catalogue (http://www.atnf.csiro.au/research/pulsar/psrcat). The errors are given by $\\sqrt{N}$, where $N$ is the number of objects in each bin.} \\end{center} \\end{figure*}  The best fit model to the observations of the PMBS pulsars was determined ``by eye'' after conducting hundreds of trials. We have found this method more reliable than the standard K-S statistic. Our results are shown in Figure 2.  The initial period distribution is described by $P_a= 0.22$~s, $\\sigma_{P_0}=P_0/2$ and $\\alpha = 0.01$.  The magnetic flux distribution is described by two Gaussians of means and spreads $\\log\\Phi_1=26.2$, $\\sigma_{\\log\\Phi1}=1.0$, and $\\log\\Phi_2= 25.2$, $\\sigma_{\\log\\Phi2} =0.61$ respectively, weighted in the ratio $1:5$. The parameters for the luminosity model are $a=1.8$ and $b=2.8$.  This model reproduces the observed total number of pulsars in the PMBS, once we exclude the millisecond pulsars ($P<20$~ms). The star formation rate is normalised to $2.5\\Msun$ pc$^{-2}$ Gyr$^{-1}$ at the galacto-centric radius of the Sun, which is within the uncertainty in the Boissier et al. (2003) relationship. This corresponds to a galactic supernova rate of 2 per century.  In a broad sense, the distribution that we use for the magnetic flux has the strongest influence on the location of the peak in the magnetic field distribution of the pulsars, while the mean period $P_a$ and the width of the period distribution $\\sigma_{P_0}$ have a strong influence on the period distribution, although clearly period and field distributions are inter-linked through period evolution which depends on the magnetic field. We find that we cannot model both the field distribution and the period distribution at the same time if we assume that all pulsars are born at a period less than $50$~ms. Such models give an excess of stars at low periods which is not consistent with the results of Vranesevic et al. (2004) who find that up to 40 percent of pulsars are born with periods in the range 0.1-0.5~s. Regimbau \\& Freitas Pacheco (2001) also noted that most pulsars are born with $P>0.1$~s and conduct their population synthesis studies using a Gaussian with a mean birth period $P_0=0.29$~s and a spread of 0.1~s while Faucher-Gigu\\`ere \\& Kaspi (2005) use a Gaussian with a mean birth period $P_0=0.3$~s and a spread of 0.15~s. All these results contradict the view that all pulsars are born as fast rotators ($P_0\\le 0.1$~s) which arises from observational selection effects, since young and fast neutron stars (e.g. the Crab pulsar) would be more easily detectable as radio pulsars than the more slowly rotating ones.  We note that with our model parameters, the pulsar birth spin period increases with increasing magnetic field, but is almost independent of the field in the range $\\log B_{\\rm NS}(G)=10-13$ where the vast majority of isolated radio pulsars lie. The magnetic field dependence becomes important in the high-field radio pulsars and magnetar field range.  The observations show that the field distribution is intrinsically asymmetric with a low field tail that is more prominent than the high field tail. This general behaviour is reproduced by our assumed bi-Gaussian flux distribution. Furthermore, the model predicts 26 pulsars with fields $> 10^{13}$ G, in good agreement with the PMBS which shows 23 stars in the same field range.  There have been many previous attempts at synthesising the population of radio pulsars with different assumptions on neutron star magnetic fields, field decay and initial periods.  Some of these investigations have led to the conclusion that the neutron star field distribution cannot be described by a single Gaussian (e.g. Gonthier et al. 2002), and we find a similar result for the magnetic fluxes of their progenitor stars.  The initial-final mass radius relationship that we have assumed implies that higher mass progenitors will produce higher mass neutron stars. These stars will, on average, also have higher magnetic fluxes, and therefore tend to produce higher field neutron stars. However, the latter effect is counterbalanced by the strongly declining initial mass function which favours low mass star progenitors, which can also produce high field neutron stars in the Gaussian tail of the magnetic flux distribution. In order to investigate the mass distribution of active radio pulsars in the Galaxy with periods less than 20~s, we have divided them into two field groups ($\\log B_{NS}(G) = 10 - 13.5$) and ($\\log B_{NS}(G) > 13.5$). Our results are shown in Figure 3 for pulsars with 1374 MHZ luminosities $\\ge 0.19$ mJy~kpc$^2$. The mean neutron star masses in these groups are $1.45\\Msun$ and $1.6\\Msun$ respectively, whilst the mean masses of their progenitors are $12 \\msun$ and $19 \\msun$ respectively. The stars in the highest field bin have a high mass tail with 54 percent of the stars extending beyond $1.5\\Msun$ compared to 19 percent of stars in the same mass range in the lower field bin. We find that the total number of active pulsars with 1374 MHZ luminosities $>0.19$ mJy~kpc$^2$ is 447,000. If we apply the beaming model of equation \\ref{beam}, we find that the number of active pulsars whose beams intercept Earth reduces to 48,000.  \\begin{figure} \\begin{center} \\hspace{0.1in} \\epsfxsize=0.95\\columnwidth \\epsfbox[36 313 379 642]{massesNS1.ps} \\caption{Masses of neutron stars with rotational periods of up to 20 s in the two indicated field ranges.} \\end{center} \\end{figure}  The magnetars are neutron stars that are discovered as X-ray and/or $\\gamma$-ray sources. They are currently confined to the period range $5 - 12$~s and are young objects (e.g. Gaensler et al. 2001). If one makes the standard assumption of dipole radiation with a braking index of 3, the magnetic fields are in the range $\\log B_{\\rm NS}(G) =13.8-15.3$.  We have used our calculations to investigate if the high field tail of the neutron stars in our model can provide an explanation for the numbers of magnetars that are seen in the Galaxy. Here we use as our cutoff $B_{\\rm NS}=10^{14}$~G, since this value is just above the revised quantum critical field of Zhang \\& Harding (2000) for radio emission. If we isolate the total number of neutron stars with fields in excess of $10^{14}$~G, and ages of less than 500,000 years, we find 146 such stars, none of which would be detectable as radio-pulsars. If we now limit the age to 100,000 years, our modelling yields 24 magnetars in the Galaxy with periods from $5$ to $12$~s and 8 magnetars with periods $>12$~s, 5 of which have periods between 12 and 15 seconds. Thus, the number declines rapidly with period outside the observed period range. We note in this context that some magnetars exhibit a large discrepancy between their real ages (as determined through their association with supernova remnants) and their characteristic ages (e.g. AXPs 1E2259+586 and 1E1845-0545, Gaensler et al. 2001). This discrepancy is readily explained through our assumption of a $P-B$ relationship, which sees high field neutron stars being born preferentially at long periods ($1-10$~s).  Below $B_{\\rm NS}=10^{14}$~G, there is an area of overlap where a couple of magnetars have fields and periods that are comparable to those of the high field radio pulsars (HBRPs) (e.g. Kaspi \\& McLaughlin 2005).  This suggests that there could exist lower field magnetars ($<10^{14}$~G) that eventually evolve into X-ray silent HBRPs of the type currently observed if their radio-beams intercept the Earth. We note that it should also be possible to discover ``hybrid'' young objects with $B_{\\rm NS}<10^{14}$~G exhibiting both magnetar and radio pulsar characteristics. In our analysis, we have counted all these objects as radio-pulsars if they pass through our PMBS filtering program.  Finally, we return to the recent evidence for high progenitor masses in the magnetars 1E~1048-5937 (Gaensler et al. 2005) and CXO~J164710.2-455216 (Muno et al. 2005). In the dynamo model, rapid spin ($\\sim 3$~ms) is required to generate super-strong fields and Heger et al. (2005) argue that it is only the cores of the more massive stars that have this property. This has been used as an argument in favour of the dynamo model, at least for the magnetars.  We note that these findings also support our fossil hypothesis, since we predict that magnetars, and more generally high field neutron stars, originate preferentially from high mass progenitors.  We conclude by noting that the fossil field hypothesis as formulated by us, and which does not allow for magnetic flux loss in the post main-sequence evolution, requires a very specific distribution of magnetic fields for massive Main Sequence stars. This is shown in Figure 4 for an assumed dipolar field structure. This distribution has a peak at $\\sim 46$~G and low and high field wings extending from $\\sim 1$~G to $\\sim 10,000$~G with 8 percent of stars having fields in excess of $\\sim 1,000$~G. The latter group are the progenitors of the highest field magnetars. The detailed shape of the Main Sequence distribution shown in Figure 4 is potentially an observable quantity, and a prediction of our fossil field model.  \\begin{figure} \\begin{center} \\hspace{0.1in} \\epsfxsize=0.98\\columnwidth \\epsfbox[14 312 385 507]{MSfields1.ps} \\caption{Predicted magnetic field distribution of massive stars ($8-45\\Msun$) on the Main Sequence.} \\end{center} \\end{figure}"
    },
    "0601/astro-ph0601291_arXiv.txt": {
        "abstract": "We present an integral field spectroscopy survey of rich clusters of galaxies aimed at studying their lensing properties.  Thanks to knowledge of the spectroscopic caracteristics of more than three families of multiple images in a single lens, one is able in principle to derive constraints on the geometric cosmological parameters. We show that this ambitious program is feasible and present some new results, in particular the redshift measurement of the giant arc in A2667 and the resdshift confirmation of the counter-image of the radial arc in MS2137--23. Prospects for the future of such program are presented. ",
        "introduction": "Gravitational lensing has been deeply renewed since 15 years now, and many extensive reviews are available in the litterature \\cite{schneider92}. So we simply remind here the basic lensing equation, which corresponds to a mapping of the image plane on the source plane \\[ \\vec \\beta = \\vec \\theta - \\vec \\nabla \\psi (\\vec \\theta) \\qquad \\qquad \\psi (\\vec \\theta) = \\frac{2}{c^2} \\frac{D_{LS}}{D_{OS}D_{OL}} \\varphi (\\vec \\theta) \\] where $\\varphi (\\vec \\theta)$ is the newtonian gravitational potential of the deflecting lens. The distances are angular-distances between the lens, the source and the observer. The Einstein radius, defined as \\[ R_E = \\sqrt{\\frac{4 G M}{c^2} \\ \\frac{D_{OL} D_{LS}}{D_{OS}}} = 2.6'' \\left(\\frac{\\sigma^2}{c^2}\\right) \\frac{D_{LS}}{D_{OS}} \\] represents the radius of the circular image in the ideal case of a source perfectly aligned behind a point source of mass $M$ or behind a singular isothermal sphere with velocity dispersion $\\sigma$. For any lensing mass, it represents the typical angular scale of the lensed images. Numerically, it is about $1-3''$ for a galaxy lens and $10-30''$ for a cluster of galaxies. In practice, gravitational lensing has many astrophysical applications, related to either the determination of the lensing potential and the mass distribution in galaxies and in clusters of galaxies, or the study of the distant galaxies thanks to the magnification of their images \\cite{mellier99}.  But it also depends on the geometrical cosmological parameters and represents an alternative possibility to constrain these parameters, independantly from other methods like supernovae or CMB fluctuations. This method was explored recently in details \\cite{golse02}.  It requires several sets of multiple images from different sources at different redshifts through the same lens. For each source the lens equation can be written as \\[ \\vec \\theta_S = \\vec \\theta_{I} - \\frac{2 \\sigma_0^2}{c^2} \\vec f(\\vec \\theta_I, \\theta_C, \\alpha, \\ldots ) \\times E(\\Omega_M, \\Omega_\\Lambda, z_L, z_S) \\] where the $f$ function includes all the caracteristics of the lensing potential and $E=D_{LS}/D_{OS}$ includes the cosmological dependence [$(\\Omega_M, \\Omega_\\Lambda)$ or $(\\Omega_M, w)$ for an euclidian universe] as well as the lens and source redshifts. Provided $z_L$ and $z_S$ are well known for two sets of multiple images, one can in principle solve the lens equation for both sets and constrain the E-term, with a well-defined degeneracy between the two geometrical cosmological parameters used \\cite{golse02}. However the cosmological dependence in the E-term is weak and better constraints require more than two sets of multiple images, spread in redshift as well as a very accurate mass determination in the lens (Fig. \\ref{a2218}).  \\begin{figure} \\centering \\includegraphics[height=4cm]{soucailF1a.ps} \\includegraphics[height=4cm]{soucailF1b.ps} \\caption{Left: Evolution of the E-term with redshift for different cosmological models, for a lens at $z=0.3$. Right: Confidence levels obtained from the modeling of the cluster-lens Abell 2218 \\cite{soucail04}} \\label{a2218}       % \\end{figure}  This method was first attemped in the cluster Abell 2218 ($z_L=0.17$,), a strong gravitational lens close to the ideal case: it displays at least 4 sets of multiple images with redshifts ranging from $z_{S1}=0.702$ to $z_{S4}=5.576$. An accurate mass model was developped \\cite{kneib96}, including 2 major mass components located on the two brightest galaxies, as well as the contribution of about 17 galaxy masses scaled by their luminosity.  This model prescription was constrained more accurately by including the 4 families of multiple images, and the likelihood distribution in the $(\\Omega_M, \\Omega_\\Lambda)$ plane finally displayed the expected degeneracy (Fig. \\ref{a2218}), reinforcing the validity of the mass model \\cite{soucail04}. ",
        "conclusions": "We have started an ambitious observing programme, well suited to 3D spectroscopy and aimed at studying in details a sample of strong gravitational lenses. The main difficulty is the faintness of the arcs which requires to push the present spectrographs to their limits. The results are therefore very sensitive to the data quality at the output of the telescope and to the data reduction which must be done with great care. These fundamental steps are in progress and we have demonstrated that faint objects 3D spectroscopy is feasible in rich environments like clusters of galaxies. Several scientific outputs from these data are presented, others are in Covone's talk (these proceedings) concerning the properties of the cluster galaxies."
    },
    "0601/astro-ph0601544_arXiv.txt": {
        "abstract": "We propose a novel parameterization of the dark energy density. It is particularly well suited to describe a non-negligible contribution of dark energy at early times and contains only three parameters, which are all physically meaningful: the fractional dark energy density today, the equation of state today and the fractional dark energy density at early times. As we parameterize $\\od(a)$ directly instead of the equation of state, we can give analytic expressions for the Hubble parameter, the conformal horizon today and at last scattering, the sound horizon at last scattering, the acoustic scale as well as the luminosity distance. For an equation of state today $w_0 < -1$, our model crosses the cosmological constant boundary. We perform numerical studies to constrain the parameters of our model by using Cosmic Microwave Background, Large Scale Structure and Supernovae Ia data. At $95\\%$ confidence, we find that the fractional dark energy density at early times $\\ode < 0.06$. This bound tightens considerably to $\\ode < 0.04$ when the latest Boomerang data is included. We find that both the gold sample of Riess et. al. and the SNLS data of Astier et. al. when combined with CMB and LSS data mildly prefer $w_0 < -1$, but are well compatible with a cosmological constant. ",
        "introduction": "Current observations \\cite{Astier:2005qq,Riess:2004nr,Spergel:2003cb,Readhead:2004gy,Dickinson:2004yr,Tegmark:2003ud} favor some form of dark energy that today comprises roughly 70\\% of the energy density of our Universe. One fundamental issue is whether dark energy is a true cosmological constant or time evolving \\cite{Wetterich:fm,Ratra:1987rm,Caldwell:1997ii}. In recent years, the notion of an evolving dark energy has been cast in various parameterizations \\cite{Cooray:1999da,Huterer:2000mj,Corasaniti:2002vg,Linder:2002et,Chevallier:2000qy,Upadhye:2004hh,Wang:2004py} of the equation of state $w(a)=p/\\rho$ of dark energy. Yet such parameterizations are ill-suited to catch an intriguing possible feature of dark energy, namely that it could be present at an observable level even from as early as Big Bang Nucleosynthesis on.  Such models would leave their imprints on the Cosmic Microwave Background \\cite{Doran:2000jt}, cosmic structure \\cite{Doran:2001rw,Bartelmann:2005fc} and maybe even Big Bang Nucleosynthesis \\cite{Wetterich:1994bg,Bean:2001wt,Muller:2004gu}. There are some good points in favor of this scenario:  if one links the presently small energy density of dark energy to the age of the Universe, one is led to attractor solutions \\cite{Wetterich:fm,Ratra:1987rm}. Such scenarios occur in attempts to solve the cosmological constant problem from the point of view of dilatation symmetry \\cite{Wetterich:fm} and also in string theories.  Instead of parameterizing $w(a)$, we parameterize $\\od(a)$ directly. This will prove advantageous for two reasons: firstly, the amount of dark energy at early times is then a natural parameter and not inferred by integrating $w(a)$ over the entire evolution. Secondly, since the Hubble parameter is given by \\ee\\label{eqn::hubble} \\frac{H^2(a)}{H^2_0} = \\frac{\\om^0 a^{-3} + \\ore^0 a^{-4}}{1 - \\od(a)}, \\eee a simple, analytic expression for $\\od(a)$ enables us to compute many astrophysical quantities analytically. In the above,  $\\ore^0$ is the fractional energy density of relativistic neutrinos and photons today, $\\om^0$ is the matter (dark and baryonic) fractional energy density and we assumed a flat Universe. For any parameterization of $\\od(a)$, the equation of state can of course be inferred from $\\od(a)$ via an analytic relation (see Equation \\eqref{eqn::connectw}). \\begin{figure} \\includegraphics[scale=0.30] {omegastep_illustration2} \\caption{Evolution of the fractional dark energy density $\\od(z)$ and the equation of state $w(z)$ as a function of redshift. The solid (black) curve depicts the behavior for a $\\Lambda$CDM cosmological constant model, in which $w_0 =-1$ by definition and the amount of dark energy at early times tends to zero. In contrast, the dashed (blue) and dotted (red) curves are early dark energy models described by our parameterization \\protect\\eqref{eqn::param}. For the models depicted, we chose $w_0=-1$ and $\\ode= 0.01$ (blue, dashed), $\\ode=0.07$  (red, dotted) respectively. In addition, we plot the equation of state $w$ for $\\ode=0.01$ (blue, dashed-dotted) and $\\ode=0.07$ (red, dashed-double-dotted) \\label{fig::evolution} } \\end{figure} ",
        "conclusions": "We introduced a new parameterization of dark energy. It is particularly well suited to describe models in which the dark energy density is non-negligible in the early Universe. Working with $\\od(a)$ instead of $w(a)$ facilitates the computation of quantities such as horizons, the luminosity distance and the acoustic scale. Using Markov Chain Monte Carlo simulations, we constrained the parameters of our model using the latest CMB, Sne Ia and LSS data.  While the CMB is more sensitive to $\\ode$, the opposite is true for Sne Ia, which are more sensitive to $w_0$. Adding the recent high multipole Boomerang data tightened the upper bound on $\\ode$ considerably to $\\ode < 0.04$. However, our analysis only included linear fluctuations.  As early dark energy models lead to more non-linear structure at higher redshifts compared to a cosmological constant \\cite{Bartelmann:2005fc}, it might well be that the upper bound turns into a detection in the future.  {\\bf Acknowledgments} We would like to thank Ben Gold for his helpful comments on the preprint."
    },
    "0601/gr-qc0601029_arXiv.txt": {
        "abstract": "In coalescing neutron star binaries, \\textit{r}-modes in one of the stars can be resonantly excited by the gravitomagnetic tidal field of its companion. This post-Newtonian gravitomagnetic driving of these modes dominates over the Newtonian tidal driving previously computed by Ho and Lai. To leading order in the tidal expansion parameter $R$/$r$ (where $R$ is the radius of the neutron star and $r$ is the orbital separation), only the $l=2, |m|= 1$ and $|m| = 2$ \\textit{r}-modes are excited. The tidal work done on the star through this driving has an effect on the evolution of the inspiral and on the phasing of the emitted gravitational wave signal. For a neutron star of mass $M$, radius $R$, spin frequency $f_{\\rm spin}$, modeled as a $\\Gamma =2$ polytrope, with a companion also of mass $M$, the gravitational wave phase shift for the $m=2$ mode is $\\sim 0.1 \\,\\text{radians}\\,(R/10 \\,\\text{km})^4(M/1.4M_\\odot)^{-10/3}(f_{\\rm spin}/100 \\,\\text{Hz})^{2/3}$ for optimal spin orientation. For canonical neutron star parameters this phase shift will likely not be detectable by gravitational wave detectors such as LIGO, but if the neutron star radius is larger it may be detectable if the signal-to-noise ratio is moderately large. The energy transfer is large enough to drive the mode into the nonlinear regime if $f_{\\rm spin} \\agt 100 \\, {\\rm Hz}$. For neutron star - black hole binaries, the effect is smaller; the phase shift scales as companion mass to the -4/3 power for large companion masses. The net energy transfer from the orbit into the star is negative corresponding to a slowing down of the inspiral. This occurs because the interaction reduces the spin of the star, and occurs only for modes which satisfy the Chandrasekhar-Friedman-Schutz instability criterion.  A large portion of the paper is devoted to developing a general formalism to treat mode driving in rotating stars to post-Newtonian order, which may be useful for other applications.  We also correct some conceptual errors in the literature on the use of energy conservation to deduce the effect of the mode driving on the gravitational wave signal. ",
        "introduction": "\\subsection{Background and motivation} \\label{sec:background}     One of the most promising sources for ground-based gravitational wave observatories such as LIGO \\cite{1992Sci...256..325A} and VIRGO \\cite{virgo} are coalescences of compact binary systems \\cite{Cutler:2002me}. The first searches for these systems have already been completed \\cite{Abbott:2003pj}, and future searches will be more sensitive. For neutron star-neutron star (NS-NS) binaries, the most recent estimate of the Galactic merger rate is $83^{+209}_{-66} \\, {\\rm Myr}^{-1}$ \\cite{Kalogera:2003tn,Kalogera:2003tne}; extrapolating this to the distant Universe using the method of Ref.\\ \\cite{Kalogera:2001dz} yields a rate of $28^{+70}_{-22} \\, {\\rm yr}^{-1}$ within a distance of $200 \\, {\\rm Mpc}$. The range of next-generation interferometers in LIGO is expected to be $\\sim 300 \\, {\\rm Mpc}$ \\cite{Cutler:2002me}. Coalescing binaries may well be the first detected sources of gravitational waves.    The gravitational waveforms will carry a variety of different types of information \\cite{Cutler:2002me,1993PhRvL..70.2984C}: First, in the early, low frequency phase of the signal, the binary can be treated to a good approximation as consisting of two spinning point masses.  The corresponding waveforms give detailed information about the deviations of general relativity from Newtonian gravity, and they have been computed to post-3.5-Newtonian order \\cite{Blanchet:2004ek}. Second, toward the end of the inspiral (frequencies of order several hundred Hertz, see Fig.\\ \\ref{fig:noisecurves}) the internal degrees of freedom of the bodies start to appreciably influence the signal, and there is the potential to infer information about the nuclear equation of state.  This has prompted many numerical simulations of the hydrodynamics of NS-NS mergers using various approximations to general relativity (see, eg, Ref.\\ \\cite{Baumgarte:2002jm} and references therein).  Equation of state information can also potentially be extracted from the waves energy spectrum \\cite{2002PhRvL..89w1102F}, and from the NS tidal disruption signal for neutron star-black hole (NS-BH) binaries \\cite{2000PhRvL..84.3519V}.  \\begin{figure} \\begin{center} \\epsfig{file=fig1.eps,width=8.5cm} \\caption{Noise curves $h_{\\rm rms}(f) = \\sqrt{f S_h(f)}$ for initial and next-generation LIGO interferometers \\protect{\\cite{noisecurveref}} (shown in thin lines) together with the signal strength $h_c(f)$ for an inspiralling binary of two $1.4 M_\\odot$ neutron stars at a distance of $200 \\, {\\rm Mpc}$ (thick line) \\protect{\\cite{hcdef}}.  The signal will terminate near the frequency of the innermost stable circular orbit, which is $\\sim 850$ Hz as shown here if the neutron stars have radii of $10$ km and are modeled as $\\Gamma =2$ polytropes \\protect{\\cite{Marronetti:2003hx}}. Resonantly excited modes can give rise to small but potentially measurable corrections to the phase of the waveform in the early portion $10\\,{\\rm Hz} \\lesssim f \\lesssim 100 \\, {\\rm Hz}$ of the signal.} \\label{fig:noisecurves} \\end{center} \\end{figure}   A third type of information potentially carried by the gravitational waves is the effect of the internal structure of the bodies on the early, low frequency ($10\\, \\text{Hz} \\lesssim f \\lesssim 100 \\, \\text{Hz}$) portion of the signal, via the excitation of internal modes of oscillation of one of the neutron stars by the tidal gravitational field of its companion.  In this regime the waveform's phase evolution is dominated by the point-mass dynamics, and the perturbations due to the internal modes are small corrections.  Nevertheless, if the accumulated phase shift due to the perturbations becomes of order unity or larger, it could impede the matched-filtering-based detection of NS-NS signals \\cite{1993PhRvL..70.2984C}. Alternatively the detection of a phase perturbation could give information about the neutron star structure.  For these reasons the excitation of neutron star modes has been studied in detail.   A given mode can be treated as a damped, driven harmonic oscillator, described by the equation \\begin{equation} {\\ddot q} + \\gamma {\\dot q} + \\omega_{\\rm m}^2 q = A(t) \\cos[ m \\phi(t)]. \\label{eq:eom00} \\end{equation} Here $q$ is a generalized coordinate describing the mode, $\\omega_{\\rm m}$ is the mode frequency, $\\gamma$ is a damping constant, $\\phi(t)$ is the orbital phase of the binary, $m$ is the azimuthal quantum number of the mode, and $A(t)$ a slowly varying amplitude. As the inspiral of the binary proceeds, the orbital frequency $\\omega = {\\dot \\phi}$ gradually increases with time.   Sufficiently early in the inspiral, therefore, the mode will be driven below resonance: $m \\omega \\ll \\omega_{\\rm m}$. Because the inspiral is adiabatic we also then have ${\\dot A} / A \\ll \\omega_{\\rm m}$ and ${\\dot \\omega} / \\omega \\ll \\omega_{\\rm m}$.  Using these approximations Eq.\\ (\\ref{eq:eom00}) can be solved analytically in this pre-resonance regime, and one finds that the total energy absorbed by the mode up to time $t$ is \\be E(t) = \\frac{A(t)^2}{2 \\omega_{\\rm m}^2} + \\gamma \\int_{-\\infty}^t dt' \\frac{\\omega(t')^2 A(t')^2 }{ \\omega_{\\rm m}^4/m^2 + \\omega(t')^2 \\gamma^2}. \\label{eq:energyabsorbed} \\ee This energy absorption causes the orbit to inspiral slightly faster and changes the phase of the gravitational wave signal.  Mode excitations can have three different types of effects on the gravitational wave signal:  \\begin{itemize}  \\item {\\it Dissipative:} There is a phase perturbation due to the second term in Eq.\\ (\\ref{eq:energyabsorbed}), corresponding to energy dissipated by the mode.  This is a cumulative effect.  It is dominated by the $l=2, m=\\pm 2$ fundamental modes of the neutron star, and its magnitude depends on the size of the damping constant $\\gamma$, which is determined by the fluid viscosity.  This effect was examined by Bildsten and Cutler \\cite{1992ApJ...400..175B}, who showed that the phase perturbation is negligible for all physically reasonable values of the viscosity.   \\item {\\it Adiabatic:} There is a phase perturbation due to the first term in Eq.\\ (\\ref{eq:energyabsorbed}).  This corresponds to the instantaneous energy present in the mode, as it adjusts adiabatically to sit at the minimum of its slowly-evolving potential.  The perturbation to the waveform is again dominated by the $l=2,m=\\pm 2$ fundamental modes \\cite{1995MNRAS.275..301K}. This effect has been studied analytically in Refs.\\ \\cite{1992ApJ...400..175B,1992ApJ...398..234K,1993ApJ...406L..63L, 1998PhRvD..58h4012T,Mora:2003wt} and numerically in Refs.\\ \\cite{1995MNRAS.275..301K,1994PThPh..91..871S,2001PhRvD..64j4007G,2002PhRvD..65j4021P, 2002PhRvD..66f4013B}.  It can also be inferred from sequences of 3D numerical quasi-equilibrium models of NS-NS binaries \\cite{ 2002PhRvL..89w1102F,Marronetti:2003gk,Bejger:2004zx}. The phase perturbation grows like $\\omega(t)^{5/3}$ and is of order a few cycles by the end of inspiral.  It may be marginally detectable for some nearby events \\cite{Hinderer:2005}.    \\item {\\it Resonant:} The frequencies of internal modes of a neutron star are typically of order $\\omega_{\\rm m} \\sim \\sqrt{M / R^3}$, where $M$ and $R$ are the stellar mass and radius.  (We use units with $G = c = 1$ throughout this paper).  For a NS-NS binary, this is of order the orbital frequency at the end of the inspiral, $\\sim 1000\\, {\\rm Hz}$. Therefore, most modes are driven at frequencies below their natural frequencies throughout the inspiral, and are never resonant.  However, some mode do have frequencies $\\omega_{\\rm m}$ with $\\omega_{\\rm m} \\ll \\sqrt{M/R^3}$, and these modes are resonantly excited as the gradually increasing driving frequency $m \\omega$ sweeps past $\\omega_{\\rm m}$.  During resonance the effect of damping is negligible, and the width $\\Delta \\omega$ of the resonance is determined by the inspiral timescale. The energy deposited in the mode is thus enhanced compared to the equilibrium energy [the first term in Eq.\\ (\\ref{eq:energyabsorbed})] by the ratio of the inspiral timescale to the orbital period.  This can be a significant enhancement.  \\end{itemize}  Resonant driving of modes in NS-NS binaries has been studied in Refs.\\ \\cite{1994ApJ...426..688R,1994MNRAS.270..611L,1999MNRAS.308..153H, 1995MNRAS.275..301K,2001PhRvD..64j4007G, 2002PhRvD..65j4021P,2002PhRvD..66f4013B}, and in the context of other binary systems in Refs.\\ \\cite{Willems:2002na,2003MNRAS.339...25R,Rathore:2004gs,Rathorethesis}. One class of modes with suitably small frequencies are $g$-modes; however the overlap integrals for these modes are so small that the gravitational phase perturbation is small compared to unity \\cite{1995MNRAS.275..301K,1994ApJ...426..688R,1994MNRAS.270..611L}. Another class are the $f$ and $p$-modes of rapidly rotating stars, where the inertial-frame frequency $\\omega_{\\rm m}$ can be much smaller than the corotating-frame frequency $\\sim \\sqrt{M/R^3}$.  Ho and Lai \\cite{1999MNRAS.308..153H} showed that the phase shifts due to these modes could be large compared to unity. However, the required NS spin frequencies are several hundred Hz, which is thought to be unlikely in inspiralling NS-NS binaries.  A third class of modes are Rossby modes ($r$-modes) \\cite{1981A&A....94..126P,1982ApJ...256..717S}, for which the restoring force is dominated by the Coriolis force.  For these modes the mode frequency $\\omega_{\\rm m}$ is of order the spin frequency of the star, and thus can be suitably small [$10 \\, \\text{Hz} \\lesssim \\omega_\\text{ m} / (2 \\pi) \\lesssim 100 \\, \\text{Hz}$].  Ho and Lai \\cite{1999MNRAS.308..153H} computed the Newtonian driving of these modes, and showed that the phase shift is small compared to unity.  In this paper we compute the phase shift due to the post-Newtonian, gravitomagnetic resonant driving of the $r$-modes.  Normally, post-Newtonian effects are much smaller than Newtonian ones.  However, in this case the post-Newtonian effect is larger.  The reason is that the Newtonian driving acts via a coupling to the mass quadrupole moment \\begin{equation} Q_{ij} = \\int d^3 x \\rho (x_i x_j - \\frac{1}{3} \\delta_{ij} {\\bf x}^2) \\end{equation} of the mode, where $\\rho$ is the mass density. For $r$ modes, this mass quadrupole moment vanishes to zeroth order in the angular velocity $\\Omega$ of the star; it is suppressed by a factor $\\sim \\Omega^2 R^3 / M$. By contrast, the gravitomagnetic driving acts via a coupling to the current quadrupole moment \\begin{equation} S_{ij} = \\frac{1}{2} \\int d^3 x \\rho \\left[ ({\\bf x} \\times {\\bf v})_i x_j + ({\\bf x} \\times {\\bf v})_j x_i \\right], \\end{equation} where ${\\bf v}$ is the velocity of the fluid. This current quadrupole moment is nonvanishing to zeroth order in $\\Omega$, and the gravitomagnetic coupling is therefore enhanced over the Newtonian tidal coupling by a factor of $\\sim \\Omega^2 R^3/M$. This enhancement factor can exceed the usual factor $\\sim v^2/c^2$ by which post-Newtonian effects are suppressed, where $v$ is a velocity scale.  It is for this same reason that the radiation-reaction induced instability of $r$-modes in newly born neutron stars \\cite{1998ApJ...502..708A,1998ApJ...502..714F} is dominated by the gravitomagnetic coupling and not the mass quadrupole coupling.     \\subsection{Effect of resonance on gravitational wave signal}  We now discuss the observational signature of a mode resonance in the gravitational wave signal.  We denote by $\\Phi(t)$ the phase of the waveform, which is twice the orbital phase $\\phi(t)$.  We denote by $\\Phi_{\\rm pp}(t)$ the phase that one obtains from a point particle model, neglecting the effects of the internal modes of the stars. To the leading post-Newtonian order, this point-particle phase is given from the quadrupole formula by the differential equation \\cite{Peters:1963ux} \\begin{equation} {\\dot \\omega} = \\frac{96}{5} \\mu M_{\\rm t}^{2/3} \\omega^{11/3}, \\label{eq:inspiral0} \\end{equation} where $\\omega = {\\dot \\phi}$ is the orbital angular velocity, and $M_{\\rm t}$ and $\\mu$ are the total mass and reduced mass of the binary.  Consider now the phase $\\Phi(t)$ including the effect of the mode. We neglect the phase shift caused by the adiabatic response of the mode at early times before the resonance [the first term in Eq.\\ (\\ref{eq:energyabsorbed})], as this phase shift is small compared to the effect of the resonance.  We also neglect the cumulative phase shift due to damping.  Therefore at early times we have $\\Phi(t) = \\Phi_{\\rm pp}(t)$.  The duration of the resonance is short compared to the inspiral timescale (see Sec.\\ \\ref{sec:estimates} below).  After the resonance, the mode again has a negligible influence on the orbit of the binary, and Eq.\\ (\\ref{eq:inspiral0}) applies once more.  However, the two constants of integration that arise when solving Eq.\\ (\\ref{eq:inspiral0}) need not match between the pre-resonance and post-resonance waveforms.  Therefore, the phase can be written as\\footnote{The sign of the $\\Delta \\Phi$ term in Eq.\\ (\\protect{\\ref{eq:signature}}) is chosen for later convenience.} \\begin{equation} \\Phi(t) = \\left\\{ \\begin{array}{ll} \\Phi_{\\rm pp}(t) & \\mbox{ $t - t_0 \\ll - t_{\\rm res}$,}\\\\ \\Phi_{\\rm pp}(t + \\Delta t) - \\Delta \\Phi & \\mbox{ $t-t_0 \\gg t_{\\rm res}$.}\\\\ \\end{array} \\right. \\label{eq:signature} \\end{equation} Here $t_0$ is the time at which resonance occurs, and $t_{\\rm res}$ is the duration of the resonance.\\footnote{At intermediate times $|t-t_0| \\sim t_{\\rm res}$, the phase $\\Phi(t)$ is not related in a simple way to $\\Phi_{\\rm pp}(t)$.  However, since the duration of the resonance is typically a few tens of cycles (short compared to the $\\sim 10^4$ cycles of inspiral in the detector's frequency band), the phase perturbation in this intermediate stage will not be observable.}  The effect of the resonance is thus to cause an overall phase shift $\\Delta \\Phi$, and also to cause a time shift $\\Delta t$ in the signal.  If energy is absorbed by the mode, then the inspiral proceeds more quickly and $\\Delta t$ is positive.    It turns out that the two parameters $\\Delta t$ and $\\Delta \\Phi$ characterizing the effect of the resonance are not independent. To a good approximation they are related by \\begin{equation} \\Delta \\Phi = {\\dot \\Phi}_{\\rm pp}(t_0) \\Delta t, \\label{eq:constraint} \\end{equation} as argued by Reisenegger and Goldreich \\cite{1994ApJ...426..688R}, and as derived in detail in Sec.\\ \\ref{sec:phasing} below. Inserting this relation into Eq.\\ (\\ref{eq:signature}) and expanding to linear order in the small parameter $\\Delta t$ yields for the late-time phase \\begin{equation} \\Phi(t) = \\Phi_{\\rm pp}(t) + \\left[ {\\dot \\Phi}_{\\rm pp}(t) - {\\dot \\Phi}_{\\rm pp}(t_0) \\right] \\Delta t. \\label{eq:latetimephase} \\end{equation} The physical meaning of the condition (\\ref{eq:constraint}) is thus that the resonance can be idealized as an instantaneous change in frequency at $t=t_0$ with no corresponding instantaneous change in phase. Rewriting the phase (\\ref{eq:latetimephase}) in terms of $\\Delta \\Phi$ using (\\ref{eq:constraint}) and using ${\\dot \\Phi}_{\\rm pp} = 2 \\omega$ gives \\begin{equation} \\Phi(t) = \\Phi_{\\rm pp}(t) + \\left[ \\frac{\\omega(t)}{\\omega_0} -1 \\right] \\Delta \\Phi, \\label{eq:signature1} \\end{equation} where $\\omega_0 = \\omega(t_0)$ is the orbital frequency at resonance. The phase perturbation due to the resonance therefore grows with time after the resonance, and can become much larger than $\\Delta \\Phi$ when $\\omega \\gg \\omega_0$.     One can alternatively think of the phase correction as being non-zero before the resonance, and zero afterwards \\cite{1994ApJ...426..688R}. This is what would be perceived if one matched a template with the portion of the observed waveform after the resonance.  In other words, if we define a point-particle waveform phase ${\\tilde \\Phi}_{\\rm pp}(t)$ with different conventions for the starting time and starting phase via \\begin{equation} {\\tilde \\Phi}_{\\rm pp}(t) \\equiv \\Phi_{\\rm pp}(t + \\Delta t) - \\Delta \\Phi, \\end{equation} then we obtain \\begin{equation} \\Phi(t) = \\left\\{ \\begin{array}{ll} {\\tilde \\Phi}_{\\rm pp}(t) - \\left[ \\frac{\\omega(t)}{\\omega_0} -1 \\right] \\Delta \\Phi & \\mbox{\\ \\  $t - t_0 \\ll - t_{\\rm res}$,}\\\\ {\\tilde \\Phi}_{\\rm pp}(t) & \\mbox{\\ \\  $t-t_0 \\gg t_{\\rm res}$.}\\\\ \\end{array} \\right. \\label{eq:signature2} \\end{equation} The phase correction $\\Phi - {\\tilde \\Phi}_{\\rm pp}$ is now a linear function of frequency, interpolating between $\\Delta \\Phi$ at $\\omega =0$ to zero at resonance $\\omega = \\omega_0$.  Since this phase correction never exceeds $\\Delta \\Phi$, we see that the detectability criterion for the resonance effect is $\\Delta \\Phi \\gtrsim 1$.  Thus, the phase shifts $\\gg \\Delta \\Phi$ predicted by Eq.\\ (\\ref{eq:signature1}) are fictitious in the sense that they are not measurable.   \\subsection{Order of magnitude estimates} \\label{sec:estimates}  We now turn to an order of magnitude estimate of the mode excitation and the phase shift $\\Delta \\Phi$, both for the Newtonian driving studied previously and for the gravitomagnetic driving.   We consider a star of mass $M$ and radius $R$ with a companion of mass $M'$. The orbital angular velocity at resonance is $\\omega_0$; this is also the mode frequency (up to a factor of the azimuthal quantum number $m$ which we neglect here).  The separation $r$ of the two stars at resonance is then given by $\\omega_0^2 \\sim M_{\\rm t}/r^3$, where $M_{\\rm t} = M + M'$ is the total mass.   Let the magnitude of the tidal acceleration be $a_{\\rm tid}$, and let the duration of the resonance be $t_{\\rm res}$.  The mode can be treated as a harmonic oscillator of mass $\\sim M$ and frequency $\\omega_0$.  During the resonance, the mode absorbs energy at the same rate as a free particle would, so the total energy absorbed is \\be \\Delta E \\sim M a_{\\rm tid}^2 t_{\\rm res}^2. \\label{eq:DeltaE} \\ee Now the gravitational wave luminosity ${\\dot E}_{\\rm gw}$ is to a good approximation unaffected by the resonance.  Therefore the loss of energy $\\Delta E$ from the orbit decreases the time taken to inspiral by an amount $\\Delta t = \\Delta E / {\\dot E}_{\\rm gw}$. From Eq.\\ (\\ref{eq:constraint}) the corresponding phase shift is $\\Delta \\Phi \\sim \\Delta t/t_{\\rm orb}$, where $t_{\\rm orb} \\sim 1/\\omega_0$ is the orbital period.  This gives \\be \\Delta \\Phi \\sim \\frac{t_{\\rm rr} \\Delta E } {t_{\\rm orb} E_{\\rm orb}}, \\label{eq:DeltaPhi11} \\ee where $E_{\\rm orb}$ is the orbital energy and $t_{\\rm rr}$ is the radiation reaction timescale.  This phase shift is just the energy absorbed divided by the energy radiated per orbit.   Next, using $E_{\\rm orb} \\sim M M'/r$ and the formula (\\ref{eq:DeltaE}) for $\\Delta E$ gives \\be \\Delta \\Phi \\sim \\frac{r}{M'} a_{\\rm tid}^2 \\frac{t_{\\rm res}^2 t_{\\rm rr} }{ t_{\\rm orb}}. \\ee The resonance time $t_{\\rm res}$ is the geometric mean of the orbital and radiation reaction times $t_{\\rm orb}$ and $t_{\\rm rr}$.  This follows from the fact that the orbital phase near resonance can be expanded as \\be \\phi \\sim \\phi_0 + \\omega_0 (t - t_0) + \\frac{\\omega_0}{t_{\\rm rr}} (t-t_0)^2 + \\label{eq:phaseexpand} \\ldots, \\ee from the definition of the radiation reaction timescale $t_{\\rm rr}$. The mode is resonant when the quadratic term in Eq.\\ (\\ref{eq:phaseexpand}) is small compared to unity, so that the force on the mode is in phase with the mode's natural oscillation. This gives $t_{\\rm res} \\sim \\sqrt{t_{\\rm orb} t_{\\rm rr}}$, and we get \\be \\Delta \\Phi \\sim \\frac{r}{M'} a_{\\rm tid}^2 t_{\\rm rr}^2. \\ee Finally, we use the scaling $t_{\\rm rr} \\sim r^4 / (\\mu M_{\\rm t}^2)$ for the radiation reaction time given by the quadrupole formula, where $\\mu$ is the reduced mass. This gives \\be \\Delta \\Phi \\sim \\frac{r^9 }{ M' \\mu^2 M_{\\rm t}^4} a_{\\rm tid}^2. \\label{eq:phaseans0} \\ee   We now consider various different cases.  For the Newtonian driving of a mode via a tidal field of multipole of order $l$, we have \\be a_{\\rm tid} \\sim \\frac{M' }{ r^2} \\left({\\frac{R}{r}} \\right)^{l-1}. \\label{eq:atidal} \\ee [The leading order, quadrupolar driving is $l=2$.]  Inserting this into Eq.\\ (\\ref{eq:phaseans0}) and eliminating $r$ using $\\omega_0^2 r^3 \\sim M_{\\rm t}$ gives \\be \\Delta \\Phi \\sim \\left( \\frac{ M^3 }{\\mu M_{\\rm t}^2} \\right) \\left( \\frac{R}{M} \\right)^5 \\left( \\frac{\\omega_0^2}{M/R^3} \\right)^{(l-5)/3}, \\label{eq:phasefinal} \\ee which agrees with previous analyses \\cite{1994ApJ...426..688R,1994MNRAS.270..611L} for $g$ and $f$ modes.  For the Newtonian driving of $r$-modes, the leading order driving is $l=3$ rather than $l=2$ \\cite{1999MNRAS.308..153H}.  Also the tidal acceleration $a_{\\rm tid}$ should be multiplied by the suppression factor $\\omega_0^2 R^3/M$ discussed in Sec.\\ \\ref{sec:background} above, and the final phase shift (\\ref{eq:phasefinal}) should be multiplied by the square of this factor.  This gives in the equal mass case $M = M'$ \\be \\Delta \\Phi \\sim \\frac{R^{10} \\omega_0^{10/3}}{M^{20/3}}. \\ee Now the frequency $m \\omega_0$ of the dominant, $l=m=2$ $r$-mode is related to the spin frequency $f_{\\rm spin}$ of the neutron star by $2 \\omega_0 = 8 \\pi f_{\\rm spin}/3$, so we obtain \\be \\Delta \\Phi = k \\frac{R^{10} f_{\\rm spin}^{10/3}}{M^{20/3}}, \\label{eq:rmodeNestimate} \\ee where $k$ is a dimensionless constant of order unity.  Equation (\\ref{eq:rmodeNestimate}) agrees with the results of Ho and Lai \\cite{1999MNRAS.308..153H}, who show that $k \\approx 0.4$ for a $\\Gamma=2$ polytrope, and for the $l=m=2$ $r$-mode.\\footnote{See equation (4.13) of Ref.\\ \\cite{1999MNRAS.308..153H}, specialized to the equal mass case and maximized over the angle $\\beta$ between the spin of the star and the orbital angular momentum.} This gives \\bea && \\Delta \\Phi_{\\rm r-mode,N} = 8 \\times 10^{-5} R_{10}^{10} f_{s100}^{10/3} M_{1.4}^{-20/3}, \\label{eq:HoLai} \\eea where the subscript N denotes Newtonian and we have defined $R = 10 R_{10} \\, {\\rm km}$, $M = 1.4 M_{1.4} M_\\odot$, and $f_{\\rm spin} = 100 f_{s100} \\, {\\rm Hz}$.   For the gravitomagnetic driving of $r$-modes, the tidal acceleration is given by \\be a_{\\rm tid} \\sim \\left( \\frac{M' R}{r^3} \\right) v_{\\rm spin} v_{\\rm orb}. \\ee Here the first factor in brackets is tidal acceleration (\\ref{eq:atidal}) for Newtonian quadrupolar $l=2$ driving.  Since the interaction is gravitomagnetic, it is suppressed by two powers of velocity relative to the Newtonian interaction.  One power of velocity is associated with the source of the gravitomagnetic field, and is the orbital velocity $v_{\\rm orb} \\sim \\omega_0 r$.  The second power of velocity is that coming from the Lorentz-type ${\\bf v} \\times {\\bf B}$ force law, and is the internal fluid velocity in the neutron star due to its spin, $v_{\\rm spin} \\sim \\omega_0 R$.  Combining these estimates with Eq.\\ (\\ref{eq:phaseans0}) and eliminating $\\omega_0$ in favor of $f_{\\rm spin}$ gives \\be \\Delta \\Phi \\sim \\frac{R^4 f_{\\rm spin}^{2/3}}{M' M^2 M_{\\rm t}^{1/3}}. \\label{eq:ans12} \\ee In the equal mass case $M = M'$ this gives \\bea && \\Delta \\Phi_{\\rm r-mode,PN} \\sim 3.4 \\ R_{10}^{4} f_{s100}^{2/3} M_{1.4}^{-10/3}. \\label{eq:ans2} \\eea Here the subscript PN denotes post-Newtonian.             \\subsection{Results and implications}  \\begin{figure} \\begin{center} \\epsfig{file=fig2.eps, angle=0, width=8.5cm} \\caption{The magnitude of the gravitational-wave phase shift $\\Delta \\Phi$ due to the $r$-mode resonance as a function of the spin frequency $f_{\\rm spin}$ of the neutron star, for both the Newtonian and gravitomagnetic driving, for a binary with masses $M = M' = 1.4 M_\\odot$, and for a neutron star radius of $R = 10 \\, {\\rm km}$.  The star is modeled as a $\\Gamma=2$ polytrope. The mode plotted is $l=2$, $m=2$ in each case.  The angle $\\psi$ between the spin and the orbital angular momentum has been chosen in each case to maximize the phase shift; $\\psi$ is $34.4^\\circ$ in the Newtonian case and $60^\\circ$ in the post-Newtonian case. The resonance occurs at an orbital frequency of $4 f_{\\rm spin} / 9$ in the Newtonian case and at $4 f_{\\rm spin}/3$ in the post-Newtonian case. The corresponding gravitational wave frequencies are $8 f_{\\rm spin} / 9$ and $8 f_{\\rm spin}/3$ respectively. It can be seen that the gravitomagnetic driving dominates.} \\label{fig:result} \\end{center} \\end{figure}   The fact that the order-of-magnitude estimate (\\ref{eq:ans2}) of the phase shift is of order unity, and thus potentially detectable in the gravitational wave signal via matched filtering \\cite{1993PhRvL..70.2984C}, is the primary motivation for the computation of this paper.  We compute the dimensionless coefficient in Eq.\\ (\\ref{eq:ans12}), and find for the case of a $\\Gamma=2$ polytrope for the $l=2$, $m=2$ mode [cf.\\ Eq.\\ (\\ref{eq:ans2ssa}) below, multiplied by two to convert from orbital phase shift $\\Delta \\phi$ to gravitational wave phase shift $\\Delta \\Phi$] \\be \\Delta \\Phi_{\\rm r-mode,PN} = -0.06 \\ R_{10}^4 f_{s100}^{2/3} M_{1.4}^{-10/3}. \\label{eq:ans2ss} \\ee This is somewhat smaller than the estimate (\\ref{eq:ans2}). Figure \\ref{fig:result} compares this result to the Newtonian phase shift (\\ref{eq:HoLai}). In Table \\ref{table:1} we show the phase shift $\\Delta \\Phi_{22} = 2 \\Delta \\phi_{22}$ that will occur at coalescence for the five known double neutron star binaries.\\footnote{The orbital motion of these five systems is eccentric today, but radiation reaction will have circularized the orbits by the time resonance occurs. It is thus consistent to use our circular-orbit results for these binaries.} For B1913+16 we used the known value $\\psi=22^\\circ$ of the angle between the pulsar spin and the orbital angular momentum \\cite{1998ApJ...509..856K}, and for the remaining binaries for which $\\psi$ is unknown we took an average of the angular factors appearing in the phase lag (\\ref{eq:deltaphi22ans}).  We assumed a $\\Gamma =2$ polytrope equation of state and a radius of 10 km for all pulsars.     \\begin{table} \\begin{tabular}{|c|c|c|c|c|} \\hline System & spin period& $M$ & $M^\\prime$ & phase lag $\\Delta \\Phi_{22}$ \\\\ \\hline B1913+16 &59.0&1.441&1.387& -0.005 \\\\ B1534+12 &37.9&1.339&1.339& -0.009 \\\\ B2127+11C &30.5&1.358&1.354& -0.010 \\\\ J0737-3039 & 22.7&1.34&1.25& -0.013\\\\ J1756-2251 & 28.5&1.40&1.18& -0.011\\\\ \\hline \\end{tabular} \\caption{List of gravitational wave phase lag parameters that will occur for the $l=2,m=2$ resonance for the known double neutron star binaries when they coalesce. For B1913+16, the known value for the misalignment angle $\\psi$ ($22^\\circ$) was used; for the others, the average value of the angular factors was used. The pulsar mass $M$ and companion mass $M^\\prime$ are in solar masses and the pulsar spin period is in milliseconds. The parameter values were taken from Ref.\\ \\cite{2005LRR.....8....7L}.} \\label{table:1} \\end{table}   The phase shift (\\ref{eq:ans2ss}) is sufficiently small that it is unlikely to be detectable in the gravitational wave signal for most detected binary inspirals.  Even for the most favorable conceivable neutron star parameters of $M = 1.0 \\, M_\\odot$ and $R = 15 \\, {\\rm km}$, the phase shift is $\\sim 0.8$ radians.  An analysis of the detectability of mode resonances indicates that the minimum detectable value of $\\Delta \\Phi$ is $\\Delta \\Phi \\sim 2$ for $f_{\\rm spin} = 100 \\, {\\rm Hz}$ at a signal to noise ratio of $S/N = 10$ \\cite{Balachandran:2005}. Therefore the minimum signal to noise ratio necessary for detection of the effect is between $\\sim 30$ and $\\sim 300$, depending on the neutron star parameters.  It is conceivable that the effect could be detected in some of the closest detected binary inspirals with advanced interferometers.     One of the interesting features of the interaction is that the phase shift is negative, corresponding to energy being transferred from the star to the orbit rather than the other way around.  This arises because the star is rotating and the interaction reduces the spin of the star, tapping into its rotational kinetic energy.\\footnote{This follows from the fact that at fixed angular momentum, the energy of a star is minimized at uniform rotation.  Therefore the only way the total energy of a rotating star can be reduced is via a reduction of the spin.}  In Sec.\\ \\ref{sec:evol2} below we show that this sign of energy transfer is only possible for modes which satisfy the Chandrasekhar-Friedman-Schutz mode instability criterion \\cite{1970PhRvL..24..762C,1978ApJ...222..281F} of being retrograde in the corotating frame and prograde in the inertial frame.  Although the $r$-mode resonance may not be detectable, it is nevertheless probably the strongest of all the mode resonances in the early phase of the gravitational wave signal.  The energy transfer corresponding to the phase shift (\\ref{eq:ans2ss}) is from Eq.\\ (\\ref{eq:DeltaPhi11}) \\be \\Delta E \\sim (1.4 \\times 10^{48} \\, {\\rm erg})\\, R_{10}^4 f_{s100}^3, \\ee which is fairly large.  The corresponding dimensionless amplitude $\\alpha$ of the $r$-mode in the notation of Ref.\\ \\cite{prl...80...4843} is \\be \\alpha \\sim 0.3 \\, R_{10}^2 f_{s100} M_{1.4}^{-1}. \\ee This amplitude is sufficiently large that nonlinear mode-mode interactions should be important, as in the case of unstable $r$-modes in newly born neutron stars and in low mass X-ray binaries \\cite{prl...80...4843}.  In those contexts the mode-mode coupling causes the mode growth to saturate at $\\alpha \\sim 10^{-4}$ \\cite{Schenk:2001zm,Morsink:2002ut,Arras:2002dw,Brink:2004bg,Brink:2004qf,Brink:2004kt}. Here however the $r$-mode is driven on much shorter timescales, and the mode-mode coupling likely does not have time to operate efficiently. It is possible that the type of nonlinear effects studied numerically in Refs.\\ \\cite{Gressman:2002zy,Lindblom:2001sg} (generation of differential rotation, shocks, wave breaking) could occur.     \\subsection{Organization of this paper}  The remainder of this paper is organized as follows.  In Secs.\\ \\ref{sec:prelim}, \\ref{sec:modeevolution} and \\ref{sec:phasing} we analyze resonant driving of modes in Newtonian gravity, building on the previous studies of Refs.\\ \\cite{1994ApJ...426..688R,1994MNRAS.270..611L,1999MNRAS.308..153H,Rathore:2004gs,Rathorethesis}. The time delay parameter $\\Delta t$ can be derived using energy conservation which we discuss in Appendix \\ref{sec:energycons}, but the relation (\\ref{eq:constraint}) between the phase shift parameter $\\Delta \\Phi$ and $\\Delta t$ requires an analysis of the orbital equations of motion, which we give in Sec.\\ \\ref{sec:phasing}. In Sec.\\ \\ref{sec:rmodedriving} and Appendix \\ref{app:Y} we show that the gravitomagnetic driving of modes can be computed to leading order using Newtonian stellar perturbation theory supplemented by a gravitomagnetic external acceleration, and we deduce the phase shifts generated by the driving of $r$-modes. ",
        "conclusions": "In this paper we have studied in detail the gravitomagnetic resonant tidal excitation of {\\it r}-modes in coalescing neutron star binaries.  We first analyzed the simpler case of Newtonian resonant tidal driving of normal modes. We showed, by integrating the equations of motion of the mode-orbit system, that the effect of the resonance to leading order is an instantaneous shift in the orbital frequency of the binary at resonance. The effect of the modes on the orbital motion away from resonance is a small correction to this leading order effect. We then showed that for the case of gravitomagnetic tidal driving of {\\it r}-modes, the equations of motion for the mode-orbit system can be manipulated into a form analogous to the Newtonian case. We then made use of the Newtonian results to compute the instantaneous change in frequency due to the driving of the {\\it r}-modes. This shift in frequency is negative, corresponding to energy being removed from the star and added to the orbital motion.  This sign of energy transfer occurs only for modes which satisfy the Chandrasekhar-Friedman-Schutz instability criterion. We also argued that the resonances are at best marginally detectable with LIGO."
    },
    "0601/astro-ph0601634_arXiv.txt": {
        "abstract": "We present the spatial correlation function analysis of non-stellar X-ray point sources in the {\\it Chandra} Large Area Synoptic X-ray Survey of Lockman Hole Northwest (CLASXS).  Our 9 ACIS-I fields cover a contiguous solid angle of 0.4 deg$^{2}$ and reach a depth of $3 \\times 10^{-15}$~\\ergpcmsqps in the 2--8 keV band. We supplement our analysis with data from the Chandra Deep Field North (CDFN). The addition of this field allows better probe of the correlation function at small scales. A total of 233 and 252 sources with spectroscopic information are used in the study of the CLASXS and CDFN fields respectively.  We calculate both redshift-space and projected correlation functions in comoving coordinates, averaged over the redshift range of $0.1<z<3.0$, for both CLASXS and CDFN fields for a standard cosmology with $\\Omega_{\\Lambda} = 0.73, \\Omega_{M} = 0.27$, and $h = 0.71$ ($H_0 = 100h$~km~s$^{-1}$~Mpc$^{-1}$). The correlation function for the CLASXS field over scales of 3~Mpc$<s<$~200~Mpc can be modeled as a power-law of the form $\\xi(s) = (s/s_0)^{-\\gamma}$, with $\\gamma = 1.6^{+0.4}_{-0.3}$ and $s_0 = 8.0^{+1.4}_{-1.5}$~Mpc. The redshift-space correlation function for CDFN on scales of 1~Mpc$<s<$~100~Mpc is found to have a similar correlation length $s_0 = 8.55^{+0.75}_{-0.74}$~Mpc, but a shallower slope ($\\gamma = 1.3 \\pm 0.1$). The real-space correlation functions derived from the projected correlation functions, are found to be $r_0 = 8.1^{+1.2}_{-2.2}$~Mpc, and $\\gamma = 2.1 \\pm 0.5$ for the CLASXS field, and $r_0 = 5.8^{+1.0}_{-1.5}$~Mpc, $\\gamma = 1.38^{+0.12}_{-0.14}$ for the CDFN field. By comparing the real- and redshift-space correlation functions in the combined CLASXS and CDFN samples, we are able to estimate the redshift distortion parameter $\\beta = 0.4 \\pm 0.2$ at an effective redshift $z = 0.94$. We compare the correlation functions for hard and soft spectra sources in the CLASXS field and find no significant difference between the two groups. We have also found that the correlation between  X-ray luminosity and clustering amplitude is weak, which, however, is fully consistent with the expectation using the simplest relations between X-ray luminosity, blackhole mass, and dark halo mass.  We study the evolution of the AGN clustering by dividing the samples into 4 redshift bins over 0.1~Mpc$<z<$3.0~Mpc. We find a very mild evolution in the clustering amplitude, which show the same evolution trend found in optically selected quasars in the 2dF survey.  We estimate the evolution of the bias, and find that the bias increases rapidly with redshift ($b(z=0.45) = 0.95 \\pm 0.15$and $b(z = 2.07)= 3.03 \\pm 0.83$). The typical mass of the dark matter halo derived from the bias estimates show little change with redshift. The average  halo mass is found to be  $\\log~(M_{halo}/M_{\\sun}) \\sim 12.1$. ",
        "introduction": "Structure formation and evolution in the universe and the formation and growth of supermassive black holes (SMBHs) are two fundamental problems in astronomy which are still not fully understood. While recent progresses in the cosmic microwave background, the high redshift type Ia supernovae survey, and the large optical surveys have significantly improved our understanding of the evolution of large scale structure, there are still several gaps in the picture of structure formation. The data at redshift of $\\sim 1$, where most of the cosmic star formation might have taken place, is still very limited. On scales of galaxies and cluster of galaxies, the feed back process from galaxies or AGNs could significantly alter structure formation models where gravitation is the only driving force. The clustering of active galactic nuclei (AGNs) provides unique path to the solution of these problems because (1) the AGNs are often bright compared to normal galaxies and are easily seen at large cosmological distance; (2) AGNs trace the violent growth phase of SMBHs and hence their clustering properties provide a link between the dark matter halo to the AGN activity. Large scale AGN surveys have been traditionally carried out in the optical band with dedicated telescopes. The most recent of these are the Sloan Digital Sky Survey (SDSS, {Schneider} {et~al.} 2004) and the Two Degree Field Survey (2dF, {Croom} {et~al.} 2005, C05 hereafter).  These surveys have demonstrated that the clustering of AGNs can be used to measure cosmological parameters ({Croom} {et~al.} 2004; {Outram} {et~al.} 2004), and to constrain gravitational lensing ({Myers} {et~al.} 2003).  However, as has been found in recent {\\it Chandra} and {\\it XMM-Newton} deep surveys, a large fraction of X-ray detected AGNs show little or no activity in optical observations (e.g. {Barger} {et~al.} 2005), most probably due to obscuration ({Fabian} \\& {Iwasawa} 1999) but possibly with some contribution from light dilution of the host galaxy ({Moran}, {Filippenko}, \\& {Chornock} 2002), though Barger et al. (2005) argue this is not a dominant factor. This results in a large fraction of AGNs being missed in the optical surveys. On the other hand, hard X-rays ($>2$~keV) is almost unaffected by obscuring column densities $N_{H}< 10^{24}$~cm$^{-2}$, and the X-ray emission from the host galaxies is low compared to the AGNs. Thus hard X-rays are at present the best energy band to find AGNs ({Mushotzky} 2004).  Recent optical follow-ups of {\\it Chandra} deep surveys have revealed that the hard  X-ray sources are mostly found around $z \\sim 1$ instead of $z \\sim 2$ as seen in optically selected quasar samples. The low redshift population is dominated by non-broadline objects, while broadline AGNs are found mostly at higher redshifts ({Steffen} {et~al.} 2003). Given these new discoveries, it is important to know how the clustering properties of X-ray and optical selected AGN differ.  The most extensive X-ray AGN surveys so far performed used the {\\it ROSAT} telescope ({Mullis} {et~al.} 2004). Because the telescope is not sensitive above 2 keV, {\\it ROSAT} misses a large fraction of hard X-ray sources . The relatively poor spatial resolution of ROSAT also limits the accuracy of optical identifications. Most of the {\\it ROSAT} detected AGNs show broad emission lines in their optical spectra.  Both the optical quasar surveys and the ROSAT surveys suggest that AGNs have correlation properties similar to the local galaxies. The results seem to be independant of the sample medium redshifts. This result is puzzling because AGNs are believed to be preferentially form in high density peaks where interactions between galaxies are more common, and interactions in turn are thought to be crucial in AGN fueling ({Di Matteo}, {Springel}, \\&  {Hernquist} 2005). The mass of the dark matter halos that host AGNs are hence likely to be more massive.  The clustering results on hard X-ray AGNs are so far contradictory. Earlier studies of a small number of individual {\\it Chandra} fields seem to indicate that the hard band number counts in these small fields has fluctuations larger than expected from Poisson noise ({Cowie} {et~al.} 2002; {Manners} {et~al.} 2003) but the result is contradicted with larger samples of {\\it Chandra} fields ({Kim} {et~al.} 2004). {Basilakos} {et~al.} (2004) found a 4$\\sigma$ clustering signal in hard X-ray sources at $f_{2-8 keV} > 10^{-14}$~\\ergpcmsqps using angular correlation functions on a XMM detected AGN sample from a 2 deg$^{2}$ survey. A similar result was also found earlier in our 0.4~deg$^{2}$ {\\it Chandra} field (see below) using the count-in-cells technique ({Yang} {et~al.} 2003). Using the Limber equation Basilacos et al. (2004) argue that the hard X-ray sources are likely to be more strongly clustered than the optically selected AGNs. {Gilli} {et~al.} (2003) reported the detection of large angular-redshift clustering in the Chandra Deep field South, which seems to be dominated by hard X-ray sources. Using the projected correlation function for the optically identified X-ray sources from the Chandra Deep field North (CDFN) and South (CDFS), {Gilli} {et~al.} (2005) found that the average correlation amplitude in the CDFS is higher than that in the CDFN, and the latter is consistent with the correlation amplitude found in optically detected quasars.  In this paper, we report our spatial correlation function analysis of the optically identified X-ray sources in the {\\it Chandra} Large Area Synoptic X-ray Survey of the ISO Lockman Hole Northwest region (CLASXS). CLASXS is so far the largest contiguous {\\it Chandra} deep field with a high level of spectroscopic identifications. The size of the field is chosen to reduce the cosmic variance in the X-ray background to $\\sim 10\\%$ (Yang et al. 2004). For comparison, we have also analyzed the correlation functions for the CDFN field, using the published X-ray catalog by {Alexander} {et~al.} (2003) and the optical catalog of {Barger} {et~al.} (2003). Because the two surveys use basically the same optical instruments in the follow-up observations, and thus have the same accuracy in redshift measurements, the comparison is relatively straight forward. The LogN-LogS of the CDFN agrees well with that of CLASXS, also indicating that the CDFN is a ``typical'' field. The depth of the CDFN is very useful in probing the correlation function at small separations. In \\S~\\ref{sec_obs_data} we summarize our observations and data analysis. In \\S~\\ref{sec_methods} we discuss the methodology we use in the clustering analysis. The results of the correlation functions are presented in \\S~\\ref{sec_results}. The evolution of AGN clustering is presented in \\S~\\ref{sec_evolution}. In \\S~\\ref{sec_discussion} we discussion the implications of our results. Finally summarize our results in \\S~\\ref{sec_conclusion}. Throughout this paper, unless noted otherwise, we assume $H_{0}=71$ and a flat universe with $\\Omega_{M}=0.27$ and $\\Omega_{\\Lambda} = 0.73$. ",
        "conclusions": "} In this paper we study the clustering of the X-ray selected AGNs in the 0.4~deg$^{2}$ {\\it Chandra} contiguous survey of the Lockman Hole Northwest region. Based on our previous study, the size of the CLASXS field is large enough that the cosmic variance should not be important. We supplement our study with the published data from the CDFN. The very similar LogN-LogS of the CLASXS and CDFN suggests that the cosmic variance should not be important when the CDFN is included in the analysis. The very deep CDFN gives a better probe of the correlation function at small separations. A total of 233 and 252 non-stellar sources from CLASXS and CDFN respectively are used in this study. We use the correlation function in the redshift-space as a major tool in our clustering analysis. For the whole sample, we have also performed an analysis using the projected correlation function. This allows us to quantify the effect of redshift distortion.  We summarize our results as follows:  1. We calculated the redshift-space correlation function for sources with $0.1<z<3.0$ in both the CLASXS and CDFN fields, assuming constant clustering in comoving coordinates. We found a $6.7\\sigma$ clustering signal for pairs within $s<20$~Mpc in the CLASXS field. The correlation function over scale of 3~Mpc$<s<$~200~Mpc is found to be a power-law with $\\gamma = 1.6^{+0.4}_{-0.3}$ and $s_0 = 8.0^{+1.4}_{-1.5}$~Mpc. The redshift-space correlation function for CDFN on scales of 1~Mpc$<s<$~100~Mpc is found to have similar correlation length $s_0 = 8.55^{+0.75}_{-0.74}$~Mpc, but the slope is shallower ($\\gamma = 1.3 \\pm 0.1$).  2. We study the projected correlation function of both CLASXS and CDFN. The best-fit parameters for the real-space correlation functions are found to be $r_0 = 8.1^{+1.2}_{-2.2}$~Mpc, $\\gamma = 2.1 \\pm 0.5$ for CLASXS field, and $r_0 = 5.8^{+1.0}_{-1.5}$~Mpc, $\\gamma = 1.38^{+0.12}_{-0.14}$ for CDFN field. Our result for the CDFN shows perfect agreement with the published results from Gilli et al. (2004). Fitting the combined data from both fields gives  $r_0 = 6.1^{+0.4}_{-1.0}$~Mpc and $\\gamma = 1.47^{+0.07}_{-0.10}$.  3. Comparing the redshift- and real-space correlation function of the combined CLASXS and CDFN fields, we found the redshift distortion parameter $\\beta = 0.4 \\pm 0.2$ at an effective redshift $z = 0.94$. Under the assumption of $\\Lambda$CDM cosmology, this implies a bias parameter  $b \\approx 2.04 \\pm 1.02$ at this redshift.  4. We tested whether the clustering of the X-ray sources is dependent on the X-ray spectra in the CLASXS field. Using a hardness ratio cut at $HR = 0.7$, we found no significant difference in clustering between hard and soft sources. This agrees with previous claims.  5. With the large dynamic range in X-ray luminosity, we found  a  weak correlation between X-ray luminosity and clustering amplitude. Using a simple model based on observations that links the AGN luminosity and halo mass, we show that the observed weak correlation is consistent with the model, except at low luminosities, where sources are likely to have lower Eddington ratio. The non-detection of a strong correlation between X-ray luminosity and clustering amplitude also suggests a narrow range of halo mass.  6. We study the evolution of the clustering using the redshift-space correlation function in 4 redshift intervals ranging from 0.1 to 3.0. We found only a mild evolution of AGN clustering in both CLASXS and CDFN samples. This qualitatively agrees with the results based on optically selected quasars from 2dF survey. The X-ray samples, however, show a similar correlation amplitude as that of the 2dF sample. This is consistent with the weak correlation between AGN luminosity and the clustering amplitude found in this work.  7. We estimate the evolution of bias by comparing the observed clustering amplitude with expectations of the linear evolution of density fluctuations. The result show that the bias increases rapidly with redshift ($b(z=0.45) \\sim 0.95$ and $b(z = 2.07) \\sim 3.03$ in CLASXS field). This agrees with the findings from 2dF.  8. Using the bias evolution model for dark halos from Sheth, Mo \\& Tormen (2001), we estimated the characteristic mass of AGNs in each redshift interval. We found the mass of the dark halo changes very little with redshift. The average  halo mass is found to be  $\\log~(M_{halo}/M_{\\sun}) \\sim 12.11$.  Our results have demonstrated that deep X-ray surveys are a very useful tool in studying how AGNs trace the large scale structure. Such knowledge provides an unique window to the understanding of AGN activity, and its relation to structure formation. The higher spatial density and much better completeness compared to current optical surveys allows us to study clustering on scales only accessible to very large optical surveys such as the 2dF and the SDSS. The high spatial resolution and positional accuracy of {\\it Chandra} is critical for unambiguous optical identifications. Since our results on the evolution of AGN clustering could still be affected by a small number of large scale structures, as seen in Chandra Deep Field South, larger {\\it Chandra} fields are still needed to improve the measurements."
    },
    "0601/astro-ph0601128_arXiv.txt": {
        "abstract": "Extrasolar planet transit surveys will also detect eclipsing binaries. Using a comprehensive binary population synthesis scheme, we investigate the statistical properties of a sample of eclipsing binaries that is detectable by an idealised extrasolar planet transit survey with specifications broadly similar to those of the SuperWASP (Wide Angle Search for Planets) project.  In this idealised survey the total number of detectable single stars in the Galactic disc is of the order of $10^6-10^7$, while, for a flat initial mass ratio distribution, the total number of detectable eclipsing binaries is of the order of $10^4-10^5$. The majority of the population of detectable single stars is made up of main-sequence stars ($\\approx 60$\\%), horizontal-branch stars ($\\approx 20\\%$), and giant-branch stars ($\\approx 10\\%$). The largest contributions to the population of detectable eclipsing binaries stem from detached double main-sequence star binaries ($\\approx 60\\%$), detached giant-branch main-sequence star binaries ($\\approx 20\\%$), and detached horizontal-branch main-sequence star binaries ($\\approx 10\\%$). White dwarf main-sequence star binaries make up approximately 0.3\\% of the sample.  The ratio of the number of eclipsing binaries to the number of single stars detectable by the idealised SuperWASP survey varies by less than a factor of 2.5 across the sky, and decreases with increasing Galactic latitude. It is found to be largest in the direction of the Galactic longitude $l=-7.5^\\circ$ and the Galactic latitude $b=-22.5^\\circ$.  We also show that the fractions of systems in different subgroups of eclipsing binaries are sensitive to the adopted initial mass ratio or initial secondary mass distribution, which is one of the poorest constrained input parameters in present-day binary population synthesis calculations. This suggests that once statistically meaningful results from transit surveys are available, they will be able to significantly improve the predictive power of population synthesis studies of interacting binaries and related objects. ",
        "introduction": "An ever increasing number of extrasolar planets is being discovered (to date the number stands at more than 160), mostly through the detection of the reflex motion of their parent star around the planetary system's centre of mass. Planets with suitably oriented orbits in space may also reveal their presence when they transit the disc of their parent star and the transit is accompanied by a measurable decrease in the observed stellar flux. When combined with high-precision radial-velocity measurements, the transits yield a wealth of information on the planet, such as its mass, radius, and mean density.  The first successful planetary transit detections were made independently by Charbonneau et al. (2000) and Henry et al. (2000), who observed a clear drop in the lightcurve of the star HD\\,209458 at times consistent with the ephemeris of its already known planetary-mass companion.  Since then, the search for transits has rapidly evolved into a fully-fledged and widely used technique to hunt for extrasolar planets.  In anticipation of upcoming space-based exoplanet transit surveys such as COROT (scheduled to launch in 2006) and Kepler (scheduled to launch in 2007), ground-based surveys can pave the way towards obtaining a statistically significant sample of planets in a relatively short amount of time. Such a sample will be of great importance in improving our understanding of the formation and evolution of extrasolar planets as well as for detecting possible correlations between the presence of planets and the properties of their host stars. For an overview of some existing exoplanet transit surveys, we refer to the reviews by Horne (2002, 2003) and Charbonneau (2003). Despite the multitude of exoplanet transit surveys, so far only six confirmed planets have been discovered by direct observations of transits: five by the Optical Gravitational Lensing Experiment (OGLE) (Konacki et al. 2003; Konacki et al. 2004; Bouchy et al. 2004; Pont et al. 2004; Konacki et al. 2005), and one by the Trans-Atlantic Exoplanet Survey (TrES) network (Alonso et al. 2004).  Besides planetary transit detections, wide-field photometric transit searches will also yield a wealth of photometric data on all kinds of new and existing variable stars, including eclipsing binaries. Brown (2003) estimated the rate of false planetary detections due to eclipsing binaries with main-sequence and giant-type component stars to be almost an order of magnitude larger than the rate of true planetary detections.  Examples of astrophysical false positives in wide-field transit searches are grazing incidence eclipses and the blending of the light of an eclipsing binary with the light of a third star (see Charbonneau et al. 2003 for details).  Our aim in this exploratory paper is to perform a binary population synthesis study to predict the statistical properties and the Galactic distribution of eclipsing binaries detectable by an extrasolar planet transit survey, and to investigate their dependency on the initial mass ratio distribution of the component stars.  Specifically, we study the detectable sample for an idealised survey that is broadly representative of the ground-based SuperWASP Wide Angle Search for Planets (Kane et al. 2003, 2004; Street et al. 2003, Christian et al. 2004, Pollacco 2005). The SuperWASP setup currently consists of 5 ultra-wide field lens cameras backed by high-quality CCDs, but is designed to hold up to a total of 8 such cameras and CCDs (the upgrade to 8 cameras is expected to be completed in 2005). The photometric precision of the CCDs is largest in the 7--13\\,mag magnitude range, where it is aimed to be better than 10\\,mmag (see Kane et al. 2004 for a detailed discussion of the photometric accuracy of the prototype WASP0).  Each camera on the SuperWASP setup furthermore has an approximate field of view of $8\\degr \\times 8\\degr$, so that, accounting for some overlap, the total area covered by 4 cameras operating simultaneously in a rectangular configuration is about $15\\degr \\times 15\\degr$. The total field of view for the full 8-camera setup will be $30\\degr \\times 15\\degr$. The large field of view is expected to yield precise photometry of approximately 5\\,000 to 10\\,000 stars at a time per operational camera, and thus 40\\,000 to 80\\,000 stars at a time for 8 cameras operating simultaneously.  In the following sections, we study in some detail the relative numbers and types of eclipsing binaries detectable in the idealised SuperWASP survey and examine their dependence on the initial mass ratio or secondary mass distribution. By considering the ratio of the number of detectable eclipsing binaries to the number of detectable single stars we also investigate if some regions in the Galaxy are expected to produce more false planetary detections due to eclipsing binaries than others. Our method of calculation is outlined in some detail in Section 2, while the results are presented in Sections 3 and 4. The final section is devoted to concluding remarks. ",
        "conclusions": "In this exploratory study we used the BiSEPS binary population synthesis code to estimate the number of eclipsing binaries and single stars detectable in the Galactic disc by an idealised SuperWASP exoplanet transit survey. The Galactic disc is modeled by a superposition of a young (thin) and old (thick) component, both of which are described by means of a double exponential distribution function characterised by different scale lengths and scale heights. The models yield a strong concentration of eclipsing binaries and single stars towards the Galactic plane and the Galactic centre.  The ratio of the number of detectable eclipsing binaries to the number of detectable single stars varies by less than a factor of 2.5 across the sky. It is largest in the direction of the Galactic longitude $l=-7.5^\\circ$ and the Galactic latitude $b=-22.5^\\circ$, so that the number of false exoplanet transit detections due to eclipsing binaries can be expected to be largest in this direction.  More favourable regions on the sky aimed at minimizing false planetary detections are therefore located away from this area. The best conditions are found at Galactic longitudes $75^\\circ \\la l \\la 285^\\circ$ and, for Galactic longitudes $|l| \\la 75^\\circ$, at Galactic latitudes $|b| \\ga 15^\\circ$.  Depending on the adopted survey model, the total number of single stars in the magnitude range detectable by the idealised SuperWASP survey is of the order of $10^6$--$10^7$. The total number of detectable eclipsing binaries varies from a few times $10^2$--$10^5$, with the exact number depending on the adopted survey model (limiting magnitude and eclipse amplitude) and on the initial mass ratio or initial secondary mass distribution. In view of the large number of stars and binaries observed by SuperWASP each night, the lower end of the predicted number of eclipsing binaries (resulting from the initial mass ratio distribution $n(q) \\propto q^{-0.99}$) could soon be confronted with interesting and potentially stringent observational constraints, assuming that our simplifications still capture the essence of the real SuperWASP set-up.  Provided that complete samples of the most abundant types of binaries can be compiled by SuperWASP, the relative contribution of each type of binary to the total population can be used to further constrain the initial mass ratio or initial secondary mass distribution (see Fig.~\\ref{pop}). In particular, for our standard survey model, the relative number of double main-sequence star binaries increases from 45\\% in the case of the initial mass ratio distribution $n(q) \\propto q$ to 80\\% in the case of the initial mass ratio distribution $n(q) \\propto q^{-0.99}$, and 90\\% in the case where the initial secondary mass is distributed independently according to the same IMF as the primary mass. The contribution of white dwarf main-sequence star binaries increases from 0.1\\% in the case of the initial mass ratio distribution $n(q) \\propto q$ to 0.7\\% in the case of the initial mass ratio distribution $n(q) \\propto q^{-0.99}$, and 1.3\\% in the case where the initial secondary mass is distributed independently according to the same IMF as the primary mass. In contrast, the relative contributions of the other major groups of binaries tend to decrease when the initial mass ratio distribution is changed from $n(q) \\propto q$ to $n(q) \\propto q^{-0.99}$ or an independent initial secondary mass distribution. Similar diagnostics for the initial mass ratio or initial secondary mass distribution can be obtained by considering ratios of numbers of different types of eclipsing binaries detectable by the transit survey rather than their fractional contributions to the total population.  In view of the large variety of eclipsing binaries found, we did not consider the effects of varying some of the binary evolution parameters as is customary in binary population synthesis studies (see, e.g., Willems \\& Kolb 2002, 2003, 2004). In particular, the masses and orbital periods of the eclipsing binaries may be affected by uncertainties in the treatment of mass transfer if either of the binary components at some point of its evolution filled its critical Roche lobe. Additional uncertainties affecting the binary separation (and thus the possibility of interactions between the component stars as well as the probability to observe the system as an eclipsing binary) are introduced by uncertainties in systemic orbital angular momentum losses through, e.g., magnetic braking and stellar winds. The latter can furthermore also have a strong impact on the stars themselves, especially during or after evolutionary phases that are accompanied by high wind mass-loss rates such as the giant branch, asymptotic giant branch, and Wolf-Rayet stages of stellar evolution. The impact of these uncertainties on the relative numbers and statistical properties of different types of eclipsing binaries will be addressed in follow-up studies that are tailored more closely to represent the precise properties of different ongoing and future exoplanet transit surveys.  The results of our exploratory study reveal very encouraging differential trends in the collective properties of the eclipsing binary population that can be probed by extrasolar planet transit surveys. This suggests that transit surveys will be able to deliver stringent constraints on parameters describing the formation and evolution of binaries once statistically meaningful results from such surveys are available. In future applications of this work we will consider a fine-tuned transit survey model that more closely matches the actual parameters of the SuperWASP survey, and that also addresses the fact that the discovery probability of eclipses will depend on the eclipse properties. We will also extend our study to other transit surveys."
    },
    "0601/astro-ph0601402_arXiv.txt": {
        "abstract": "{The evolution of the cluster abundance with redshift is known to be a powerful cosmological constraint when applied to X-ray clusters. Recently, the evolution of the baryon mass function has been proposed as a new variant that is free of the uncertainties present in the temperature-mass relation. A flat model with $\\Omega_{\\rm M} \\simeq 0.3$ was shown to be preferred in the case of a standard cold dark matter scenario.} {We compared the high redshift predictions of the baryon mass in clusters with data for a more general class of spectra with a varying $shape$ factor $\\Gamma$ without any restriction to the standard cold dark matter scenario in models normalized to reproduce the local baryon mass function.} {Using various halo mass functions existing in the literature we evaluated the corresponding baryon mass functions for the case of the non-standard power spectra mentioned previously.} {We found that models with $\\Omega_{\\rm M} \\simeq 1$ and $\\Gamma \\simeq 0.12$ reproduce high redshift cluster data just as well as the concordance model does.} {Finally, we conclude that the baryon mass function evolution alone does not efficiently discriminate between the more general family of flat cosmological models with non-standard power spectra.} ",
        "introduction": "The growth of structure in the Universe is believed to be the result of gravitational collapse generated by the existence of tiny departures from homogeneity and isotropy (presumably generated during inflation) in the primordial distribution of matter (see e.g., Peacock 1999). The overdensity of matter in a particular comoving scale $l$ evolves according to linear theory until it reaches a value of $\\delta_k \\sim 1$, where we have used $k \\equiv |\\bf{k}|$ with $\\bf{k}$ the comoving {\\it wavevector} that satisfies the relation $|\\bf{k}|$=$2\\pi/l$. For later times the evolution is highly nonlinear and the formation of bound structures like galaxies and  galaxy clusters takes place. It is believed that this hierarchical process of structure formation is still at work today. Press \\& Schechter (1974, hereafter PS) developed a semianalytical formulation to deal with this regime and eventually predict the number of collapsed objects (often called {\\it virialized} objects) of a given mass $M$ (associated with the scale $l$) at a given redshift $z$, i.e. to determine the so-called {\\it mass function}. Interest in the PS approach has grown in recent years because it appears to reproduce the results of  numerical simulations well (e.g. Efstathiou et al. 1988; White, Efstathiou \\& Frenk 1993; Lacey \\& Coley 1994) where the nonlinear regime can be tested unambiguously inside the limits imposed by the resolution of the simulation. Sheth \\& Tormen (1999) obtained a different expression for the mass function by assuming an {\\it ellipsoidal} collapse instead of a {\\it spherical} one (as assumed in the PS theory) and found better agreement between the model and the numerical results. Sheth, Mo, \\& Tormen (2001, hereafter SMT) present an improved version of their 1999 work. Instead of working out a semianalytical approach Jenkins et al. (2001, hereafter J01) have introduced several fits to the numerical simulations using different algorithms for the halo finder and for several kinds of cosmologies. Further analyses have been made of these matters such as, White (2002) and Warren et al. (2005). Of particular interest is the evolution of the mass function with redshift that has been shown to be sensitive primarily to the mass density of the Universe (Blanchard \\& Bartlett 1998). This high sensitivity allows us to use the evolution of the abundance of X-ray clusters as a powerful cosmological test (Oukbir \\& Blanchard 1992). Recently, the {\\it baryon} mass function (i.e., the number density of comoving objects with a given baryon mass) has been advocated by Vikhlinin et al. (2003, hereafter V2003) as an useful alternative and cosmological constraints were derived in the context of standard cold dark matter (hereafter CDM) spectrum. V2003 have consequently concluded that cluster data favors a concordance-like Universe. This analysis seems to conflict with the study made by Blanchard et al. (2000) and with the  recent $XMM-Newton~\\Omega$-project (Vauclair et al. 2003), which is somewhat surprising because Sadat et al. (2005) found that gas fraction in distant clusters within the $XMM-Newton~\\Omega$-project was consistent with a high-density Universe and not with a concordance model. These two sets of analyses therefore suggest that the baryon mass fraction should also  be consistent with a high-density Universe. The present paper aims to clarify this issue. Here we study the more general class of spectra with a varying {\\it shape} factor $\\Gamma$, without any restriction to the standard CDM model. In Sect. 2 we briefly review the PS formalism and recall the results of SMT and J01. In Sect. 3 we introduce the temperature-mass (hereafter $T-M$) relation, and we use the different available expressions for the baryon mass function, normalized to local cluster data, to compare them with the high redshift observations provided by V2003. Finally, we discuss our main conclusion that data on the baryon mass function of clusters can also be reproduced in a critical Universe with $\\Omega_{\\Lambda} \\simeq 0$, indicating that there is no discrepancy in both approaches. ",
        "conclusions": "We have found that a non-standard ($\\Gamma$-varying) spectra can reproduce the observed baryon mass distribution function both for local and high redshifts in the case of an Einstein-de Sitter cosmology just as well as the concordance model does. For this, we used several formulas for the baryon mass function and found that the results are essentially insensitive to the expression used. It is important to recall that the V2003 analysis is based on the assumption of standard CDM power spectra for matter and on a particular value of the Hubble constant. This assumption fixes the value of the shape factor essentially in $\\Gamma \\simeq \\Omega_{\\rm M} h$ (e.g. Peacock \\& Dodds 1994). Furthermore, as they have used $h \\simeq 0.65$, i.e. a Hubble constant value near the Hubble Key Project result (Freedman et al. 2001), they would get a shape factor of $\\Gamma \\simeq 0.65$ in an Einstein-de Sitter Universe, which is very far from the value we get here for the same cosmology, i.e. $\\Gamma = 0.12$.  It is well known that the standard CDM model is ruled out in an Einstein-de Sitter cosmology, which is the reason for needing a non-standard CDM spectra in the $\\Omega_{\\rm M} \\simeq 1$ case. This difference explains the apparently conflicting conclusions between V2003 and the present work. While the Einstein-de Sitter and concordance models produce distinct cluster number counts for a flux limited survey (Vauclair et al. 2003) or for the temperature distribution function evolution (Blanchard et al. 2000), once models are properly normalized at low redshift, the baryon mass function does not differentiate among flat cosmologies, because the inferred baryon masses are different, and this difference accidentally and roughly compensates for the effect of number evolution. This kind of non-standard spectra with lower $\\Gamma$ on cluster scales could have originated from different hypotheses on the dark matter content and/or from the existence of some unknown phase in the evolution of the Universe (like hot dark matter or quintessence) that could drive different power spectra for matter. The initial power spectrum could also be altered from a single power law by physics at the inflation period. We are therefore lead to the final conclusion that present baryon cluster data can be described equally well by either a concordance or an Einstein-de Sitter (with a non-standard $\\Gamma$ value on cluster scales) Universe implying that the baryon mass function is not as effective as the evolution of the temperature function. This removes the apparent discrepancies between the conclusions inferred from the $XMM-Newton~\\Omega$-project (Vauclair et al. 2003, Sadat et al. 2005) and V2003."
    },
    "0601/astro-ph0601459_arXiv.txt": {
        "abstract": "We are interested in electrons kinetics in a stellar atmosphere to validate or invalidate the usually accepted hypothesis of thermalisation of electrons. For this purpose, we calculate the velocity distribution function of electrons by solving the kinetic equation of these particles together with the equations of radiative transfer and statistical equilibrium. We note that this distribution can deviate strongly from a Maxwell-Boltzmann distribution if non-LTE  effects are important. Some results and astrophysical consequences are examined. ",
        "introduction": "Our work is concerned with the kinetics of electrons in a stellar atmosphere, modelled as a parallel-plane slab irradiated on a face. Our models of atmospheres start in the deep layers of stars, where the radiative field can be described in the diffusion approximation, and end with the layers of minimal temperature, before the chromospheric raise whose effects are ignored. The free electrons are characterized by their velocity distribution function : the electron distribution function (edf), which is calculated with the other thermodynamical quantities of the atmosphere. Our main objective is to understand the mechanism leading to the thermalisation of electrons, where the edf tends toward the Maxwell-Boltzmann distribution. It is accepted, in stellar atmospheres theory, that the thermalisation of electrons is effective as long as elastic collisions dominate inelastic interactions of electrons with the plasma, a rather well verified hypothesis for electrons having energies greatly below the first excitation energies of atoms and ions composing the atmosphere. This hypothesis is not necessarily correct for faster electrons. Our work follows the line drawn by some plasma physicists at the beginning of the 70s (Peyraud 1968, 1970; Peyraud \\etal\\, 1969; Oxenius 1970, 1974; Shoub 1977). Their work demonstrated the important role played by inelastic (collisional or radiative) processes in the equilibrium reached by electrons, which can deviate considerably from the maxwellian equilibrium at high energies. We present below a stellar atmosphere model which is not in local thermodynamical equilibrium (LTE), confirming the results anticipated in the 70s on the basis of mainly theoretical developments. ",
        "conclusions": "} Astrophysical consequences of this work are numerous. In general the deviation of the edf from a maxwellian distribution has a direct effect on all thermodynamic quantities involving the edf, \\emph{e.g.} collisional transition rates or spectral lines profiles. It has an indirect effect on all other characteristics of the atmosphere, because of the coupling of all equations. Transition rates are used to solve the equations of statistical equilibrium, which lead to the populations and ionization degree of the atmosphere. Inversion techniques of spectral lines observed by spectroscopy are also sensitive to the edf shape, so that temperatures or densities calculated with these techniques are \\apriori\\, affected by deviation of the edf from a maxwellian distribution (\\cf \\, Shoub 1983, Owocki \\etal\\, 1983, Salzmann \\etal\\, 1995). Finally the non thermodynamical equilibrium of electrons may be at the origin of physical processes which are still not very well understood at present, for example the heating of the sun corona (\\cf \\, Scudder 1992,1994). \\newline The results presented in this article are based on a very simple atmosphere model, which guarants a numerically stable solution. In non-LTE regions close to the surface, the edf shows important deviations from the maxwellian distribution in the fast electrons tail. Our model confirm, by means of numerical codes accurate enough to handle this complex problem, most of the physical ideas advanced thirty years ago. Our main contribution consists in the construction of a selfconsistent model of a stellar atmosphere with \\apriori\\, non thermalized electrons. Also we have used very accurate radiative transfer codes. It remains to make this model more realistic for comparison with observations."
    },
    "0601/astro-ph0601490_arXiv.txt": {
        "abstract": "We investigate the in-spiraling timescales of globular clusters in dwarf spheroidal (dSph) and dwarf elliptical (dE) galaxies, due to dynamical friction. We address the problem of these timescales having been variously estimated in the literature as much shorter than a Hubble time. Using self-consistent two-component (dark matter and stars) models, we explore mechanisms which may yield extended dynamical friction timescales in such systems in order to explain why dwarf galaxies often show globular cluster systems. As a general rule, dark matter and stars both give a comparable contribution to the dynamical drag. By exploring various possibilities for their gravitational make-up, it is shown that these studies help constrain the parameters of the dark matter haloes in these galaxies, as well as to test alternatives to dark matter. Under the assumption of a dark halo having a central density core with a typical King core radius somewhat larger than the observed stellar core radius, dynamical friction timescales are naturally extended upwards of a Hubble time. Cuspy dark haloes yield timescales $\\lesssim$ 4.5 Gyr, for any dark halo parameters in accordance with observations of stellar line-of-sight velocity dispersion in dwarf spheroidal galaxies. We confirm, after a detailed formulation of the dynamical friction problem under the alternative hypothesis of MOND dynamics and in the lack of any dark matter, that due to the enhanced dynamical drag of the stars, the dynamical friction timescales in MOND would be extremely short. Taking the well-measured structural parameters of the Fornax dSph and its globular cluster system as a case study, we conclude that requiring dynamical friction timescales comparable to the Hubble time strongly favours dark haloes with a central core. ",
        "introduction": "The dynamical friction (DF) timescale for globular clusters (GCs) to sink to the centre of most of dwarf spheroidals (dSph) and dwarf elliptical (dE) galaxies yields a fraction of the Hubble time, when dark halo parameters are taken straight from measurements of the stellar distributions (e.g., Tremaine 1976; Hernandez \\& Gilmore 1998b; Oh, Lin \\& Richer 2000; Lotz et al.~2001). Both the stellar bulk of the dwarf galaxy and the dark matter contribute to the DF. In some dE, the DF timescale with the dark matter alone is shorter than one Hubble time, in some dSph, the dynamical friction due to the stars alone would suffice to decay the GC orbits in a few Gyr, under the above assumption. Tremaine (1976) first noticed that using the preliminary values for the radius and mass of the second most luminous of the $10$ dSph satellites of the Milky Way, Fornax, this time is $\\sim 1$--$2$ Gyr, very short as compared to absolute ages estimated for these clusters ($\\sim 14.6\\pm 1.0$ Gyr for clusters $1$--$3$ and $5$, $\\sim 11.6$ Gyr for cluster $4$, see Buonanno et al.~1998 and Mackey \\& Gilmore 2003). One would expect the formation of bright nuclei from the merger of orbitally decayed GCs. Paradoxically, Fornax contains five GCs, an unusually high GC frequency for its dynamical mass, and does not present a central nucleus. Although one of the clusters is observed located near the centre, it presents a radial velocity comparable to the mean velocity of the field stars (Dubath et al. 1992). Hence, this globular cluster has not been really dragged to the Fornax centre through dynamical friction.  Lotz et al.~(2001) found that the nuclei in $M_{V}>-14$ dE's are several magnitudes fainter than expected\\footnote{Note that Fornax has a magnitude $M_{V}=-12.5$.}. Indeed, there are dE with faint nuclei and high globular cluster specific frequencies. Therefore, either the DF timescale is being underestimated or, in these galaxies, some mechanisms are working against DF to prevent the GC system from collapsing into a bright nucleus (see Lotz et al.~2001 for a discussion of the different possibilities).   In this paper we reconsider the problem of the orbital decay of globular clusters in dwarf galaxies. In the case of central cluster galaxies, a variety of processes may contribute to altering the kinematics of GCs which are beyond the scope of this paper. Using Fornax as a case study, we examine the orbital evolution of GCs in both the conventional dark matter scenario and under MOdified Newtonian Dynamics (MOND; Milgrom 1983). Although our work is aimed at exploring possible explanations for the survival of GCs in dSph and dE, our analysis will be focused mainly on the Fornax dSph because the recent detailed determinations of the structural and dynamical parameters of both Fornax and its GC population (i.e.~position, age, mass and radial velocity of the GCs are known), allow us to test our assumptions. Fornax is the prototypical dSph with a high GC frequency, has no nucleus, and short DF timescale estimates.  Since DF is a gravitational effect, our analysis is relevant for any massive perturber, such as black holes, star clusters or any kinematically cold massive substructure (e.g., Kleyna et al.~2004) and not only for GCs. Studies of the survival of GC systems and density substructure can provide useful constraints on the dark halo profile in dSph and dE (e.g., Hernandez \\& Gilmore 1998b), and could shed light on the debate of cuspy/cored haloes, which is currently a test for the standard $\\Lambda$CDM paradigm. On the other hand, since modified Newtonian dynamics (MOND; Milgrom 1983) can explain the dynamics of spiral galaxies without any dark matter (e.g., Sanders 1996; McGaugh \\& de Blok 1998), and perhaps the dynamics of dSph (Lokas 2001, 2002), one could ask whether MOND offers a more natural solution to the problem of the rapid orbital evolution of the globular clusters in dwarf galaxies. Since MOND is solely determined by the luminous material, the DF timescale can be used as a test of modified dynamics.  In \\S 2 we give a general statement of the problem, reviewing estimates of the DF timescales for the GCs in dwarf galaxies, and discussing the processes that may be working against the orbital decay of GC systems and their associated difficulties. Section 3 presents a solution to the dynamical-friction problem in a self-consistent two-component dynamical system, dark halo and stars, giving analytical approximate solutions and full numerical integrations for a range of possible dark halo profiles. In \\S 4 we explore the resulting in-spiraling timescales in the MOND scenario, for various relevant limits of such theory. We find that the problem of the sedimentation of GCs has a simple solution in the dark matter scenario, whereas it appears problematic in the MOND scenario. Our conclusions are summarized in \\S 5. ",
        "conclusions": "Long before relevant data was available and using a simple but transparent scaling of the dynamical friction timescale, Tremaine (1976) warned that nuclei should be formed in the centre of dSph galaxies by coalescence of GCs. He noticed that NGC 185, NGC 147 and Fornax all have GCs but no nucleus. More recently, and for a sample of $51$ dwarf elliptical galaxies in nearby clusters, Lotz et al.~(2001) examined the radial distribution of GCs and found that for the fainter dE's, the predicted nuclei was several magnitudes brighter than observed. Since the composite radial distribution of the GCs follows the exponential stellar profile of dE's at present day, the simplest hypothesis is that the GCs initially followed the exponential profile of the underlying stellar population and that the dynamical friction force was partly inhibited. Note that the drag should not be completely suppressed in order to explain the significant deficit of bright, massive GCs in the inner regions of the sample of Lotz et al.~(2001). In Local Group dSph galaxies, the problem of the survival of enhanced density substructure with cold kinematics against dynamical friction has been revived recently (Kleyna et al.~2004; Walker et al.~2006). For the less massive galaxies, the tidal field could play a role in heating their GCs but they must be tidally disrupted and suffer significantly mass loss ($50$--$90\\%$) in order to counter the drag of dynamical friction. If so, these galaxies should present a rising projected velocity dispersion profile beyond a critical radius caused by tidal stripping (Read et al.~2006). In contrast, the Local Group dSph galaxies show a very flat or falling projected velocity dispersion, which indicates that tidal stripping is unimportant interior to $\\sim 1$ kpc.  In this paper we have shown that all these features can be explained by considering dark haloes as having density cores. Taking Fornax as a reference case, we have explored under which structural parameters of the dark halo is the orbital decay of GCs minimized (see also Goerdt et al.~2006 for an independent and complementary study). Using self-consistent dynamical models where both the stars and dark matter contribute to the frictional drag, we have computed the radial evolution of a massive GC initially on a circular orbit. We find that the contributions to dynamical friction of stars and dark matter are both comparable in dwarf galaxies. The survival, against dynamical friction, of GCs requires a King core radius for the dark halo slightly larger than that of the underlying stellar population. In the particular case of Fornax, we infer a dark halo with the following parameters: a King core radius of $\\sim 1.5$ kpc (or $r_{c,-0.1}=0.3$ kpc), a central velocity dispersion of $\\sim 30$ km s$^{-1}$ and a total mass of $1.2\\times 10^{9}$ M$_{\\odot}$. In a dark matter halo with a cuspy inner profile, the dynamical friction timescale is $\\lesssim 4.5$ Gyr. The indirect evidence of a central density core in dSph has important implications for the formation of dSph galaxies and for cosmology, and gives new insight to the ongoing debate on the inner slope of the dark halo profiles in low surface brightness galaxies. The present study puts constraints on the epoch of core formation. If cored haloes are solely responsible for the survival of the present-day GC configuration, the formation of the core should have proceeded in an early phase of the galaxy's formation, at least $6$ Gyr ago, to prevent excessive drag of Fornax GCs.  Under the MOND hypothesis, dynamical friction timescales for dwarf galaxies are necessarily short, $\\sim 1$ Gyr. Therefore, alternative explanations must be found for the observed GC systems. No such entirely satisfactory alternative has been presented to date, making old GC systems in dwarfs a strong objection to MOND under its current formulation."
    },
    "0601/astro-ph0601173_arXiv.txt": {
        "abstract": "{ We present observations of XRF 050406, the first burst detected by Swift showing a flare in its X-ray light curve. During this flare, which peaks at $t_{\\rm peak} \\sim 210$\\,s after the BAT trigger, a flux variation of $\\delta F / F \\sim 6$ in a very short time $\\delta t / t_{\\rm peak} \\ll 1$ was observed. Its measured fluence in the 0.2--10 keV band was $\\sim 1.4 \\times 10^{-8}$ erg cm$^{-2}$, which corresponds to 1--15\\% of the prompt fluence. We present indications of spectral variations during the flare. We argue that the producing mechanism is late internal shocks, which implies that the central engine is still active at $210$\\,s, though with a reduced power with respect to the prompt emission. The X-ray light curve flattens to a very shallow slope with decay index of $\\sim 0.5$ after $\\sim 4400$\\,s, which also supports continued central engine activity at late times. This burst is classified as an X-ray flash, with a relatively low fluence ($\\sim 10^{-7}$ erg cm$^{-2}$ in the 15--350 keV band, $E_{\\rm iso} \\sim 10^{51}$ erg), a soft spectrum (photon index 2.65), no significant flux above $\\sim 50$ keV and a peak energy $E_{\\rm p} < 15$ keV. XRF 050406 is one of the first examples of a well-studied X-ray light curve of an XRF. We show that the main afterglow characteristics are qualitatively similar to those of normal GRBs. In particular, X-ray flares superimposed on a power-law light curve have now been seen in both XRFs and GRBs. This indicates that a similar mechanism may be at work for both kinds of events. ",
        "introduction": "} The Swift Gamma-Ray Burst Explorer (\\citealt{SWIFT}) was successfully launched on 2004 Nov 20. Its payload includes one wide-field instrument, the gamma-ray (15--350 keV) Burst Alert Telescope (BAT; \\citealt{BAT}), and two narrow-field instruments, the X-Ray Telescope (XRT; \\citealt{XRT2}) and the Ultraviolet/Optical Telescope (UVOT; \\citealt{UVOT}). BAT detects the bursts, calculates their position to $\\sim 1$--$4 \\arcmin$ accuracy and triggers an autonomous slew of the observatory to point the two narrow-field instruments. The XRT, which operates in the 0.2--10\\,keV energy range, can provide $\\sim 5\\arcsec$ positions, while the UVOT, which operates in the 1700--6000\\,\\AA{} wavelength range, can further refine the afterglow localization to $\\sim 0\\farcs5$. With its unique fast re-pointing capabilities Swift set out to investigate the very early phases of gamma-ray burst (GRB) afterglows, beginning as early as one minute after the BAT trigger. During the initial activation and calibration phases, which ended on 2005 Apr 5, Swift discovered 25 GRBs. The narrow-field instruments were re-pointed towards seven of them within a few hundred seconds, and such is the case for GRB 050406.  On 2005 Apr 6 at 15:58:48.40 UT, BAT triggered on GRB 050406 (trigger 113872; \\citealt{GCN3180}), and located it at RA(J2000$) = 02^{\\rm h}  17^{\\rm m}  53^{\\rm s}$, Dec(J2000$) =-50^{\\circ}  10\\arcmin  52\\arcsec$, with an uncertainty of 3 arcmin (95\\% containment; \\citealt{GCN3183}). The derived value for the time during which 90\\% of the burst fluence is observed was $T_{90}=5\\pm1$\\,s in the 15--350 keV band. In the 15--25 keV band the light curve peak had a fast-rise, exponential decay (FRED) profile, while in the 25--50 keV band, the shape was more symmetric, with the peak starting $\\sim 2$\\,s earlier (\\citealt{GCN3183}). Both the peak and time-averaged spectra were well fit by a simple power-law with a time-averaged spectrum photon index of 2.38$\\pm$0.34 (90\\% confidence; \\citealt{GCN3183}). The fluence in the 15--350 keV band was $9.0 \\times 10^{-8}$ erg cm$^{-2}$. The gamma-ray characteristics of this burst, i.e.\\ the softness of the observed spectrum and the absence of significant emission above $\\sim 50$ keV, classify GRB 050406 as an X-ray flash (XRF; \\citealt{Heiseea01}). From now on, we shall therefore refer to this event as XRF 050406.  Swift executed a prompt slew. The XRT imaged the BAT field only 84\\,s after the trigger but no bright X-ray source could be detected within the field of view. However, a refined on-ground analysis revealed a previously uncatalogued X-ray source (\\citealt{GCN3181,GCN3184}). From the very first examination of the down-linked data it was clear that the afterglow of this burst was peculiar. Indeed, after an initial decay, the X-ray count rate began rising, peaking at $\\approx 220$\\,s, and subsequently decaying again (\\citealt{GCN3184}).  Ground-based observations started as soon as the burst discovery was reported via the GCN network. The Magellan/Clay telescope imaged the XRT error circle with LDSS-3 in the $R$ and $i$ bands and found a single faint source ($R=22.0\\pm0.09$ mag, 7.8 hr after the burst) located at RA(J2000$)=02^{\\rm h} 17^{\\rm m} 52\\fs3$, Dec(J2000$)=-50^{\\circ} 11\\arcmin 15\\arcsec$ with an uncertainty of $\\sim 0\\farcs5$ in each coordinate (\\citealt{GCN3185,Bergerea05b}). Similarly to XRT, UVOT also imaged the field at the end of the slew (starting from $\\sim 88$ s after the trigger) and though it failed to detect the afterglow on-board (\\citealt{GCN3182}), subsequent on-ground analysis revealed a source within the XRT error circle at the 4.3- (19.0 mag), 3.0- and 2.5-$\\sigma$ detection levels in the $U$, $B$ and $V$ bands, respectively. The UVOT position was RA(J2000$)=02^{\\rm h} 17^{\\rm m} 52\\fs2$, Dec(J2000$)=-50^{\\circ} 11\\arcmin 15\\farcs8$, consistent with the Magellan one. By the time the second UVOT observation (1.3 hr later) was performed, the source was not detected in the $U$ band, confirming it  as the afterglow of XRF 050406. \\citet{Schady06} obtained an estimate of $z=2.44\\pm0.36$ from fitting the broad band spectrum (combined UVOT and XRT data).   In this paper we present observations of the first Swift burst where a flare is clearly detected in its X-ray light curve, during which the source count rate increased by a factor of $\\ga 6$. This feature had never been observed before in Swift data, and had rarely been observed before in any X-ray afterglow (\\citealt{piroea05}). This paper is organized as follows. In Sect.\\ \\ref{grb050406:dataredu} we describe our observations and data reduction; in Sect.\\ \\ref{grb050406:dataanal} we describe our spatial, timing and spectral data analysis; in Sect.\\ \\ref{grb050406:discussion} we discuss our findings. Finally, in Sect.\\ \\ref{grb050406:conclusions} we summarize our conclusions. Throughout this paper the quoted uncertainties are given at 90\\% confidence level for one interesting parameter (i.e., $\\Delta \\chi^2 =2.71$) unless otherwise stated. Times are referred to the BAT trigger $T_0$, $t=T-T_0$. The decay and spectral indices are parameterized as follows, $F(\\nu,t) \\propto t^{-\\alpha} \\nu^{-\\beta}$, where $F_{\\nu}$ (erg cm$^{-2}$ s$^{-1}$ Hz$^{-1}$) is the monochromatic flux as a function of time $t$ and frequency $\\nu$; we also use $\\Gamma = \\beta +1$ as the photon index, $N(E) \\propto E^{-\\Gamma}$ (ph keV$^{-1}$ cm$^{-2}$ s$^{-1}$). ",
        "conclusions": "} %  XRF 050406 is classified as an X-ray flash, with fluence $\\sim 1 \\times 10^{-7}$ erg cm$^{-2}$ (15--350 keV), a soft spectrum ($\\Gamma_\\gamma=2.65$), no significant flux above $\\sim 50$ keV and a peak energy $E_{\\rm p} < 15$ keV. Its main characteristics are however not qualitatively different from those of normal GRBs. As observations accumulate, it becomes clear that these two classes of phenomena share many properties, and both have afterglows with similar characteristics. This is a clue that both events may have a common origin.  XRF 050406 is the first Swift-detected burst that showed a flare in its X-ray light curve, a feature  now found in $\\sim 50$\\,\\% of the XRT afterglows. The flare peaked at $\\sim 210$ s after the BAT trigger ($\\sim 61$ s in the rest frame). The best fit of the afterglow decay is obtained with a broken power law with $\\alpha_1=1.58\\pm0.17$,  $\\alpha_2=0.50^{+0.14}_{-0.13}$, and a break at $\\sim 4400$ s after the BAT trigger. The mean photon index is $\\Gamma_{\\rm X} = 2.1\\pm0.3$. During the X-ray flare a flux variation of $\\delta F / F_{\\rm peak} \\sim 6$ in a timescale $\\delta t / t_{\\rm peak} \\ll 1$ is observed, and its measured fluence in the 0.2--10 keV band is $\\sim 1.4 \\times 10^{-8}$ erg cm$^{-2}$ [$(2.0 \\pm 1.4) \\times 10^{50}$ erg], which corresponds to 1--15\\% of the prompt fluence. We argued that the flare-producing mechanism is late internal shocks, which implies that the central engine is still active at $t \\sim 210$\\,s, though with a reduced power with respect to the prompt emission. We showed possible indications of spectral variations during the flare, and a flattening of the X-ray light curve after  $t \\sim 4400$\\,s in support of continued central engine activity at late times.  Since XRF 050406 was observed, flares have been detected by XRT in both X-ray flashes and normal GRBs, indicating that flares are linked to some common properties of both kinds of bursts, and probably tied to their central engine."
    },
    "0601/astro-ph0601345_arXiv.txt": {
        "abstract": "Elements of kinematical and dynamical modeling of elliptical galaxies are presented. In projection, NFW models resemble S\\'ersic models, but with a very narrow range of shapes ($m=3\\pm1$). The total density profile of ellipticals cannot be NFW-like because the predicted local $M/L$ and aperture velocity dispersion within an effective radius ($R_e$) are much lower than observed. Stars must then dominate ellipticals out to a few $R_e$. Fitting an NFW model to the total density profile of S\\'ersic+NFW (stars+dark matter [DM]) ellipticals results in very high concentration parameters, as found by X-ray observers. Kinematical modeling of ellipticals assuming an isotropic NFW DM model underestimates $M/L$ at the virial radius by a factor of 1.6 to 2.4, because dissipationless $\\Lambda$CDM halos have slightly different density profiles and slightly radial velocity anisotropy. In $N$-body+gas simulations of ellipticals as merger remnants of spirals embedded in DM halos, the slope of the DM density profile is steeper when the initial spiral galaxies are gas-rich. The Hansen \\& Moore (2006) relation between anisotropy and the slope of the density profile breaks down for gas and DM, but the stars follow an analogous relation with slightly less radial anisotropies for a given density slope. Using kurtosis ($h_4$) to infer anisotropy in ellipticals is dangerous, as $h_4$ is also sensitive to small levels of rotation. The stationary Jeans equation provides accurate masses out to $8\\,R_e$. The discrepancy between the modeling of Romanowsky et al. (2003), indicating a dearth of DM in ellipticals, and the simulations analyzed by Dekel et al. (2005), which match the spectroscopic observations of ellipticals, is partly due to radial anisotropy and to observing oblate ellipticals face-on. However, one of the 15 solutions to the orbit modeling of Romanowsky et al. is found to have an amount and concentration of DM consistent with $\\Lambda$CDM predictions. ",
        "introduction": "The quantity of dark matter lying in the outskirts of luminous elliptical galaxies is a hotly debated topic (see Romanowsky, Napolitano, Stoehr, in these proceedings). There is a wide consensus that, given their flat rotation curves, spiral galaxies must be embedded within dark matter halos, unless one resorts to modifying physics (e.g. MOND, see McGaugh in these proceedings). Moreover, dissipationless cosmological $N$-body simulations lead to structures, whose halos represent most (Hayashi et el. 2004) if not all \\citep{Stoehr06} spiral galaxies.  If elliptical galaxies originate from major mergers of spiral galaxies \\citep{Toomre77,Mamon92,BCF96,SWTK01}, then they too should possess dark matter halos. However, inferring the presence of dark matter halos in ellipticals is difficult, because the velocity dispersions of the usual kinematical tracer, stars, can only be measured out to $2\\,R_e$ (effective radii, containing half the projected light). Moreover, the mass distribution depends on the radial variation of the velocity anisotropy, and one cannot solve for both, unless one assumes no rotation and makes use of the 4th order moment (kurtosis) of the velocity distribution \\citep{vdMF93,Gerhard93,LM03}.  Using planetary nebulae (PNe) as tracers of the dark matter at large radii, \\cite{Romanowsky+03} found low velocity dispersions for their outermost PNe, which after some simple Jeans modeling and more sophisticated orbit modeling led them to conclude to a dearth of dark matter in ordinary elliptical galaxies. This result is not expected in the standard model of formation of structure in the Universe and of galaxies in particular.  This has led us \\citep{Dekel+05} to analyze the final outputs of $N$-body simulations of spirals merging into ellipticals (\\citealp{CPJS04}, Stoehr, in these proceedings). \\citeauthor{Dekel+05} show that the line-of-sight velocity dispersions of their simulated merger remnants are as low as the PNe dispersions measured by \\citeauthor{Romanowsky+03} thus removing an important thorn in the standard model of galaxy formation in a $\\Lambda$CDM Universe. One is led to wonder how could the kinematical analysis of \\citeauthor{Romanowsky+03} lead to a lack of dark matter, while the dynamical analysis of \\citeauthor{Dekel+05} matches the same set of observations with simulations including normal amounts of dark matter.  To explain this discrepancy, it is useful to focus first on some important aspects of kinematical modeling. ",
        "conclusions": ""
    },
    "0601/astro-ph0601035_arXiv.txt": {
        "abstract": "The southern site of the Pierre Auger Observatory is under the construction near Malargue in Argentina and now more than 60\\% of the detectors are completed. The observatory has been collecting data for over 1 year and the cumulative exposure is already similar to that of the largest forerunner experiments. The hybrid technique provides model-independent energy measurements from the Fluorescence Detector to calibrate the Surface Detector. Based on this technique, the first estimation of the energy spectrum above 3 EeV has been presented and is discussed in this paper. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601203_arXiv.txt": {
        "abstract": "We present the analysis of the first 18 months of data obtained with the COSMOSOMAS experiment at the Teide Observatory (Tenerife). Three maps have been obtained at 12.7, 14.7 and 16.3 GHz covering 9000 square degrees each with a resolution of $\\sim 1$ degree and with sensitivities 49, 59 and 115 $\\mu$K per beam respectively. These data in conjuction with the WMAP first year maps have revealed that the Cosmic Microwave Background (CMB) is the dominant astronomical signal at high galatic latitude ($|b|>40^\\circ$) in the three COSMOSOMAS channels with an average amplitude of $29.7\\pm 1.0$ $\\mu$K (68\\% c.l. not including calibration errors). This value is in agreement with the predicted CMB signal in the COSMOSOMAS maps using the   best fit $\\Lambda$-CDM model to the WMAP power spectrum. Cross correlation analysis of the 408~MHz map and the COSMOSOMAS data at high galactic latitudes give values  in the range $17.0-14.4~\\mu$K from 12.7 to 16.3 GHz. Removing detected point sources in this template, reduces the amplitude of the correlated signal to  8-9 $\\mu$K. The mean spectral index of the correlated signal between the 408 MHz desourced and the COSMOSOMAS maps is between   -3.20 and -2.94  at $|b|$$>$40$^\\circ$ which indicates that this signal is  due to synchrotron emission.  Cross-correlation of COSMOSOMAS data with the DIRBE map at 100 $\\mu$m shows the existence of a common signal with  amplitude $7.4\\pm 1.1$, $7.5\\pm 1.1$, and $6.5\\pm 2.3$ $\\mu$K in the 12.7, 14.7 and 16.3 GHz COSMOSOMAS maps at $|b|>$30$^\\circ$.  Using the WMAP data we find this DIRBE correlated signal rises from high to low frequencies flattening below $\\sim 20$ GHz.  At higher galactic latitudes the average amplitude of the correlated signal with the DIRBE maps decreases slightly. The frequency behaviour of the COSMOSOMAS/WMAP correlated signal with DIRBE is not compatible with the expected tendency for thermal dust. A study of the H$\\alpha$ emission maps do not support free-free as a major contributor to that signal. Our results provide evidence of a new galactic foreground with properties compatible with those predicted by the spinning dust models. ",
        "introduction": "In recent years many experiments have measured Cosmic Microwave Background (CMB) anisotropies at angular scales from several degrees to a few arcmin. The analysis of the resulting power spectra has tested the Lambda Cold Dark Matter ($\\Lambda$CDM) model and allowed an accurate determination of cosmological parameters ($e.\\,g.$ Spergel et al. 2003; Rebolo et al.  2004). Detailed knowledge of the galactic foregrounds (synchrotron, free-free and vibrational thermal dust emission) is required to ensure the validity of the results. The joint analysis of different templates of galatic emission with WMAP observations (Bennett et al. 2003; Hinshaw et al. 2003) has allowed a better assesment of the properties of these galactic foregrounds at the frequencies where CMB searches are conducted (see also Banday et al. 2003). Apart from the three classical diffuse components (synchrotron, free-free and vibrational dust  emission), there is increasing evidence of an additional dust correlated foreground (Casasus et al. 2004; de Oliveira-Costa et al. 1997, 1998, 1999, 2004; Finkbeiner et al. 2002; Kogut et al. 1996) which does not follow the expected thermal dust spectrum at  frequencies below $\\nu\\la$40 GHz.  The physical nature of such a component is intriguing (Lagache 2003; Leitch et al. 1997; Lim et al. 1996; Mukherjee et al. 2001; Banday et al. 2003). Although it was initially ascribed to free-free emission, the low level of related H$\\alpha$ emission makes unlikely this explanation. Alternatively, one of the most attractive scenarios to explain the origin of that emission is electric dipole radiation from spinning small dust particles as proposed by Draine \\& Lazarian (1998$a,b$). The key observational prediction of such a model is a turn-over in the spectrum at frequencies 10-20~GHz which are not covered by COBE or WMAP.  The Tenerife experiment (Guti\\'errez et al. 2000) has provided observations at 10 and 15 GHz able to test this prediction ($e.\\,g.$ de Oliveira et al 2004). The COSMOSOMAS (COSMOlogical Structures On Medium Angular Scales) experiment represents a qualitative improvement in terms of angular resolution, sensitivity and sky coverage to study any possible new foreground in the frequency range 10-17~GHz. This experiment, designed and built at the Instituto de Astrof\\'\\i sica de Canarias, is located at an altitude of 2400 m in the Teide Observatory, Tenerife (Spain). It consists of two similar instruments working at frequencies centred at 11 and 15 GHz named COSMO~11 and COSMO~15 respectively.  This paper is dedicated exclusively to the analysis of COSMO~15 data. The COSMO~11 results will be presented in a separate paper (Hildebrandt et al. 2005 in prep.). A detailed technical description of the COSMO~15 instrument has been presented in Gallegos et al. (2001) where we also reported the first results obtained with this experiment. Here, we present additional data taken up to January 2001 improving the sensitivity by a factor two. We use the COSMOSOMAS data in conjuction with the first year WMAP data and several Galactic templates to make a robust estimation of the CMB and the different components of the galactic diffuse emission in the frequency range 12-17 GHz. ",
        "conclusions": "The first instrument of the COSMOSOMAS experiment has provided temperature maps of the sky emission at frequencies 12.7, 14.7 and 16.3 GHz covering $\\sim 9000$ square degrees with sensitivities from 50 to 120$\\mu$K per beam. Known point-like radio sources with fluxes above $\\sim 0.8$ Jy are detected in the three COSMOSOMAS maps. The main results obtained in the cross correlation analysis of these maps  with templates of CMB and galactic foregrounds are:  \\begin{itemize}  \\item The correlated signal between the COSMOSOMAS maps and the highest frequency WMAP\\_V  map has an amplitude of 29.7$\\pm 1.0$ $\\mu$K ($|b|>40^\\circ$). This is in agreement with the expected amplitude of the CMB fluctuations observed through the COSMOSOMAS strategy, using the best fit  $\\Lambda$-CDM model to the WMAP power spectrum.  \\item The amplitude of the COSMOSOMAS-WMAP correlated signal increases as one goes to lower WMAP frequencies, indicating the presence of a foreground signal superimposed over the CMB fluctuations.  \\item We find evidence for a contribution of unresolved radiosources to the COSMOSOMAS maps with an amplitude of $\\sim 19 $ $\\mu$K at 14.7 GHz which we interpret  as the result of a randomly distributed population of unresolved radiosources.  \\item At high galactic latitudes significant correlations for COSMOSOMAS and the 408 MHz and 1420 MHz galactic templates have been found. The values of the synchrotron spectral index are between -3.20 and -2.94 in the  408 MHz-16.3 GHz range.  \\item The correlation analysis between the COSMOSOMAS/WMAP maps, and DIRBE channels at 100 and 240 $\\mu$m indicates also the presence of common signals at high galactic latitudes ($|b|>30^\\circ$).  The amplitude of these signals rises with a very steep index from $\\sim 40$ GHz up to  $\\sim 20$ GHz,  flattening in the range 13-17 GHz with an amplitude of  5-7$\\mu$K. This behaviour is compatible with the predictions of spinning dust models. The small correlation ($\\sim 3$ $\\mu$K) between COSMOSOMAS and the H$\\alpha$ map rules  out free-free emission as the  emission mechanism responsible of the dust correlated signal detected in these regions.  \\end{itemize}"
    },
    "0601/astro-ph0601565_arXiv.txt": {
        "abstract": "We report the results from a deep {\\em HST} NICMOS H-band imaging survey of a carefully selected sample of 33 luminous, late-stage galactic mergers at $z < 0.3$. This program is part of {\\em QUEST}, a comprehensive investigation of the most luminous mergers in the nearby universe, the ultraluminous infrared galaxies (ULIRGs) and the quasars ({\\em QUEST} = {\\em Q}uasar / {\\em U}LIRG {\\em E}volutionary {\\em S}tudy).  Signs of a recent galactic interaction are seen in all of the objects in the {\\em HST} sample, including all 7 IR-excess Palomar-Green (PG) QSOs in the sample.  Unsuspected double nuclei are detected in 5 ULIRGs.  A detailed two-dimensional analysis of the surface brightness distributions in these objects indicates that the great majority (81\\%) of the single-nucleus systems show a prominent early-type morphology.  However, low-surface-brightness exponential disks are detected on large scale in at least 4 of these sources.  The hosts of 'warm' ({\\em IRAS} 25-to-60 $\\mu$m flux ratio, $f_{25}/f_{60} > 0.2$), AGN-like systems are of early type and have less pronounced merger-induced morphological anomalies than the hosts of cool systems with LINER or HII region-like nuclear optical spectral types.  The host sizes and luminosities of the 7 PG~QSOs in our sample are statistically indistinguishable from those of the ULIRG hosts.  In comparison, highly luminous quasars, such as those studied by Dunlop et al. (2003), have hosts which are larger and more luminous. The hosts of ULIRGs and PG~QSOs lie close to the locations of intermediate-size ($\\sim$ 1 -- 2 $L^*$) spheroids in the photometric projection of the fundamental plane of ellipticals, although there is a tendency in our sample for the ULIRGs with small hosts to be brighter than normal spheroids.  Excess emission from a young stellar population in the ULIRG/QSO hosts may be at the origin of this difference. Our results provide support for a possible merger-driven evolutionary connection between cool ULIRGs, warm ULIRGs, and PG~QSOs. However, this sequence may break down at low luminosity since the lowest luminosity PG~QSOs in our sample show distinct disk components which preclude major (1:1 -- 2:1) mergers. The black hole masses derived from the galaxy host luminosities imply sub-Eddington accretion rates for all objects in the sample. ",
        "introduction": "Galaxy merging is a key driving force of galaxy evolution. In hierarchical cold dark matter models of galaxy formation and evolution, merging leads to the formation of elliptical galaxies, triggers major starbursts, and may account for the growth of supermassive black holes and the formation of quasars (e.g., Kauffmann \\& Haehnelt 2000). The importance of mergers increases with redshift (e.g., Zepf \\& Koo 1989; Carlberg, Pritchet, \\& Infante 1994; Neuschaefer et al. 1997; Khochfar \\& Burkert 2001). It is clear that dust-enshrouded starbursts and active galactic nuclei (AGN) play an extremely important role in the high-redshift Universe and are probably the dominant contributors to the far-infrared/submm and X-ray backgrounds, respectively (e.g., Pei, Fall, \\& Hauser 1999; Miyaji, Hasinger, \\& Schmidt 2000).  These luminous, merger-induced starbursts and AGN at high redshift thus provide readily observable signposts for tracing out the main epoch of elliptical galaxy and quasar formation if the above scenario is correct.  In order to assess quantitatively the physics of the merger process and its link to the epoch of elliptical and QSO formation at high redshift we must first understand the details of galaxy merging and its relationship to starbursts and AGN in the local universe. The most violent local mergers and the probable analogs to luminous high-redshift mergers are the ultraluminous infrared galaxies (ULIRGs). ULIRGs are advanced mergers of gas-rich, disk galaxies sampling the entire Toomre merger sequence beyond the first peri-passage (Veilleux, Kim, \\& Sanders 2002; hereafter referred as VKS02). ULIRGs are among the most luminous objects in the local universe, with both their luminosities ($>$ 10$^{12}$ L$_\\odot$ emerging mainly in the far-IR) and space densities similar to those of quasars (e.g., Sanders \\& Mirabel 1996). The near-infrared light distributions in many ULIRGs appear to fit a $R^{1/4}$ law (Scoville et al. 2000; VKS02).  ULIRGs have a large molecular gas concentration in their central kpc regions (e.g., Downes \\& Solomon 1998) with densities comparable to stellar densities in ellipticals. These large central gas concentrations (and stars efficiently forming from them) may be the key ingredient for overcoming the fundamental phase space density constraints that would otherwise prevent formation of dense ellipticals from much lower density disk systems (Gunn 1987; Hernquist, Spergel, \\& Heyl 1993). Kormendy \\& Sanders (1992) have proposed that ULIRGs evolve into ellipticals through merger induced dissipative collapse.  In this scenario, these mergers first go through a luminous starburst phase, followed by a dust-enshrouded AGN phase, and finally evolve into optically bright, `naked' QSOs once they either consume or shed their shells of gas and dust (Sanders et al. 1988a).  Gradual changes in the far-infrared spectral energy distributions between `cool' ULIRGs ({\\em IRAS} 25-to-60 $\\mu$m flux ratio, $f_{25}/f_{60} < 0.2$), `warm' ULIRGs, and QSOs (Sanders et al. 1988b; Haas et al. 2003) bring qualitative support to an evolutionary connection between these various classes of objects, but key elements remain to be tested. In a pilot study of a dozen ULIRGs observed with Keck and VLT, Genzel et al. (2001) and Tacconi et al. (2002) have found that ULIRGs resemble intermediate mass ellipticals/lenticulars with moderate rotation, in their velocity dispersion distribution, their location in the fundamental plane (e.g., Djorgovski \\& Davis 1987; Dressler et al. 1987; Kormendy \\& Djorgovski 1989) and their distribution of the ratio of rotation/velocity dispersion [v$_{\\rm rot}$ sin(i)/$\\sigma$]. These preliminary results therefore suggest that ULIRGs form moderate mass ($m^\\ast \\sim 10^{11}$ M$_\\odot$), but not giant (5 -- 10 $\\times$ 10$^{11}$ M$_\\odot$) ellipticals. A comparison between these ULIRGs and the sample of luminous radio-quiet QSOs from Dunlop et al. (2003) indicates that the ULIRGs are offset from the location of the hosts of these QSOs in the photometric projection of the fundamental plane. The latter fall near the locale of giant ellipticals on the fundamental plane. The black hole (BH) masses inferred for the ULIRGs from the host dynamical masses and the local BH mass to bulge velocity dispersion relationship (e.g., Gebhardt et al. 2000) are akin to those of local Seyfert galaxies. From this perspective, ULIRGs and QSOs have comparable luminosities because the ULIRGs are much more efficiently forming stars and feeding their BHs at the epoch they are observed.  This pilot study (Genzel et al. 2001; Tacconi et al. 2002) has raised important questions towards understanding the evolution of ULIRG mergers and their relation to QSOs, but is limited by the small number of galaxies. For example, giant ellipticals constitute only $\\sim$ 10 -- 15\\% of all ellipticals with mass $>$ 10$^{11}$ M$_\\odot$ -- a similar fraction of the most massive ULIRG hosts might have been missed in our small pilot sample. To better understand the evolutionary process of ULIRGs and their relation to elliptical galaxies and quasars, we are currently conducting a comprehensive study of a large sample of ULIRGs and QSOs which samples the full range of properties (L$_{\\rm IR}$, merger state, and AGN/starburst fraction). This study is called {\\em QUEST} (for {\\em Q}uasar/{\\em U}LIRG {\\em E}volutionary {\\em ST}udy).  Our program relies on {\\em HST} NICMOS imaging and VLT/Keck near-infrared spectroscopy to determine the structure and dynamics of ULIRG and QSO mergers, and on high-S/N {\\em Spitzer Space Telescope} ({\\em SST}) mid-infrared spectroscopy to quantify the energy production mechanisms in these objects, and to study their obscuration and physical conditions along the merger sequence.  The purpose of the present paper is to report the results from the {\\em HST} NICMOS imaging component of {\\em QUEST}. Part of the spectroscopy has been presented in Dasyra et al. (2005). In \\S 2, we describe the {\\em HST} sample. Next we discuss the observational strategy of our program and the methods used to obtain, reduce, and analyze the data (\\S 3, \\S 4, and \\S 5, respectively). The results are presented in \\S 6 and discussed in \\S 7. The main conclusions are summarized in \\S 8.  Throughout this paper, we adopt $H_0$ = 75 km s$^{-1}$ Mpc$^{-1}$ and $q_0$ = 0 to be consistent with previous papers in this series.  Given the proximity ($z < 0.27$) of all of our objects, adoption of the WMAP cosmological parameters ($\\Omega_M = 0.27$, $\\Omega_\\Lambda$ = 0.73, H$_0$ = 71 km s$^{-1}$ Mpc$^{-1}$) would not affect any of our conclusions. ",
        "conclusions": "Our detailed two-dimensional analysis of deep {\\em HST} NICMOS images of 26 highly nucleated 1-Jy ULIRGs and 7 PG~QSOs indicates that an important fraction of them (81\\% of the single-nucleus systems) present early-type morphology well fit by a de Vaucouleurs-like surface brightness distribution at their centers but with significant merger-induced morphological anomalies on large scale. This provides support to the picture where ULIRGs are spheroids in formation. Their positions relative to the photometric projection of the fundamental plane for normal (non-active) spheroids are similar to those of intermediate-size ($\\sim$ 1 -- 2 $L^*$) ellipticals or lenticulars. Indeed, in a few cases, our data reveal a faint disk on large scale. However, the spheroidal components of small ULIRGs tend to be brighter than inactive spheroids of the same size. Many of these results are consistent with the kinematic data of Genzel et al. (2001) and Tacconi et al. (2002).  The host sizes, magnitudes, and surface brightness distributions of the 7 PG~QSOs in our sample are statistically identical to those of the ULIRGs. They all display significant tidal features similar to those seen in Seyfert ULIRGs, but weaker than those in H~II and LINER ULIRGs. These results bring further support to the suggestion of a merger-induced evolutionary connection between ULIRGs and PG~QSOs. Earlier studies that found a poor match between the host properties of ULIRGs and those of quasars used a sample of quasars which were significantly more luminous than the ULIRGs in the 1-Jy sample. However, the evolutionary sequence may break down at low luminosities: the disk components found in the lowest luminosity QSOs in our sample cannot be easily explain in the context of major mergers.  Finally, we add a note of caution when interpreting these results. The present conclusions are based purely on photometric measurements which are subject to uncertainties from dust extinction and young stellar populations. The excess emission from a young circumnuclear starburst may help explain why ULIRGs with small hosts are brighter than inactive spheroids. We are in the process of studying the kinematic properties of the objects in this sample to test our conclusions.  Mid-infrared spectroscopic data from our on-going Cycle 1 {\\em SST} program will also complement this data set by providing quantitative measurements of the evolution of the energy source --- starburst versus AGN --- along the merger sequence. Our {\\em HST} data suggest that AGN-dominated ULIRGs do not require super-Eddington accretion rates.  \\vskip 0.1in"
    },
    "0601/astro-ph0601423_arXiv.txt": {
        "abstract": "{We investigate the thermal and kinematic configuration of a sunspot penumbra using very high spectral and spatial resolution intensity profiles of the non-magnetic \\ion{Fe}{i} 557.6 nm line. The dataset was acquired with the 2D solar spectrometer TESOS. The profiles are inverted using a one-component model atmosphere with gradients of the physical quantities. From this inversion we obtain the stratification with depth of temperature, line-of-sight velocity, and microturbulence across the penumbra. Our results suggest that the physical mechanism(s) responsible for the penumbral filaments operate preferentially in the lower photosphere. We confirm the existence of a thermal asymmetry between the center and limb-side penumbra, the former being hotter by 100-150 K on average. We also investigate the nature of the bright ring that appears in the inner penumbra when sunspots are observed in the wing of spectral lines. It is suggested that the bright ring does not reflect a temperature enhancement in the mid photospheric layers. The line-of-sight velocities retrieved from the inversion are used to determine the flow speed and flow angle at different heights in the photosphere. Both the flow speed and flow angle increase with optical depth and radial distance.  Downflows are detected in the mid and outer penumbra, but only in deep layers ($\\log \\tau_{500} \\leq -1.4$). We demonstrate that the velocity stratifications retrieved from the inversion are consistent with the idea of penumbral flux tubes channeling the Evershed flow. Finally, we show that larger Evershed flows are associated with brighter continuum intensities in the inner center-side penumbra. Dark structures, however, are also associated with significant Evershed flows. This leads us to suggest that the bright and dark filaments seen at 0\\farcs5 resolution are not individual flow channels, but a collection of them. Our analysis highlights the importance of very high spatial resolution spectroscopic and spectropolarimetric measurements for a better understanding of sunspot penumbrae. ",
        "introduction": "The fine structure of the penumbra is intimately connected to the Evershed flow, the most conspicuous dynamical phenomenon observed in sunspots (e.g., Solanki 2003; Thomas \\& Weiss 2004). Therefore, a proper observational characterization of this fine structure is essential to understand the Evershed flow and, more generally, the nature of the penumbra itself.   Studying the fine structure of the penumbra is difficult because of the very small scales involved. Recent imaging observations with the Swedish 1-m Solar Telescope (Scharmer et al.\\ 2003; Rouppe van der Voort et al.\\ 2004) have revealed that even at a resolution of 0\\farcs1 some of the filaments that form the penumbra may be unresolved. That is, different structures (each having different properties) may coexist in the resolution element, and we are not able to separate them. A proper characterization of the penumbra calls not only for diffraction-limited imaging observations, but also for more quantitative spectroscopic and spectropolarimetric analyses at the highest resolution possible.  By interpreting the shape of spectral lines we can derive the thermal, magnetic, and kinematic configuration of the penumbra via the Zeeman and Doppler effects. Unfortunately, current spectroscopic observations do not attain resolutions better than $\\sim 0\\farcs5$, and this limit is often reached at the cost of poor spectral resolution. The situation is even worse in the case of spectropolarimetric measurements: the best angular resolution of existing solar polarimeters is about 1\\arcsec.  This paper is the third of a series devoted to the analysis of 2D spectroscopic observations of a sunspot. The data were acquired with the TElecentric SOlar Spectrometer (TESOS; Kentischer et al.\\ 1998; Tritschler et al.\\ 2002) and the Kiepenheuer Adaptive Optics System (KAOS; Soltau et al.\\ 2002; von der L\\\"uhe et al.\\ 2003) at the German Vacuum Tower Telescope of Observatorio del Teide (Tenerife, Spain). These observations combine high spectral ($\\lambda/\\Delta \\lambda \\sim 250\\,000$) and spatial (0\\farcs5) resolution (Tritschler et al.\\ 2004; hereafter referred to as Paper I). In Paper II of this series (Schlichenmaier et al.\\ 2004), we investigated the properties of the Evershed flow through a bisector analysis of the observed intensity profiles.  Here we perform a full inversion of the same data set in order to study the thermal and kinematic configuration of the penumbra. Our analysis allows us to explain a number of findings reported in Paper~I and to remove some of the uncertainties associated with the methods employed in Paper~II.  At the same time, we gather new information on the thermal properties of the penumbra. This is particularly important because very few thermal studies of sunspot penumbrae have been published to date (del Toro Iniesta et al.\\ 1994; Balasubramaniam 2002; Rouppe van der Voort 2002; S\\'anchez Cuberes et al.\\ 2005).  A brief account of the observations and details of the inversion procedure are given in Sect.~\\ref{obs}. In Sect.~\\ref{results} we discuss the maps of physical quantities inferred from the inversion. The thermal and kinematic configuration of the penumbra is described in detail in Sects.~\\ref{thermal} and \\ref{kinematic}. Some implications of the results are presented in Sect.~\\ref{discusion}, and a summary of our findings is given in Sect.~\\ref{summary}. ",
        "conclusions": "\\label{discusion}  \\subsection{Origin of enhanced line-wing intensities in the inner penumbra} From an analysis of the Stokes profiles of the \\ion{Fe}{i} 1564.8 nm line, Bellot Rubio (2003) found that the intensity emerging from the inner penumbra is higher than that emerging from the outer penumbra at wavelengths close to, but not exactly at, the line center. That is, sunspots exhibit a bright ring in the inner penumbra when observed in the wing of spectral lines. The bright ring can be interpreted as a temperature enhancement in the mid-photospheric layers (as it is detected in the wings of the line but not in the continuum or line core).  Our high spatial resolution observations of NOAA 10019 confirm this behavior (Fig.\\ 8 of Paper~I) and at the same time add new information on it: the intensity enhancement occurs in the form of radial, narrow fibrils very similar to the bright penumbral filaments seen in continuum images, but with much higher contrast (Fig.\\ 3 of Paper~I). Moreover, the enhancement is observed preferentially in the inner center-side penumbra blueward of the line core, and in the inner limb-side penumbra redward of the line core. The temperature stratifications of Fig.~\\ref{t_rad} indicate that the inner penumbra is always cooler than the outer penumbra at all heights in the atmosphere, so the enhanced brightness cannot be due to a temperature effect. Heating of the mid photosphere was the mechanism originally proposed by Bellot Rubio (2003). It is also the scenario favored by S\\'anchez Cuberes et al.\\ (2005).  Here we argue that the bright ring is actually produced by (a) the small Doppler shifts and line asymmetries observed in the inner penumbra, together with (b) the increase of the line width toward the outer sunspot boundary. The left panel of Fig.~\\ref{brightring} shows an intensity map of the spot at $\\Delta \\lambda = -6.4$ pm from line center (filtergram \\#42 in Fig.~3 of Paper~I). The enhanced brightness of the inner center-side penumbra is apparent in this image. The intensity profiles emerging from the two spatial positions marked with crosses are displayed in the right panel of Fig.~\\ref{brightring}. As can be seen, these profiles possess very different properties: in the outer penumbra (solid line) they are broader and have a stronger asymmetry than in the inner penumbra (dashed line). The bright ring results from a comparison of the intensities of pixels belonging to the inner and outer penumbra at fixed wavelengths (dotted vertical lines). Consider, for example, wavelengths in the blue wing (leftmost vertical line). At this wavelength, the inner penumbra is clearly brighter than the outer penumbra. In the red wing (rightmost vertical line), the inner penumbra would still be brighter, but with a much reduced contrast. This is consistent with the observations described in Paper~I. Similar arguments apply to profiles emerging from the limb-side penumbra: because of the smaller line widths and asymmetries in the inner penumbra, the intensity at fixed wavelengths is enhanced relative to the outer penumbra, with higher contrasts redward of the line core.    \\begin{figure} \\begin{center} \\resizebox{0.49\\hsize}{!}{\\includegraphics{./linewing_blue_map.ps}} \\resizebox{0.49\\hsize}{!}{\\includegraphics{./brightring_centerside.ps}} \\end{center} \\caption{ {\\em Left:} Intensity map of the spot at $\\Delta \\lambda = -6.4$ pm from line center. Two pixels belonging to the inner and outer center-side penumbra are marked with crosses. {\\em Right:} Best-fit intensity profiles returned by the inversion at the spatial positions marked with crosses in the left panel. Solid: outer penumbra. Dashed: inner penumbra. The vertical dotted lines mark fixed wavelength positions in the blue wing, the line core, and the red wing (from left to right). The blue and red wing positions correspond to filtergrams \\#42 and \\#58 of Fig.~3 in Paper~I. \\label{brightring} } \\end{figure}   Thus, the bright ring can be explained satisfactorily in terms of the radial variation of the microturbulence and the Evershed flow. Microturbulence is an important parameter: if it were the same everywhere across the spot, redward of the line core the inner penumbra would appear {\\em darker} than the outer penumbra on the center side, contradicting the observations. To further support the idea that the bright ring is not due to temperature enhancements, we have carried out the same inversions by forcing the LOS velocity to be constant with height (i.e., we do not allow for line asymmetries). In addition, we set the microturbulence to zero. The thermal stratification resulting from these inversions is shown in Fig.~\\ref{maps_vel1nodo} at $\\log \\tau_{500} = 0$ and $-1$. While the temperature in the deep layers is essentially unchanged, the map of temperature at $\\log \\tau_{500} = -1$ shows a nice bright ring in the inner penumbra. From this we conclude that the only way to explain the enhanced line-wing intensities of the inner penumbra when no line asymmetries are considered is by an increase of the temperature in the mid photosphere.   \\begin{figure} \\begin{center} \\resizebox{0.49\\hsize}{!}{\\includegraphics{./templogtau0_map_vel1nodo.ps}} \\resizebox{0.49\\hsize}{!}{\\includegraphics{./templogtau-1_map_vel1nodo.ps}} \\vspace*{-.5em} \\end{center} \\caption{Maps of temperature at $\\log \\tau_{500}= 0$ {\\em (left)} and $\\log \\tau_{500}= -1$ {\\em (right)} resulting from an inversion of the observed profiles in which only constant velocities along the LOS are permitted. \\label{maps_vel1nodo} } \\end{figure}  Using flowless maps, Balasubramaniam (2002) detected bright structures in the inner penumbra at wavelengths close to line center (his Fig.~8). Transforming observed intensities into radiation temperatures, he concluded that these structures are hotter than their surroundings. In our opinion, however, they more likely trace regions of reduced Evershed flows and small-scale turbulence.  \\subsection{Interpretation of the LOS velocity stratifications} \\label{signreversal}  In Paper~II we demonstrated that the uncombed model proposed by Solanki \\& Montavon (1993) is able to explain the bisector shapes observed in the penumbra of NOAA 10019. In an uncombed penumbra, the Evershed flow occurs along flux tubes whose magnetic fields are more horizontal than that of their surroundings. Here we use this model to interpret the LOS velocity stratifications displayed in Fig.~\\ref{averageLOS1}.  As pointed out in Sect.\\ \\ref{inversion}, it is the magnitude and slope of the inferred velocity stratification which inform us about the flow geometry (flow speed and flow angle) and height of the Evershed channels. Figure~5 of Paper~II shows that, in the very outer penumbra, two flow channels stacked along the LOS and with different inclination angles are necessary to understand the kinks and reversals of the observed bisectors. In the inner and mid penumbra the bisectors are rather linear, and can be explained in terms of a single, deep-lying flow channel.  Based on these results, in the inner penumbra ($r=0.6R$) we consider a flow channel located in deep layers, between $\\log \\tau_{500}= -0.7$ and 0.  The flow has a speed of 8 km~s$^{-1}$ and an inclination of $87^\\circ$ with respect to the vertical. In the outer penumbra ($r=0.97R$), we assume two flow channels located in deep and mid photospheric layers. The deep-lying channel extends from $\\log \\tau_{500}= -0.2$ to 0.5. It has a flow speed of 8 km~s$^{-1}$ and an inclination of 135$^\\circ$. The upper channel extends from $\\log \\tau_{500}= -1.8$ to $-1.1$ and is assumed to have a flow speed of 7~km~s$^{-1}$ and an inclination of 85$^\\circ$. This flow configuration is similar to the one deduced in Paper II, except for the fact that our higher-lying channel is slightly inclined upward instead of being horizontal. The LOS velocities resulting from the two flux-tube geometries at the heliocentric angle of our observations ($\\theta = 23^\\circ$) are represented by solid lines in Fig.~\\ref{averageLOS1_bis}, for the center and limb-side penumbra.  To investigate whether or not these flow configurations can explain the LOS velocities retrieved from the inversion, we have computed synthetic profiles under the assumption that the resolution element contains two different atmospheres. The temperature and pressure stratifications of both atmospheres are taken to be the mean stratifications deduced from the inversion of NOAA 10019. The LOS velocity is set to zero in the first component, whereas for the second component we use the stratifications of Fig.~\\ref{averageLOS1_bis}. In the inner penumbra we adopt a filling factor of 0.2, i.e., only 20\\% of the resolution element is occupied by flows. In the outer penumbra the filling factor is 0.5, as in Paper II.  The resulting profiles are convolved with the instrumental profile of TESOS and a macroturbulent velocity of 2~km~s$^{-1}$ (inner penumbra) or 1~km~s$^{-1}$ (outer penumbra). Finally, they are inverted under the same conditions as the real data (two nodes for velocity, temperature and microturbulence, and a height-independent macroturbulence). Not unexpectedly, the simulated profiles show the same bisector shapes and bisector variations across the penumbra as the real observations.   \\begin{figure} \\begin{center} \\resizebox{0.47\\hsize}{!}{\\includegraphics{./inner_limb_inv.ps}} \\resizebox{0.47\\hsize}{!}{\\includegraphics{./inner_center_inv.ps}} \\\\ \\resizebox{0.47\\hsize}{!}{\\includegraphics{./outer_limb_inv.ps}} \\resizebox{0.47\\hsize}{!}{\\includegraphics{./outer_center_inv.ps}} \\end{center} \\caption{LOS velocity stratifications due to flux tubes along the LOS (solid lines) that would result in {\\em linear} velocity stratifications (dashed lines) when two nodes are used to retrieve the actual velocity stratification. See text for details on the flux-tube geometries adopted in the inner and outer penumbra. \\label{averageLOS1_bis}} \\end{figure}   \\begin{figure*} \\begin{center} \\resizebox{.68\\hsize}{!}{\\includegraphics{./corrIc_v_inner_biermann.ps}} \\end{center} \\caption{Continuum intensity fluctuations {\\em (solid)} and LOS velocity at $\\log \\tau_{500} = 0$ {\\em (dashed)} along an azimuthal path crossing the inner center-side penumbra as indicated on the left image. Negative velocities represent blueshifts. The vertical dotted lines mark positions of local maxima of the continuum intensity. \\label{corrI_v} } \\end{figure*}  The {\\em linear} velocity stratifications retrieved from the inversion of the simulated profiles are represented by dashed lines in Fig.~\\ref{averageLOS1_bis}.  The first thing to notice is that the inferred stratifications change sign in the mid photosphere (around $\\log \\tau_{500} = -2$) both on the center and limb sides, even though there are no structures in high layers that could be responsible for such velocity reversals. Actually, the sign reversal is the result of extrapolating the jump(s) of velocity sensed by the LOS to the upper layers, where the response of the line to plasma motions is small. This extrapolation effect was first described by Mart\\'{\\i}nez Pillet (2000) in connection with inversions of Stokes profiles of visible lines that do not explicitly account for the existence of penumbral flux tubes. The same effect has been observed by Mathew et al.\\ (2003) in their one-component inversions of near-infrared lines.   The second thing to note is that the azimuthally averaged LOS velocities inferred from the inversion of NOAA 10019 (Fig.~\\ref{averageLOS1}) are in excellent accord with the {\\em linear} velocity stratifications resulting from the flux tube configurations displayed in Fig.~\\ref{averageLOS1_bis}. Specific features that are well reproduced include the following: \\begin{itemize} \\item On the limb side, the slope of the stratifications is larger in the outer penumbra as compared with the inner penumbra, whereas on the center side it is slightly larger in the inner penumbra. \\item The slope of the stratifications in the inner penumbra is larger on the center side. \\item The LOS velocities corresponding to the outer center-side penumbra are negative all along the atmosphere. \\end{itemize} From this we conclude that the presence of flux tubes in the penumbra can explain the observed LOS velocity stratifications. In particular, the seemingly weird run with depth of the LOS velocity in the center-side penumbra at normalized radial distances larger than 0.9 (with zero or small redshifts at $\\log \\tau_{500} = 0$ and blueshifts increasing with height, see Fig.~\\ref{averageLOS1}) can be understood in terms of two flux tubes stacked along the LOS. The tube lying deeper contributes with positive velocity (i.e., a {\\em redshift}). This redshift is necessary to explain the bisector reversals described in Paper II. In order to produce a redshift in the center-side penumbra, the deep-lying flux tube must be inclined downward with respect to the horizontal by more than 23$^{\\rm o}$ (the heliocentric angle of the spot), otherwise it would show negative velocities and the slope of the resulting LOS velocity stratification would be opposite to what we observe. Flux tubes returning to the solar surface have previously been inferred by, among others, Westendorp Plaza et al.\\ (1997, 2001), Schlichenmaier \\& Schmidt (2000), Bellot Rubio et al.\\ (2003, 2004) and Borrero et al.\\ (2004, 2005).  The tube located higher in the atmosphere may be horizontal (as in Paper II) or slightly inclined upwards (less than 23$^\\circ$ from the horizontal). The later configuration would still produce redshifts on the limb side and blueshifts on the center side, exactly what is required to explain the velocity stratifications observed in the penumbra. Interestingly, such high-lying tubes do not seem to be necessary in the inner penumbra, only in the mid and outer penumbra.  It is tempting to associate them with those described by Westendorp Plaza et al.\\ (2001) and Bellot Rubio et al.\\ (2004). These authors found indications of a new family of penumbral tubes in the mid penumbra and beyond. Spectropolarimetric measurements suggest that such tubes are more vertically oriented than the deeper ones, carrying a fraction of the Evershed flow to the upper photosphere.   \\subsection{Do we see individual flow channels?} Continuum intensity fluctuations and LOS velocities are very well correlated in the inner penumbra. This is shown in Fig.~\\ref{corrI_v}, where we plot normalized continuum intensity variations and LOS velocities at $\\log \\tau_{500} = 0$ along an azimuthal path crossing 8 distinct continuum filaments of the inner center-side penumbra (we do not consider the limb-side part of the azimuthal path because no clear filaments are observed on the limb side). The curves depicted in Fig.~\\ref{corrI_v} reveal that {\\em larger continuum intensities are associated with larger LOS velocities}, i.e., the Evershed flow is more intense in the brighter structures, at least in the inner center-side penumbra. This agrees with the conclusions of Bello Gonz\\'alez et al.\\ (2005) and Schlichenmaier et al.\\ (2005).  However, the most important fact demonstrated by Fig.~\\ref{corrI_v} is that also the dark structures show large LOS velocities. Stray light from the bright filaments may change the line profile in the dark structures, but only to some extent (cf.\\ Sect. 3.2 in Paper I).  Unless the stray light contamination is unreasonably large, the velocities found in the dark filaments must be real. This means that we see the Evershed flow everywhere: even at a resolution of 0\\farcs5, there is almost no position in the penumbra where the velocity drops to zero (except where it must be zero for geometrical reasons). In our opinion, this well established fact (see also Hirzberger \\& Kneer 2001 and Rouppe van der Voort 2002) has not been emphasized adequately in the past.  Since Evershed flows are ubiquitous, it is tempting to think that the bright and dark structures observed in continuum images are individual flow channels (i.e., isolated flux tubes).  But this interpretation is problematic. For example, one would have to explain why the flow channels possess different temperatures and velocities at the same radial distance and, more importantly, why there exists such a regular azimuthal pattern of temperature enhancements/coolings and larger/smaller plasma flows. An alternative explanation is that the filaments seen in continuum images are not individual flow channels, but a collection of them. If the filling factor of the tubes changes regularly in the azimuthal direction, i.e., if the number of flow channels present in the resolution element changes azimuthally, then one would observe larger LOS velocities and higher temperatures at those positions where there are more tubes (provided, of course, that the tubes are hot in the inner penumbra, as suggested by the simulations of Schlichenmaier et al.\\ 1998 and the uncombed inversions of Borrero et al.\\ 2005). The important idea here is that the intrinsic flow speed and temperature of the flux tubes would not change azimuthally at a fixed radial distance from sunspot's center, only the filling factor would change. We believe this is a natural explanation for the correlation displayed in Fig.~\\ref{corrI_v}.  Independent support for this idea has been gathered from the analysis of full Stokes profiles of near-infrared lines (Bellot Rubio 2004; Bellot Rubio et al.\\ 2004).  Our results strongly suggest that the filaments observed in Doppler maps at a resolution of about 0\\farcs5 are still unresolved. That is, they are not individual flow channels. For this reason, in Paper~II we referred to them as {\\em flow filaments}, to clearly distinguish them from the actual, smaller flow channels."
    },
    "0601/astro-ph0601109_arXiv.txt": {
        "abstract": "Irradiance variability has been monitored from space for more than two decades. Even if data are coming from different sources, it is well established that a temporal variability exists which can be set to as $\\approx$ 0.1\\%, in phase with the solar cycle. Today, one of the best explanation for such an irradiance variability is provided by the evolution of the solar surface magnetic fields. But if some 90 to 95\\% can be reproduced, what would be the origin of the 10 to 5\\% left? Non magnetic effects are conceivable. In this paper we will consider temporal variations of the diameter of the Sun as a possible contributor for the remaining part. Such an approach imposes strong constraints on the solar radius variability. We will show that over a solar cycle, variations of no more than 20 mas of amplitude can be considered. Such a variability --far from what is reported by observers conducting measurements by means of ground-based solar astrolabes-- may explain a little part of the irradiance changes not explained by magnetic features. Further requirements are needed that may help to reach a conclusion. Dedicated space missions are necessary (for example PICARD, GOLF-NG or SDO, scheduled for a launch around 2008); it is also proposed to reactivate SDS flights for such a purpose. ",
        "introduction": "Total Solar Irradiance (TSI) has been measured from various spacecrafts platforms (and some rockets) since November 16, 1978. TSI refers to the total electromagnetic energy received by a unit area per unit of time, at the mean Sun-Earth distance (1 AU). Historically, TSI has been called the {\\it ``solar constant\"} by A. Pouillet in 1837. Since then, measurements have been made by a long series of observers as late as 1942 when Abbot set the conventional value to 1.98 $\\pm$ 0.05 cal$/$min$/$cm$^2$ (i.e. 1381 $\\pm$ 35 W$/$m$^2$, close to the modern one, 1366 $\\pm$ 1.3 W$/$m$^2$ (de Toma et al., 2004). At that time, it was recognized, notably by S.P. Langley, that the solar output might not be constant, by the order of $\\pm$ 0.05$\\%$. However, until to the 70's, a great number of investigators favored the notion of an unchanging Sun. Evidence suggesting variation of the solar output was thus very controversial. The difficulties in measuring the solar constant through the ever-changing Earth atmosphere precluded a resolution of this question prior to the space age.  Today there is irrefutable evidence for variations of the total solar irradiance. Measurements made during 26 years (two and an half solar cycles) show a variability on different time-scales, ranging from minutes up to decades, and likely to last even longer. Despite of the fact that collected data come from different instruments aboard different spacecrafts, it has been possible to construct a homogeneous composite TSI time series, filling the different gaps and adjusted to an initial reference scale. Fig. \\ref{pap_eps1} shows such a composite, details of which can be found in Pap (2003). The most prominent discovery of these space-based TSI measurements is a 0.1$\\%$ variability over the solar cycle, values being higher during phases of maximum activity  (Pap, 2003 and other references therein). Since the data do not come from a unique instrument, it is not yet clear if the minimum level of TSI varies from one cycle to another one.  The physical causes of these variations still remain a subject of debate. If we know that surface magnetism might be the primary cause, we are not sure that other factors would not have a role to play. Among them, one possible could be temporal variations of the radius of the Sun, others being mainly effective temperature and solar oscillations frequencies. We will deal here merely with radius changes. However even if the Sun's radius is the parameter implied in standard solar models, we prefer to speak in terms of the Sun's shape. In any case, the physical link between this geometrical shape and radiance is more direct than the simple correspondence between radius and radiance. We are from the start very careful on this question, and we will later explain why.  \\begin{figure}[t] \\begin{center} \\includegraphics[width=.7\\textwidth]{pap_oleron_nb} \\end{center} \\caption[]{The various total irradiance time series are presented on the upper panel, the composite total solar irradiance is shown on the lower panel (updated from Fr\\\"ohlich, 2000; courtesy J. Pap, 2003). } \\label{pap_eps1} \\end{figure}  Curiously, concerning the current situation of a possible temporal variation of the radius, the situation is the same as that which prevailed for the solar constant at the time of Abbot. The lack of consensus on the subject is mainly due to measurements made by means of ground-based instruments suggesting a temporal variation, but contaminated by atmospheric artefacts. Several space missions to measure simultaneously the TSI and the shape of the Sun have been proposed, but so far, none has been launched yet.  Understanding and modeling TSI is important for at least three main reasons. First, for the reconstruction of the TSI before the satellite era. Second, for predictions, and third, linked to the others, for the radiative forcing on the Earth's troposphere. A successful model of TSI is therefore of particular interest both for {\\it space weather} (subject of this volume) and for climatic impacts.  Based on the outlines mentioned above, this paper is divided into three sections. In the first one, we will briefly recall our present understanding on TSI changes, emphasizing observed differences between irradiance and magnetic field variations. In the second section, we will describe our concept of the Sun's shape. Then we will show how observed irradiance changes put strong constraints on temporal radius variations. We conclude by recommanding  a satellite mission dedicated to measuring in real time the solar asphericity-luminosity parameter $w$, as described by Sofia (1979) or Lefebvre and Rozelot (2004). ",
        "conclusions": "In this study we have shown that part of the TSI variations could be explained by variations in the solar radius (more precisely in the solar shape changes). Even if a modulation of 200 mas may explain the quasi-totality of the irradiance trend, such a large variability is already excluded. However, an amplitude of a maximum (peak to peak) 20 mas over one solar cycle explains 10\\% of this trend, and 10 mas, $\\approx$ 5\\%. It is thus suggested that so far non explained variations of TSI could draw their origins from non-magnetic mechanisms. Such an approach has also been studied by Callebaut, Makarov and Tlatov (2002) and Li et al. (2003). The first authors developed a self-consistent model to explain cyclic variations of solar radius variations: the decrease in luminosity during half a solar cycle, balanced by gravitational energy, yields radius variations of $\\approx$ 11 mas. The latter authors constructed models of the solar structure in which cyclic radius variations have been taken into account. At last, recent space-based data obtained with SOHO/MIDI by Kuhn et al. (2004) place limits on solar radius variations, $dR < 7$ mas, and $w < 0.1$, still leaving room for such changes.  While a considerable amount of data has been gathered on TSI variations since 1978, we still significant lack of information on the possible temporal changes of the solar diameter (and of the Sun's shape). In-flight missions are necessary, particularly dedicated missions which could measure in real time and simultaneously, irradiance and the  asphericity-luminosity parameter $w$, (for instance aboard PICARD, GOLF-NG or the next SOHO satellite, called SDO). Until such missions take place, though not scheduled before 2008, we propose to reactivate annual SDS balloon flights based upon a joint European-Aerican program (Rozelot and Sofia, 2004).  Finally, we would like to emphasize the need for a better understanding of TSI variability to discern its role in space-weather and its impact on the terrestrial climate.  \\vspace{0.5cm} \\noindent \\Large{\\bf{Acknowledgements}}\\\\  \\noindent \\small{The authors wish to thank P. Scherrer for useful suggestions on the manuscript.}"
    },
    "0601/astro-ph0601615_arXiv.txt": {
        "abstract": "We present far-ultraviolet (FUV) spectral images, measured at \\ion{C}{4} $\\lambda$1550, \\ion{He}{2} $\\lambda$1640, \\ion{Si}{4}+\\ion{O}{4}] $\\lambda$1400, and \\ion{O}{3}] $\\lambda$1664, of the entire Cygnus Loop, observed with the Spectroscopy of Plasma Evolution from Astrophysical Radiation ({\\em SPEAR}) instrument, also known as {\\em FIMS}. The spatial distribution of FUV emission generally corresponds with a limb-brightened shell, and is similar to optical, radio and X-ray images. The features found in the present work include a ``carrot'', diffuse interior, and breakout features, which have not been seen in previous FUV studies. Shock velocities of 140--160 km s$^{-1}$ is found from a line ratio of \\ion{O}{4}] to \\ion{O}{3}], which is insensitive not only to resonance scattering but also to elemental abundance. The estimated velocity indicates that the fast shocks are widespread across the remnant. By comparing various line ratios with steady-state shock models, it is also shown that the resonance scattering is widespread. ",
        "introduction": "The Cygnus Loop is one of the most well-studied supernova remnants (SNRs) in our galaxy because of its large apparent angular size, its high surface brightness, and its low reddening. It is generally considered the prototypical ``middle-aged'' SNR. Global features of the Cygnus Loop have been studied in great detail at optical, X-ray, radio, and infrared wavelengths \\citep{Levenson1998, AL1999, LRB1997, ADL1992}. X-rays are emitted from hot (temperature $T\\sim10^6$ K) gas heated by shocks with velocities $v_s \\sim400$ km s$^{-1}$, while optical emission arises from slower shocks in which the postshock region cools to $T\\sim10^4$ K. Observations of far-ultraviolet (FUV) emission from intermediate temperatures gas, mostly limited to small spatial regions, have played an important role for the reliable estimation of shock speeds and elemental abundances. Early studies concentrated on bright optical filaments and compared FUV spectra with steady-flow shock models \\citep{Benvenuti1980, Raymond1980, Raymond1981, Raymond2001, Blair1991}. Subsequent investigations, focused on fainter filaments \\citep{Raymond1983,Long1992,Sankrit2000,Sankrit2002}, have shown the complexities involved including the effects of geometry, shock incompleteness, resonance line scattering, dust grain destruction.  FUV observations of the entire Cygnus Loop are of great interest because they provide a global understanding of the interaction between the X-ray and optical emitting components. \\cite{CP1976} obtained images from a portion of the Cygnus Loop in two broad bands (1250--1600 and 1050--1600\\AA). \\cite{Blair1991} and \\cite{RM1992} have produced skymaps with nearly complete coverage of the Cygnus Loop, measured at \\ion{O}{6} $\\lambda\\lambda$1032, 1038. Observations of five 40$\\arcmin$ diameter fields around the Cygnus Loop rim with higher spatial resolution were obtained with the UIT \\citep{Cornett1992, Danforth2000}.  We report the first FUV spectral line images of the entire Cygnus Loop, These images, with spatial resolution of 3$\\arcmin$ to 30$\\arcmin$, were obtained with {\\em SPEAR} (The Spectroscopy of Plasma Emission from Astrophysical Radiation), also known as {\\em FIMS} (Far-ultraviolet IMaging Spectrograph). We describe how these data reveal physical conditions in the remnant and complement earlier observations. ",
        "conclusions": "The {\\em SPEAR} FUV spectral images of the Cygnus Loop offer a new view of the SNR physical environment and conditions. We find that the remnant's projected diffuse interior, ``carrot'', breakout region, and periphery show correlated FUV, optical, X-ray, and radio emission. The FUV emission can be well explained by the scenario wherein the Cygnus Loop has been created by a cavity explosion \\citep{Levenson1997}, with the presence of large clouds around its periphery \\citep{HC1986, HRB1994, Leahy2002}. Further study of these {\\em SPEAR} data will provide insight into the overall shock energetics through a more detailed comparison with earlier observations. These data will be particularly important for studying the NE nonradiative shock, where an additional set of {\\em SPEAR} short-wavelength data is expected to reveal large-scale features of \\ion{O}{6} emission."
    },
    "0601/astro-ph0601437_arXiv.txt": {
        "abstract": "We present BIMA and SCUBA observations of the young cluster associated with \\bdstar in the dense molecular gas tracer CS~J~=~2~$\\rightarrow$~1 and the continuum dust continuum at $\\lambda$~=~3.1~mm and  $\\lambda$~=~850~$\\mu$m. The dense gas and dust in the system is aligned into a long ridge morphology extending $\\sim$0.4 pc with 16 gas clumps of estimated masses ranging from 0.14 - 1.8 M$_\\odot$. A north-south variation in the CS center line velocity can be explained with a two cloud model. We posit that the \\bdstar stellar cluster formed from a cloud-cloud collision. The largest linewidths occur near V1318~Cyg-S, a massive star affecting its natal environment. In contrast, the dense gas near the other, more evolved, massive stars displays no evidence for disruption; the material must either be processed into the star, dissipate, or relax, fairly quickly. The more evolved low-mass protostars are more likely to be found near the massive stars. If the majority of low-mass stars are coeval, the seemingly evolved low-mass protostars are not older: the massive stars have eroded their structures. Finally, at the highest resolution, the $\\lambda$ = 3.1 mm dust emission is resolved into a flattened structure 3100 by 1500 AU with an estimated mass of 3.4 M$_\\odot$. The continuum and CS emission are offset by 1$\\farcs$1 from the southern binary source. A simple estimate of the extinction due to the continuum emission structure is A$_V$ $\\sim$ 700 mag. From the offset and as the southern source is detected in the optical, the continuum emission is from a previously unknown very young, intermediate-mass, embedded stellar object. ",
        "introduction": "Over the last few decades, a good broad-brush picture of isolated, low-mass, Sun-like star formation \\citep[e.g.,][]{shuppiii} has been developed. However, these theories are unlikely to explain properly the formation of our Sun or indeed the majority of other stars in the Galaxy. For our Sun in particular, the isotopic evidence strongly suggests that it formed in a location that was enriched either by a stellar wind or a recent supernova \\citep[e.g.,][]{lee1977}, hence in or near a cluster environment. New results show that the early solar system contained significant amounts of the short-lived $^{60}$Fe radionuclide \\citep[e.g.,][]{tachibana2003}, probably from a nearby supernova \\citep[$\\sim$ 1.5 pc distance,][]{supernova}.  The suggestion that our Sun might have been born in a cluster environment, should not be too surprising, as there is good evidence that the majority of stars in the Galaxy formed in clusters \\citep[e.g.,][]{carpenter2000,ladalada2003}.  Giant molecular clouds in the Galaxy provide the ideal environment for star formation, comprising dense molecular gas and dust that provide the raw material \\citep[e.g.][]{Zuckerman74}.  But how do we expect low-mass star formation to proceed in a cluster?  According to the standard model \\citep[e.g.][]{shuadaliz87}, it is a relatively long process requiring a core to exceed its Jeans mass, initiate self-similar collapse, followed by multiple fragmentation down to star formation with the Initial Mass Function (IMF) stellar distribution.  In this model, higher mass stars form quickly and low mass stars form slowly.  However, in clusters the environment is much more complicated. Clusters are often associated with one or many massive stars accompanied by a large number of lower mass stars.  In fact, the typical star forming system is actually a low stellar density cluster, including young intermediate to high mass star or stars (3-20 $M_{\\odot}$) surrounded by ten to a few hundred of low-mass protostars. This trend for massive stars to be more and more gregarious, is first associated with stars at the Herbig Ae/Be stage \\citep[e.g.][]{lynne95,palla95}.  In such an environment, the standard model has to be modified for the large influence that the massive star will have on low-mass star formation and evolution. The low-mass star may evolve faster than in isolation. It has been posited that stellar formation in a large cluster, with a combined mass of a few hundred $M_\\odot$, evolves faster by minimizing the effects of magnetic fields \\citep{shuadaliz87}.  Are low-mass stars formed coevally with massive stars during a trigger event?  Are low-mass stars triggered by the high mass star?  In order to better understand the differences between isolated low-mass star formation and low-mass star formation in a cluster with massive stars, it is important to understand the most simple case of small clusters, or groups, around intermediate-mass stars. One of the main points is to develop a better picture of how the most massive members of the cluster influence the formation of the lower-mass members by probing the dense molecular gas, dust, and stellar population. An excellent example is the small cluster surrounding the intermediate-mass/massive ($\\sim$ 9 $M_\\odot$), young \\citep[$\\sim$ 1 Myr;][]{lynne95} star \\bdstar, one of the original Ae/Be stars studied by \\cite{herbig}.  Although the main star has a spectral type of B2 Ve and is hot enough ($T_{\\rm eff} > 22\\,000$~K), there is no observational evidence for an \\hii region \\citep{skinner93}. Located in the Cygnus arm at a distance of about one kiloparsec, \\bdstar is the optically brightest member of a small group of young stars that include V1686~Cyg and the binary V1318~Cyg, all of which have infrared excesses \\citep{strom72} and are very young \\citep{cohen72}.  \\bdstar defines the center of a partially embedded aggregate of at least 33 stars, 80\\% of which are located within a projected distance of $0.15{\\rm\\,pc}$ from the central star \\citep{lynne95}. Single-dish maps of CO, $^{13}$CO, C$^{18}$O and CS, \\citep{lynne95,palla95} and NH$_3$ \\citep{fuente90}, show that this relatively isolated and dense ($n \\ga 1000\\,{\\rm stars\\,pc^{-3}}$) cluster is still associated with more than $\\sim 300\\,M_\\odot$ of clumpy, dense molecular material. Therefore, ongoing star-formation is still possible in this region.  In fact, \\cite{lynne95} concluded that high- and  low-mass stars are still forming, with the balance in favor of high-mass stars.  This small and young group of stars is interesting for two main reasons. First, the parent molecular core is isolated from any large star-forming complex, and, therefore, this region is an ideal site to investigate the conditions in which a single high-mass star, capable of becoming a supernova ($M_\\ast \\ga 8\\,$\\msun) forms, and its role, if any, in the formation of its lower-mass siblings. Second, \\cite{lynne95} concluded that star formation in this region was triggered by an external event, based on morphological arguments. However, no evidence can be found for any of the standard mechanisms for induced star formation (e.g. nearby expanding \\hii regions, supernova remnants, or nearby OB associations).  To unravel the star-formation history of this region it is first necessary to identify the sites of ongoing-- and potentially future-- star formation and to study the kinematic signatures and physical conditions of these sites. The observations focused on the CS~($2\\rightarrow 1$) transition, as a tracer of dense gas. Due to its high critical density, \\citep[$n_{H_2}\\sim10^{5}\\rm\\, cm^{-3}$, e.g.][]{irvine87}, CS~($2\\rightarrow 1$) is thermalized in the dense gas of star-forming cores.  In other words, CS emission is a good method for revealing dense, star forming clumps. To date, the available observations at millimeter wavelengths are of moderate spatial resolution-- they were all obtained with single-dish telescopes. In this paper, we present high spatial- and spectral-resolution data obtained with the BIMA array in the $J = 2\\rightarrow 1$ transition of CS and dust continuum emission that probe the dense gas and dust in the young cluster on spatial scales of 1000's to 10,000's of AU. ",
        "conclusions": "In this paper, we have probed the dense gas and dust in the small cluster associated with the intermediate-mass young star \\bdstar. The dense gas, as traced by CS J = 2 $\\rightarrow$ 1, consists of $\\sim$16 clumps that have the morphology of a dense molecular ridge about $\\sim$ 0.4 pc in size. By a simple LVG analysis of the clumps, we find that the typical estimated mass of the clumps is $\\sim$ 0.5 M$_\\odot$. Although it is impossible with this data to predict which, if any, of these clumps may contain a young, deeply embedded population of protostars, it seems likely that some clumps do.  In particular, we resolve a flattened dust structure, 3100 $\\times$ 1500 AU in size that is 1.1$\\arcsec$ offset from the southern component of the V1318 Cyg binary system, but well aligned with the water maser position. By modeling the circumstellar emission as arising from optically thin isothermal dust, a mass of 3.4 $M_\\odot$ is estimated for the structure. The V1318 Cyg binary system is optically visible; however, the extinction is expected to be around A$_V$ $\\sim$ 700 mag. With so high an extinction, the central source should not be visible.  If the optically visible V1318 S is the source of the continuum emission, it would require an unusual geometry that is not evident in the continuum emission. Based on the offset and the visibility of V1318 Cyg-S, we suggest that this object is a hereto unknown young embedded object and not V1318 Cyg-S. The structure, based on its mass, is probably a flattened circumstellar envelope which contains an intermediate mass, deeply embedded protostar. Indeed, star formation is still ongoing in this cluster.  However, we posit that, in general, star formation of this small cluster was coevally triggered by the collision of two clouds.  The dense gas detected in this paper is essentially the remnants of the two original cores.  The main reason for this conclusion is that the systematic velocity of the CS~($2\\rightarrow 1$) transition suggests the presence of two independently rotating clouds, distributed spatially: cloud 1 is the center of the southern star forming activity and cloud 2 is the northern spur.  Currently, cloud 1 is the center of the dense dust and gas in the cluster. Triggered star formation would probably explain the high star formation efficiency in the system.  Finally, there is evidence, e.g. the linewidths of the V1318 Cyg region, that the young intermediate mass stars in the core are disrupting the dense gas, which interferes with coeval lower-mass star formation. However, it must also be noted that this argument also implies that the previous stages of star formation, in particular the formation of \\bdstar and V1686 Cyg, did not disrupt the kinematics of the residual clouds significantly. In other words, the gas that was affected by their formation has been processed into the star, removed from the system, or relaxed fairly quickly. Another example of this is that the more evolved low-mass protostars are more likely to be near the more massive young stars. If the system was triggered by colliding clouds, then the majority of the low-mass stars are of similar ages. That implies that the presences of the higher mass protostars either hastens the lower-mass evolution, or that the more massive stars oblate the circumstellar structures of the lower-mass siblings, making them appear older."
    },
    "0601/astro-ph0601608.txt": {
        "abstract": "{The dominating reionization source in the young universe has not yet been identified. Possible candidates include metal poor dwarf galaxies with starburst properties.}{We selected an extreme starburst dwarf, the Blue Compact Galaxy Haro 11, with the aim of determining the Lyman continuum escape fraction from UV spectroscopy.} {Spectra of Haro 11 were obtained with the Far Ultraviolet Spectroscopic Explorer (FUSE). A weak signal shortwards of the Lyman break is identified as Lyman continuum (LyC) emission escaping from the ongoing starburst. From profile fitting to weak metal lines we derive column densities of the low ionization species. Adopting a metallicity typical of the H~II regions of Haro 11, these data correspond to a hydrogen column density of $\\sim$10$^{19}$cm$^{-2}$. This relatively high value indicates that most of the LyC photons escape through transparent holes in the interstellar medium. We then use spectral evolutionary models to constrain the escape fraction of the produced LyC photons.} {Assuming a normal Salpeter initial mass function we obtain a Lyman continuum escape fraction of f$_{esc}\\sim$ 4--10\\%. We argue that in a hierarchical galaxy formation scenario, the upper limit we derive for the escape rate allows for a substantial contribution to cosmic reionization by starburst dwarf galaxies at high redshifts.} {} ",
        "introduction": "According to WMAP data, the epoch of reionization started at a redshift of $z\\approx 20$ (Kogut et al. \\cite{kogut}). Observations of the high opacity in the Lyman continuum (hereafter LyC) emission of galaxies (Becker et al., \\cite{becker}, Fan et al. \\cite{fan2}) and the results from the recent studies of the intergalactic medium (IGM) around SDSS quasars (e.g. Mesinger \\& Haiman \\cite{mesinger}) indicate that the final stage of the epoch of reionization may have occurred at redshifts $z\\sim$6--7. There is still no consensus about what could have been the major ionization sources of the IGM during the reionization epoch -- stars, normal galaxies, AGNs or some exotic mechanism (Yan \\& Windhorst \\cite{yan}, Meiksin \\& White \\cite{meiksin}). The rapid decline of the space density of quasars at redshifts $z>$3 (Madau et al. \\cite{madau1}, Fan et al. \\cite{fan1}) makes it likely that their contribution to the LyC production rate required for the reionization is insignificant, although it has been proposed that miniquasars fed by intermediate-mass black holes may have been important (Madau et al. \\cite{madau2}).  At lower redshifts, observations of a sample of Lyman-break galaxies at $z \\sim$ 3.4 (Steidel et al. \\cite{steidel}) seem to show that galaxies could account for a significant fraction of the photon flux required for reionization. The observations indicate a contribution of LyC photons a factor of $\\sim$5 higher than for quasars at the same epoch. These results have however not been confirmed by other similar studies. On the contrary, it has been argued (Fern\\'andez-Soto et al. \\cite{fernandez-soto}) that a maximum of 4\\% of the LyC flux could escape from luminous galaxies at redshifts 1.9$<z<$3.5. This is also close to the upper limit of the escape fraction of low redshift galaxies (Leitherer et al. \\cite{leitherer1}; Hurwitz et al. \\cite{hurwitz}; Deharveng et al. \\cite{Deharveng et al.}; Heckman et al. \\cite{heckman}). The lack of a strong trend with redshift indicates that galaxies at higher redshifts also have low escape fractions, insufficient to keep the intergalactic gas ionized. Alternative ionization sources that have been discussed are massive pregalactic Pop III stars and young globular clusters. Most probably the reionization is ruled by different sources at different redshifts. It is therefore important to establish observational constraints on the LyC fluxes from the various candidates at different epochs.  Our information about the galaxy luminosity/mass function at the epoch of reionization is very meager, and it is not at all clear what mass scale would dominate the LyC photon production rate. In a galaxy formation scenario where mergers are of major importance one would expect starburst dwarf galaxies to be frequent among the first generation of galaxies. Indeed the number density of so called Luminous Compact Blue Galaxies (LCBGs) appears to increase with redshift (Lilly et al. \\cite{lilly}; Mall\\'en-Ornelas et al. \\cite{mallen-ornelas}). It is therefore important to investigate if this type of galaxy could have been an important ionization source at high redshift. Dwarf galaxies are however too faint to be observed at high redshifts. We are therefore dependent on theoretical models and observations of nearby galaxies that can be used as templates of the first generation of dwarf galaxies. Recent models do not give much hope for Lyman escape from galaxies of masses larger than 10$^7$ solar masses (Kitayama et al. \\cite{kitayama}) but these predictions have to be tested observationally.  We have searched for local dwarfs with properties similar to young dwarf galaxies, that could be used as test objects. A few such studies have already been carried out but only provided upper limits to the escape fraction (e.g., Giallongo et al. \\cite{giallongo}; Devriendt et al. \\cite{devriendt}; Shull et al. \\cite{shull}). However, predicted equivalent widths of the Balmer emission lines in starbursts are in some cases larger than observed (e.g. Bergvall \\cite{bergvall1}, Bresolin et al., \\cite{bresolin}, Mas-Hesse \\& Kunth \\cite{mashesse}, Moy et al. \\cite{moy}, Stasinska et al. \\cite{stasinska}). Several possible explanations to this were discussed by Mas-Hesse \\& Kunth (\\cite{mashesse}). One possibility is that the most massive stars remain hidden in their parental clouds until most of the fuel has been consumed. This would mimic a truncation of the IMF at high masses. Mas-Hesse and Kunth also discussed the possibility that a large fraction of the ionizing photons are absorbed by dust in the H II regions. To obtain agreement with the observations, a large amount of the photons, $\\sim$ 60\\% would have to be absorbed. This possibility has recently been investigated by Hirashita et al. (\\cite{hirashita}) where the escape fraction of the LyC photons from the H II region itself is estimated from model comparisons. Based on input data from a sample of star forming galaxies they find that, in the mean about 40\\% of the photons are trapped, which is less than what would be needed according to Mas-Hesse and Kunth. But the scatter is high so in some cases this certainly would be a possibility. Hirashita et al. seem to find that the escape fraction is higher is galaxies with properties similar to Haro 11 (e.g. the BCG ESO 338-IG04). Therefore an alternative, or complementary, explanation for the low flux of the recombination lines, may be that a significant fraction of the LyC photons is in fact escaping from the galaxy. This is the possibility that we study in this paper.  H~II regions in nearby galaxies appear to leak 40--50\\% of the LyC photons produced by the hot stars into the diffuse interstellar medium (Ferguson et al. \\cite{ferguson}; Oey et al. \\cite{oey}). Although it cannot be completely excluded that the gas may be ionized by single O stars located in situ, a comparison of the observed emission line strengths with theoretical models (Iglesias-P\\'aramo \\& Mu\\~noz-Tu\\~n\\'on \\cite{iglesias-paramo}; Wood \\& Mathis \\cite{wood}) give support to the leakage interpretation. An even more extreme case of a density bounded H~II region around a massive star cluster was reported by Leitherer et al. (\\cite{leitherer2}). They estimate that about 75\\% of the LyC radiation is leaking into the diffuse ambient medium. Depending on the distribution of the gas, the location where these photons are absorbed may be very different from case to case. In some situations the diffuse ionized gas may be much more extended than the H~II region producing the photons and will be ignored by the observer who consequently will underestimate the star formation rate.  The other alternative for the discrepancy between theory and observations of the Balmer emission line strengths is that the dust is unevenly distributed across the starburst region so that the younger regions are more reddened than the older ones (Calzetti et al. \\cite{calzetti}). This could reduce the equivalent width of H$\\alpha$ with up to a factor of $\\approx 2$.  The only way to really distinguish between a global leakage into to intergalactic medium and the other two alternatives is to search for leaking photons shortward of the Lyman limit.  In previous studies of starburst dwarfs (Bergvall \\cite{bergvall1}, Bergvall \\& \\\"Ostlin \\cite{bergvall4}, \\\"Ostlin et al. \\cite{ostlin1}, \\cite{ostlin2}) we have focused on a few luminous (--19$> $M$_B>$--21) metal poor blue compact galaxies (BCGs). They have properties similar to the LBGs at intermediate redshifts (e.g. Guzman et al. \\cite{guzman}). Here we report on such a study using FUSE, the Far Ultraviolet Spectroscopic Explorer.   \\subsection{The target galaxy}  The target galaxy of our choice is Haro 11 (=ESO 350-IG38) at $\\alpha_{2000}$=00$^h$36$^m$52.5$^s$, $\\delta_{2000}$=-33$^{\\circ}$49$\\arcmin$49$\\arcsec$. From the emission-line spectrum we derive a heliocentric radial velocity of v$_0$=6180 km s$^{-1}$, corresponding to a distance of 86 Mpc, assuming a Hubble constant of 72 km s$^{-1}$ Mpc$^{-1}$. We have previously studied this galaxy over a wide frequency range, from UV to radio (Bergvall \\& Olofsson \\cite{bergvall2}, Bergvall et al. \\cite{bergvall3}, \\cite{bergvall4}, \\\"Ostlin et al., \\cite{ostlin1}, \\cite{ostlin2}). It has unique properties that makes it one of the most relevant local objects for our purpose. The chemical abundances are low, the oxygen abundance being $\\sim$20\\% solar (relative to the revised solar oxygen abundance by Asplund et al. \\cite{asplund}).  An estimate of the typical hydrogen mass of a BCG can be obtained from the correlation between gas mass and major diameter (Gordon \\& Gottesman \\cite{gordon}). Haro 11 has a major diameter of 11 kpc at $\\mu_B$=25 mag arcsec$^{-2}$ indicating that the H~I mass should be approximately 2$\\times$10$^9$ \\sma. Despite several efforts at VLA, Nancay and Parkes (all unpublished) we have not been able to detect H~I in this galaxy. We estimate the upper limit to \\ma (H~I)$\\sim$ 10$^8$ \\sma. The extraordinary strong [C~II]$\\lambda$158$\\mu$ line, observed with ISO (Bergvall et al. \\cite{bergvall3}), shows that a large part of the H~I gas is located in photodissociation regions (PDRs). We have estimated the contribution from H$_2$ and H~II gas to between 10$^8$ and 10$^9$ \\sma~ each. This would leave a remarkably small fraction of the total gas mass in atomic form. A possible explanation is that the bulk of the gas in the halo that normally is in neutral state has become ionized by the starburst. Our H$\\alpha$ observations show that the main body is indeed surrounded by an extended halo of ionized gas. The conditions for LyC photon escape thus appear to be unusually favourable. It is interesting to note that the nearby BCG Pox 186, sometimes called 'ultracompact', is also H I quiet (Begum \\& Chengalur \\cite{begum}).  HST GHRS UV spectroscopic observations of Haro 11 were carried out by Kunth et al. (\\cite{kunth}) who found a quite complex absorption system with velocity components of Ly$\\alpha$ in absorption on both sides of the emission line. Thus it appears that there are both outflows and infall of gas in the centre direction. From the absorption lines Kunth et al. estimate column densities and find values as high as log n$_H \\sim$ 20 cm$^{-2}$. The aperture used is small however (2$\\arcsec$ circular diameter) and covers only the very central part of the galaxy. Our optical spectroscopy (\\\"Ostlin et al. \\cite{ostlin05}) reveals velocity gradients of both gas and stars as large as 200 km s$^{-1}$ across two of the bright condensations in the very centre, indicating an extremely high mass density in this region. ",
        "conclusions": "We report on the discovery of a weak signal in the Lyman continuum of the luminous starburst galaxy Haro 11. It is the first time Lyman continuum leakage is found in a galaxy in the local environment. Using the derived column densities of low ionization species we estimate the column density of neutral hydrogen to be $\\sim$ 10$^{19}$ cm$^{-2}$, which is too high to allow Lyman continuum photon escape if the gas is diffusely distributed. We conclude that the Lyman continuum radiation is escaping through transparent windows of the ISM. Assuming a Salpeter IMF, we estimate a total escape fraction of 0.04$<f_\\mathrm{esc}<$0.10. More extreme IMFs allow a lower escape fraction but the major uncertainty is due to the problematic background subtraction at short wavelengths. The upper limit is below the model estimates of the minimum escape rate required for reionization. Still, this single case demonstrates that a significant fraction of the required Lyman continuum photons may be produced by starburst dwarfs, in particular if we also consider that Pop III stars, that should dominate the first generation galaxies, are more efficient than Pop II stars in producing Lyman continuum photons."
    },
    "0601/astro-ph0601571_arXiv.txt": {
        "abstract": "AMANDA-II is the largest neutrino telescope collecting data at the moment, and its main goal is to search for sources of high energy extra-terrestrial neutrinos. The detection of such sources could give non-controversial evidence for the acceleration of charged hadrons in cosmic objects like Supernova Remnants, Micro-quasars, Active Galactic Nuclei or Gamma Ray Bursts. No significant excess has been found in searching for neutrinos from both point-like and non-localized sources. However AMANDA-II has significantly improved analysis techniques for better signal-to-noise optimization. The km$^3$-scale IceCube telescope will enlarge the observable energy range and improve the sensitivities of high energy neutrino searches due to its 30 times larger effective area. ",
        "introduction": "\\label{s:intro}  The detection of extra-terrestrial multi-TeV neutrinos would significantly contribute to the understanding of the origin, acceleration and propagation of high energy cosmic rays. It would also influence the interpretation of results of the most recent gamma ray experiment results, like those of H.E.S.S. \\cite{hess}.  Neutrinos can propagate through the Universe without being affected by interactions or magnetic fields. Therefore they are ideal high energy cosmic messengers and, for the same reason, unfortunately very difficult to detect. Even the so called \"guaranteed\" neutrino fluxes, i.e. those associated with known cosmic accelerators with an identified pion production target, typically require a $\\sim$ km$^3$ scale detector. AMANDA-II \\cite{amanda,amanda2}, with its $\\sim 0.016$ km$^3$ instrumented volume, is currently the largest operating neutrino telescope. Its sensitivity is improving because of the larger data sample collected, better apparatus understanding, and more efficient background rejection techniques.  The main background for neutrino telescopes is the intense flux of downward penetrating cosmic muons which can mimic the upward neutrino-induced event signature. A total rejection factor of $10^6$ is required to isolate muons arising from neutrinos created in atmospheric cosmic ray showers. AMANDA-II has improved reconstruction algorithms~\\cite{reco} and analysis techniques and currently it is able to select about four upward going muons per day \\cite{icrc05}. After rejecting the downward muon background, these atmospheric neutrinos remain as an irreducible background in searches for extraterrestrial neutrino fluxes.  The following sections describe the searches for localized neutrino sources and for a diffuse neutrino flux from many weak sources, respectively. ",
        "conclusions": ""
    },
    "0601/astro-ph0601092_arXiv.txt": {
        "abstract": "We present a method that employs the secondary eclipse light curves of transiting extrasolar planets to probe the spatial variation of their thermal emission. This technique permits an observer to resolve the surface of the planet without the need to spatially resolve its central star. We evaluate the feasibility of this technique for the \\hdtwo{} system by first calculating the secondary eclipse light curves that would result from several representations of the planetary emission, and then simulating the noise properties of observations of this signal with the \\spitzer{} Infrared Array Camera (IRAC).  We consider two representations of the planetary thermal emission; a simple model parameterized by a sinusoidal dependence on longitude and latitude, as well as the results of a three-dimensional dynamical simulation of the planetary atmosphere previously published by Cooper \\& Showman. We find that observations of the secondary eclipse light curve are most sensitive to a longitudinal offset in the geometric and photometric centroids of the hemisphere of the planet visible near opposition.  To quantify this signal, we define a new parameter, the ``uniform time offset,'' which measures the time lag between the observed secondary eclipse and that predicted by a planet with a uniform surface flux distribution.  We compare the predicted amplitude of this parameter for \\hdtwo{} with the precision with which it could be measured with IRAC.  We find that IRAC observations at \\iraca{} of a single secondary eclipse should permit sufficient precision to confirm or reject the Cooper \\& Showman model of the surface flux distribution for this planet.  We quantify the signal-to-noise ratio for this offset in the remaining IRAC bands (\\iracb{}, \\iracc{}, and \\iracd{}), and find that a modest improvement in photometric precision (as might be realized through observations of several eclipse events) should permit a similarly robust detection. ",
        "introduction": "\\label{sec:intro}  The identification of the first transiting extrasolar planet \\hdtwob{} \\citep{charb00,henry00,mazeh00} initiated a flurry of investigations into the properties of the planetary body that are not possible for non-transiting objects.  Nine transiting extrasolar planets have now been identified; for a review of their properties, see Charbonneau et al.\\ 2006. In addition to HD~209458b, three of these (TrES-1; Alonso et al.\\ 2004, HD~149026b; Sato et al.\\ 2005, and HD~189733b; Bouchy et al.\\ 2005) orbit stars that are sufficiently close and hence bright enough ($V < 12$) to permit a direct study of their atmospheric absorption and emission features through a variety of techniques.  One such method is that of transmission spectroscopy, whereby stellar spectra gathered outside and inside times of planetary transit are compared to search for additional absorption features in the latter due to the presence of certain atoms or molecules in the planetary atmosphere.  The only detections of this effect to date have been achieved with the STIS spectrograph aboard the \\emph{Hubble Space Telescope}:  \\citet{charb02} observed \\hdtwo{} in visible light and detected the absorption from gaseous atomic sodium in the planetary atmosphere, and \\citet{vidal03} observed the same system in the ultraviolet to detect absorption resulting from a large cloud of atomic hydrogen escaping from the planet.  Numerous ground-based observational efforts (Bundy \\& Marcy 2000; Moutou et al.\\ 2001, 2003; Brown et al. 2002; Winn et al. 2004; Narita et al. 2005) have yielded only upper limits, albeit useful ones.  Most recently, \\citet{deming05a} placed a stringent upper limit on the presence of gaseous CO from observations near 2.3~\\micron. Along with the sodium detection, these observations place tight constraints on the distribution of condensates in the upper atmosphere (Fortney 2005).  A complementary technique that promises to be at least as powerful is that of occultation photometry and spectroscopy.  This method subtracts observations gathered during secondary eclipse (i.e. when the planet passes \\emph{behind} the star) with those gathered just before or after this time (when the planet is unocculted), to search for any excess emission attributable to the planet itself.  For a hot Jupiter planet orbiting a Sun-like star, the relative size of this excess at infrared wavelengths is a few parts in one thousand. Ground-based attempts to observe thermal radiation from transiting planets have been frustrated by the variability of the telluric opacity over short timescales, and the large, ambient thermal background, and thus have resulted in only upper limits \\citep{wiedemann01,lucas02,richardson03a,richardson03b}.  Recently, two groups have succeeded in detecting the planetary thermal emission, and once again this feat was enabled by a space-based observatory, in this case the \\spitzer{} (Werner et al.\\ 2004).  \\citet{charb05} used the Infrared Array Camera (IRAC; Fazio et al. 2004) to detect the thermal emission of TrES-1 in two band passes, \\iracb{} and \\iracd{}. \\citet{deming05b} used the Multiband Imaging Photometer (MIPS; Reike et al.\\ 2004) to detect the thermal emission from \\hdtwob{} at 24~\\micron, and \\citet{deming06} employed the Infrared Spectrograph (IRS; Houck et al. 2004) to detect the emission from HD~189733b at 16~\\micron.  Each of the data sets consisted of a rapid cadence, high signal-to-noise ratio (SNR) photometric light curve spanning a predicted time of secondary eclipse, in which a decrement in the total system flux of the expected depth and duration was clearly detected at the anticipated time.  An intriguing possibility permitted by such observations is that of resolving the surface of the planet through high SNR photometry during the times of ingress and egress, when different portions of the planetary surface are occulted by the star. The purpose of this paper is to explore this effect in detail, and to evaluate the likelihood of its detection with the \\spitzer{}.  It is not currently possible to directly image exoplanetary surfaces, nor is this ability anticipated for any planned facility, including the NASA \\emph{Terrestrial Planet Finder}. Nonetheless, the spatial-dependence of the planetary photosphere is accessible to an observer through careful monitoring of the structure in a secondary-eclipse light curve.  During the ingress and egress phases of secondary eclipse, which last approximately $15 - 30$ minutes depending on the orbital geometry and planetary radius, the portion of the visible planetary hemisphere (the ``dayside'') that remains unocculted varies smoothly as a function of time. From the known system parameters (namely the orbital period, phase, inclination, and radii of the planet and star), it is then possible to invert the observed light curve to recover some aspects of the flux distribution across the dayside of the planet.  This technique is not new; a similar approach has been used, for example, to produce surface maps of Pluto and Charon (for a review for such observations, see Stern 1992). But only with the \\emph{Spitzer} detections of the past year has it been feasible to consider applying this technique to extrasolar planets.  This possibility is particularly interesting because several recent dynamical simulations have predicted the presence of a large flux contrast across the dayside of a hot Jupiter planet. Independent studies of \\hdtwob{} by \\citet{showman02} and \\citet{cooper05} predict a strong eastward jet.  The CS05 supersonic jet pushes the atmosphere's hottest region downstream by 60 degrees from the planetary substellar point.  The planetary circulation pattern leads to global temperature variations of $\\sim$500 K at photospheric pressures.  Simulations by \\citet{cho03} find three broad east-west jets and polar vortex motions.  They find that temperature variations may reach up to $\\sim$1000 K in certain circumstances.  \\citet{burkert05} find similarly large temperature variations. They highlight how atmospheric opacity, which controls the penetration depth of stellar flux, affects the circulation problem. Rapid cadence light curves during ingress and egress may allow for observational tests of these circulation models. The analysis of such light curves would provide insight into the dynamical flows of hot Jupiters and also illuminate the energy budgets of these planets.  In \\S\\ref{sec:motivation}, we further motivate the use of a secondary eclipse light curve to study a planetary surface flux distribution. In \\S\\ref{sec:lightcurves}, we describe the software that we developed for simulating light curves, which forms the basis of our subsequent investigations. In \\S\\ref{sec:fits}, we introduce a parameter, the ``uniform time offset,'' which characterizes the longitudinal flux contrast of a planetary emission. We then probe the behavior of this parameter with Monte Carlo simulations of hypothetical light curves. Finally, in \\S\\ref{sec:discussion}, we discuss our results and near-future applications of this technique. In this \\selfname, we restrict our focus to upcoming observations of \\hdtwob{} with the \\emph{Spitzer} IRAC instrument. It should be emphasized, however, that our technique can be applied to other instruments and transiting extrasolar planets. ",
        "conclusions": "\\label{sec:discussion}  In this \\selfname, we have explored how the secondary eclipse light curves of an extrasolar planet may be used to learn about the spatial variation of its emission. This technique permits an observer to gain access to this important information without the need to directly image to surface of the planet. We have explored this technique with simulations of \\emph {Spitzer} IRAC observations of the planet \\hdtwob{} and considered the resulting value of the uniform time offset parameter $t_{offs}$ that would be observed if the emission of the planet is consistent with a family of simple models whereby the flux varies sinusoidally in longitude and latitude, as well as the results of a three-dimensional dynamical calculation of the atmosphere of this planet as previously published by Cooper \\& Showman (2005).  We found that physically reasonable longitudinal flux contrasts could plausibly be detected in the \\hdtwo{} system with planned \\emph {Spitzer} IRAC observations. Specifically, sinusoidal models with a latitudinal parameter $|B| > 0.5$ yield a value of $P_0 < 0.25$ in all but the 5.6~\\micron{} bands. From the definition of the sinusoidal model, we see that $B = 0.5$ corresponds to a flux contrast of a factor of 3 between the leading and trailing edges of the planet. In the IRAC \\iracb{} band this contrast would correspond, for example, to two regions emitting as blackbodies with temperatures of approximately 1300K and 900K. We find that the results of the CS05 model are readily testable with \\emph{Spitzer} IRAC observations. In particular, observations in the \\iraca{} band would already be able to exclude this model with a high degree of confidence, should a value of $t_{offs} = 0$ be observed.  Of course, interpretation of such data will undoubtedly be more complex that the approach presented here.  For example, the presence of high-altitude clouds could decrease the atmospheric pressure corresponding to the photosphere, which would serve to mask the underlying dynamics of the atmosphere \\citep{showman05}, where the effect of winds on redistributing the energy of the incident stellar flux is much more prominent.  It should be noted that the value of the orbital eccentricity $e$ and longitude of periastron $\\omega$ also serve to affect the time of secondary eclipse, which could mimic the effect discussed here. Tidal circularization is expected to reduce the eccentricity of hot Jupiter orbits to virtually zero, in which case observations of the primary eclipse (e.g.  \\citet{brown01}) of such systems can generally constrain the predicted time of secondary eclipse to several seconds, which contributes a negligible source of error to the value of $t_{offs}$. Direct observational constraints on the orbital eccentricity at the required level of precision are likely not feasible, however. The current upper limit on the orbital eccentricity of HD~209458 from the radial velocity observations alone is $e<0.02$ (\\citet{laughlin05}), which could induce an offset in the time of secondary eclipse as large as 65 minutes, more than two orders of magnitude greater than the effect we describe. Rather, the validity of the assumption of a circularized orbit can be bolstered by checking for the presence of transit timing variations \\citep{agol05,holman05} using a sequence of transit observations. The observed lack of such variations (Knutson et al. 2006) indicate the absence of perturbers in the \\hdtwo{} system, thus permitting tidal effects to complete the circularization of the orbit.  We also note that the timing offset due to a non-zero orbital eccentricity is fundamentally achromatic, in contrast to the signal we describe here (e.g. Table~1).  By observing a single secondary eclipse simultaneously in multiple bandpasses (such as the 4 bands of the IRAC instrument), an observer can search for variations in the observed value of $t_{offs}$ between the various wavelength bands, a sure sign that the signal is not due to residual orbital eccentricity.  In particular, comparison between the values of $t_{offs}$ between the 3.6~$\\mu$m and 8.0~$\\mu$m bands would show the largest effect.  There are several ways in which the predictions presented here can be refined. More precise values for the orbital parameters and an observational determination of the eclipse depths could refine our theoretical predictions. But the greatest source of uncertainty in our results is the model of the planetary atmosphere and its emission. As mentioned before, clouds might dramatically affect the simulation results. As alluded to in \\S\\ref{ssec:temptoflux}, limb darkening is another effect that we do not take into account. Planets other than \\hdtwob{} may even exhibit limb brightening: this possibility has been investigated with regards to the planet HD 149026b by \\cite{fortney06a}. Depending on their assumed model parameters, they produce physically-plausible atmospheres that exhibit either limb brightening or limb darkening. In either case, however, the phenomenon is radially-symmetric and does not significantly alter the longitudinal flux contrast of the planet, and so the effect on $t_{offs}$ should be modest. We have confirmed this intuition by performing Monte Carlo simulations with the CS05 model modified to have its emission scaled according to the classical limb-darkening law $I'(r) = I(r) * (1 - c[1 - \\sqrt{1 - r^2}])$, where $r \\in [0, 1]$ is the normalized projected radial distance of a point from the center of the planetary disk and $c$ sets the magnitude of the effect. With $c = 0.25$, equivalent to a temperature decrement of $\\sim$300~K on the limb of the planet (which is much larger than what is calculated for \\hdtwob{}), the resulting change in $t_{offs}$ for all IRAC bands is $\\approx3$ s.  In the near future, there will be many opportunities to employ this technique for probing exoplanets. This \\selfname{} has considered the specific case of observations of \\hdtwob{} with the \\spitzer's IRAC. Such observations should yield unprecedentedly precise secondary eclipse light curves, and we conclude that they can realistically  be expected to probe the surface flux distribution of the day side of the planet. They would also determine the secondary eclipse depths of \\hdtwob{} in the 4 IRAC bands, which would be of immediate interest in their own right. Even more promising would be observations of the recently-discovered extrasolar planet HD 189733b \\citep{bouchy05}, which is a mere 19 pc away and has a very favorable planet-to-star flux ratio, owing to the relatively large planet-to-star surface area ratio, and the large equilibrium temperature of the planet. \\emph{Spitzer} has already observed HD~189733 with all 3 instruments, and we encourage a search for the effects described in this paper. In the longer term, the {\\em James Webb Space Telescope} (JWST), with infrared detectors and a larger aperture than {\\em Spitzer}, will be an extremely powerful tool for applying this technique. As instrumentation and models improve, and as an increasing number of nearby transiting extrasolar planet systems are discovered, the prospects for resolving the surfaces of these distant worlds grows ever brighter."
    },
    "0601/astro-ph0601217_arXiv.txt": {
        "abstract": "% Recent high-resolution observations of the central region of Galactic globular clusters have shown the presence of a large variety of exotic stellar objects whose formation and evolution may be strongly affected by dynamical interactions. In this paper I  review the main properties of two classes of exotic objects: the so-called Blue Stragglers stars and the recently identified optical companions to Millisecond pulsar. Both these  class of objects are invaluable tools to investigate the binary evolution in very dense environments and   are powerful tracers of the dynamical history of the parent cluster. ",
        "introduction": "Ultra-dense cores of Galactic Globular Clusters (GCs)  are very efficient ``furnaces'' for generating exotic objects, such as low-mass X-ray binaries, cataclysmic variables, millisecond pulsars (MSPs), blue stragglers (BSS), etc. Most of these objects are thought to result from the evolution of various kinds of binary systems originated and/or hardened by stellar interactions. The nature and even  the existence of binary by-products can be strongly affected by the cluster core dynamics, thus   serving as a diagnostic  of the dynamical evolution of GCs.  This topic has received strong impulse in the recent years and many studies have been devoted to investigate the possible link between the dynamical  evolution of clusters and the evolution of their stellar population.  In particular two aspects can be investigated: (1) the environment effects on canonical evolutionary sequences (as for example the possible  effect of different environments on the  blue tail extension of the Horizontal Branch - HB, see Fusi Pecci et al 1992, Buonanno et al 1997); (2) the creation of artificial sequences as BSS and other exotic objects (see Bailyn 1995 and reference therein).  In this paper I will review the main properties of two anomalous sequences in the color-magnitude diagram (CMD):  \\begin{itemize} \\item the most  known {\\it anomalous sequence} in GCs: the so-called BSS sequence, indeed the very first sequence of exotic objects discovered in the CMD of GCs; \\item the most recently discovered {\\it anomalous sequence}: the one defined by MSP companions. \\end{itemize} ",
        "conclusions": ""
    },
    "0601/astro-ph0601398_arXiv.txt": {
        "abstract": "We have obtained near-IR spectra of five Active Galactic Nuclei that exhibit double-peaked Balmer Emission Lines (NGC~1097, Pictor~A, PKS~0921--213, 1E~0450.30--1817, and IRAS~0236.6--3101). The stellar velocity dispersions of the host galaxies were measured from the \\caii~$\\lambda\\lambda 8494,8542,8662~\\ang$ absorption lines and were found to range from 140 to 200~$\\kms$. Using the well-known correlation between the black hole mass and the stellar velocity dispersion, the black hole masses in these galaxies were estimated to range from $4\\times 10^{7}$ to $1.2\\times 10^{8}~\\msol$. We supplement the observations presented here with estimates of the black holes masses for five additional double-peaked emitters (Arp~102B, 3C~390.3, NGC~4579, NGC~4203, and M81) obtained by other authors using similar methods. Using these black hole masses, we infer the ratio of the bolometric luminosity to the Eddington luminosity, ($L_{bol}/\\ledd$). We find that two objects (Pictor~A and PKS~0921--213) have $L_{bol}/\\ledd \\sim\\;$0.2, whereas the other objects have $L_{bol}/\\ledd \\lesssim 10^{-2}$ (nearby, low-luminosity double-peaked emitters are the most extreme, with $L_{bol}/\\ledd \\lesssim 10^{-4}$).  The physical time scales in the outer regions of the accretion disks (at $r\\sim 10^3\\; GM/c^2$) in these objects were also estimated and range from a few months, for the dynamical time scale, to several decades for the sound crossing time scale.  The profile variability in these objects are typically an order of magnitude longer than the dynamical time, but we note that variability occurring on the dynamical time scale has not been ruled out by the observations. ",
        "introduction": "During the past two decades, the number of Active Galactic Nuclei (AGN) known to exhibit broad, double-peaked Balmer emission lines has been steadily increasing. These lines were first observed in the broad-line radio galaxies (BLRGs) Arp~102B \\citep*{ssk83,ch89,chf89}, 3C~390.3 \\citep{o87, p88}, and 3C~332 \\citep{h90} and are reminiscent of the double-peaked emission lines observed in Cataclysmic Variables, where they are regarded as the kinematic signature of the accretion disk \\citep[e.g.,][]{hm86}. Thus it was suggested by the above authors that the double-peaked emission lines originate in the accretion disk that fuels the AGN.  A systematic survey of $z<0.4$ BLRGs and radio-loud quasars was undertaken later on and it was found that $\\sim$~20\\% of the observed objects exhibited double-peaked Balmer lines \\citep{eh94,eh03}.  More recently, \\citet{s03} found that 3\\% of the $z<0.332$ AGNs in the Sloan Digital Sky Survey \\citep[SDSS; ][]{y00} are double-peaked emitters. \\footnotemark\\footnotetext{\\citet{s03} finds that double-peaked emitters are preferentially radio-loud, thus the larger sample fraction among BLRGs is expected. See \\citet{s03} for a description of the differences between these two samples.} Double-peaked emission lines have also been observed in Low Ionization Nuclear Emission Line Regions \\citep[LINERs; ][]{h80} including NGC~1097 \\citep{sb93}, M81 \\citep{b96}, NGC~4203 \\citep{s00}, NGC~4450 \\citep{h00}, and NGC~4579 \\citep{b01}. Given the difficulty of detecting broad emission lines in LINERs \\citep{ho97} it is likely that a larger number of LINERs exhibit double-peaked emission lines. Although the known double-peaked emitters represent a small fraction of the total AGN population, they are an important and intriguing class of objects that deserve further study.  As described in detail by \\citet{eh03}, observational tests and basic physical considerations strongly suggest that the double-peaked Balmer lines originate in the outer regions of the accretion disk at distances from the black hole of hundreds to thousands of gravitational radii ($\\rg$~=~$G\\mbh/c^{2}$, where $\\mbh$ is the mass of the black hole). The {\\it profiles} of the double-peaked emission lines are observed to vary on time scales of years to decades and show occasional reversals in the relative strength of the red and blue peaks \\citep*[for examples, see][]{vz91,zvg91,g99,shap01,sps00,serg02,sb03,suvi04,l04a,l04b,lewis05,gezari05}. Intensive monitoring of 3C~390.3 shows that the time lag between variations in the continuum and the optical broad line is $\\sim$ 20 days \\citep{d98,serg02} thus the profile variations occur on much longer time scales than the light-crossing time and seem to be unrelated to fluctuations in the ionizing continuum flux that take place on time scales of a few days to a few weeks.  A similar effect is also observed in systematically monitored Seyfert galaxies \\citep[e.g.,][]{wp96,k97}. The profile variability is likely to be a manifestation of {\\it physical} changes in the outer disk and the profile variability is thus a powerful tool that can be used to test models for dynamical behavior in accretion disks. Therefore, in the early 1990s a campaign was undertaken to systematically monitor a set of $\\sim 20$ double-peaked emitters, including BLRGs from the \\citet{eh94} sample and several LINERs.  \\citet{eh03} and \\citet{s03} found that the profiles of approximately 60\\% and 40\\%, respectively, of the double-peaked emitters from their samples can be fitted by a simple, axisymmetric accretion disk model. However the long-term profile variability and the fact that in some objects the red peak is stronger than the blue peak (in contrast to the expectations of relativistic Doppler boosting) require the use of more general, non-axisymmetric models.  Some of the models that have been successfully employed to fit observed profiles include circular disks with orbiting bright spots \\citep[Arp~102B and 3C~390.3; ][]{ne97,sps00,zvg91} or spiral emissivity perturbations \\citep[3C~390.3, 3C~332, and NGC~1097; ][]{g99,sb03} and precessing elliptical disks \\citep[NGC~1097, ][]{e95,sb95,sb97,sb03}.  To successfully interpret and model the long-term profile variability, it is not sufficient to reproduce the observed sequence of profiles; the variability must occur on a physical time scale that is consistent with the chosen model. The time scales of interest are the light crossing, dynamical, thermal, and sound crossing times, which are set by the black hole mass ($\\mbh$) and are given by \\citet*{fkrbook}: \\begin{eqnarray} \\tau_{\\ell} & = & 6\\; M_8\\;\\xi_3 \\, {\\rm days} \\label{tl}\\\\ \\tau_{dyn} & = &  30\\; \\tau_{\\ell}\\;\\xi_3^{1/2}\\; = \\;6\\; M_8\\;\\xi_3^{3/2} \\, {\\rm months} \\label{tdyn}\\\\ \\tau_{th}  & = & \\tau_{dyn}/\\alpha\\; = \\;5\\; M_8\\;\\xi_3^{3/2} \\, {\\rm years}\\\\ \\tau_{s}   & = & 70\\; M_8 \\;\\xi_3\\; T_5^{-1/2} \\, {\\rm years} \\label{ts} \\end{eqnarray} \\noindent where $M_{8}=\\mbh/10^8\\;\\msol$, $\\xi_{3} = r/10^{3}\\;\\rg$, $ T_{5}= T/10^{5}\\;$K, and $\\alpha~(\\sim 0.1)$ is the Shakura-Sunyaev viscosity parameter \\citep{ss73}. The above models all predict variability on different time scales. For example matter embedded in the disk orbits over $\\tau_{dyn}$, thermal instabilities will dissipate over $\\tau_{th}$, density perturbations, such as a spiral wave, precess on time scales that are an order of magnitude longer than $\\tau_{dyn}$ up to $\\tau_{s}$, and an elliptical disk will precess over even longer time scales.  Many of the models described above yield strikingly similar sequences of profiles. In order to determine which physical mechanism is responsible for the profile variability, it is necessary to connect the observed variability time scale (in years) with one of the above {\\it physical} time scales. Using an estimate of the black hole mass in NGC~1097, \\citet{sb03} estimated the physical time scales in the outer accretion disk which favored a spiral arm model over an elliptical disk model for this object. However, in order to discriminate between the above time scales, the black hole mass must be known to better than an order of magnitude and preferably to better than a factor of three.  Another motivation for accurately estimating the black hole masses in the double-peaked emitters is to determine the accretion rate, relative to the Eddington accretion rate ($\\mdot/\\mdotedd$), which is a diagnostic of the structure of the inner accretion flow. Because the local energy dissipation in the accretion disk is not always sufficient to power the observed H$\\alpha$ luminosity, \\citet{ch89} proposed that the inner accretion flow in double-peaked emitters has the form of an ion torus \\citep{r82}, or similar vertically extended, radiatively inefficient flow. This vertically extended structure would be able to illuminate the outer accretion disk and power the observed emission lines.  For the well-studied \\citet{eh94} sample there is considerable indirect observational support of this hypothesis, discussed fully in \\citet{eh03}, which is underscored by the growing number of LINERs known to sport double-peaked emission lines. The assumption that the inner accretion flow is radiatively inefficient is fundamental to the current ideas for the formation of the double-peaked emission lines and also the relationship between double-peaked emitters and the AGN population in general.  Thus it is extremely important to test this hypothesis more directly by obtaining estimates of the black hole masses, and thus $\\mdot/\\mdotedd$, for double-peaked emitters.  Most of the double-peaked emitters are too distant to obtain black hole masses directly via spatially resolved stellar and gas kinematics. However, it is possible to determine the black hole masses indirectly through the well-known correlation between $\\mbh$ and the {\\it stellar} velocity dispersion ($\\sigma$) measured on the scale of the effective radius of the bulge \\citep{fm00,g00,trem02}. Therefore we have begun a program to measure the stellar velocity dispersions in double-peaked emitters using the \\caii~$\\lambda 8594,8542,8662$ triplet and present here the results for five objects. After obtaining estimates of the black hole masses via the \\msig~relationship, the physical time scales in the outer accretion disk and $\\mdot/\\mdotedd$ can be inferred for each object.  Here, we focus specifically on measurements of the velocity dispersion in these double-peaked emitters to enable determinations of black hole masses. Although the time scales derived here are extremely valuable in modeling the long-term profile variability in these objects, a detailed discussion of the profile modeling is deferred to forthcoming papers. Preliminary results on modeling the profile variability (with the help of the black hole masses obtained here) can be found in \\citet{lewis05} and \\citet{gezari05}.  This paper is organized as follows. In \\S\\ref{mass_data}, we describe the target selection, observations and data reductions. The analysis of the data, including an examination of the various sources of error, are described in \\S\\ref{mass_analysis} and the results and their implications are presented and discussed in \\S\\ref{mass_discussion}. In Appendix \\ref{mass_ap1}, we present a detailed description of the procedure used to correct the telluric water vapor absorptions lines in these spectra. Throughout the paper, we assume a WMAP cosmology \\citep[$H_{\\rm 0}=70~{\\rm km~s^{-1}~Mpc^{-1}}$, $\\Omega_{\\rm M}=0.27$, $\\Omega_{\\Lambda}=0.73$;][]{wmap}. ",
        "conclusions": "\\label{mass_discussion}  Using the measurements of the velocity dispersion found above, we determine the black hole masses, the Eddington ratios, and the physical timescales in the accretion disk. All of these quantities are given in Table \\ref{results_tab}. The black hole masses were estimated via the \\msig~relationship \\citep{trem02}, namely, \\begin{equation} \\log\\left({M\\over\\msol}\\right) = \\alpha + \\beta\\; \\log\\left({\\sigma\\over\\sigma_0}\\right)\\; , \\end{equation} where $\\alpha = 8.13 \\pm 0.06$, $\\beta = 4.02 \\pm 0.32$, and $\\sigma_0 = 200~\\kms$.  The scatter in the \\msig~relationship is not included in the error, but the error bars on the coefficients are. For completeness, we include in this table the double-peaked emitters with stellar velocity dispersions reported in the literature, namely Arp~102B, NGC~4203, and NGC~4579 \\citep{b02} and 3C~390.3 \\citep{nel04}. We also include M81, for which a mass estimate of $\\sim~6\\times10^{7}~\\msol$ is based on resolved stellar and kinematics \\citep[][respectively]{b00,dev03}. The inferred black hole masses for these ten objects range from $4\\times10^{7}~\\msol$ to $5\\times 10^{8}~\\msol$.  We note the estimate of the black hole mass for NGC~4203 used here ($\\mbh~\\sim~6~\\times10^{7}~\\msol$) is in disagreement with the results of \\citet{s00} and \\citet{sarzi01}, who placed an upper limit of $6\\,\\times\\,10^{6}~\\msol$ on the black hole mass from spatially resolved gas kinematics. As noted by \\citet{s00}, the gas need not lie in a flat disk and it could be subject to non-gravitational forces; in the particular case of NGC~4203, the disk is is suspected to be warped \\citep{sarzi01}. Thus, we have chosen to use the mass obtained from the stellar velocity dispersion.  The physical timescales in the outer accretion disk were computed using equations~\\ref{tdyn}--\\ref{ts}, assuming that $\\xi_{3} = 1$, $\\alpha = 0.1$ and $T_{5} = 1$. For these black hole masses, the dynamical timescales ranges from a few months to one year, the thermal timescales ranges from one year to a decade, and the sound crossing timescales ranges from a decade to 100 years. The observed profile variations occur on timescales of several years \\citep[see, e.g., ][]{lewis05,gezari05} for most objects, which is an order of magnitude longer than the dynamical timescales found here.  This strongly suggests that in general the profile variability might be due predominantly to either a thermal phenomenon or the precession of a large scale emissivity pattern (e.g. a spiral arm).  One exception is Arp~102B, in which the relative fluxes of the red and blue peaks of the profile varied sinusoidally with a period of 2 years (of the same order as the dynamical time) over an interval of four years \\citep{ne97}. A very similar variability pattern, with the same period, also appeared five years later \\citep{gezari05}.  We also note that in addition to gross profile variability that occurs on timescales of years, some objects also exhibit sporadic variability on timescales of only a few months; these variations may be related to a phenomenon that occurs on the dynamical time scale.  The light crossing time of the disk of 3C~390.3 computed from equation (\\ref{tl}) is in the range of 11 to 37 days \\citep[using the inner and outer radii of the line-emitting disk from][]{eh94}.  This affords an immediate test of the disk interpretation of the double-peaked lines. A reverberation mapping campaign by the {\\it International AGN Watch} \\citep{d98} yielded a lag between the continuum and Balmer line variations of approximately 20~days, which is in good agreement with the above value.  The Eddington ratios were computed by comparing the total bolometric luminosity to the Eddington luminosity ($\\ledd = 1.38\\times10^{38}\\;M/\\msol\\; {\\rm erg\\;s^{-1}}$). We used the available Spectral Energy Distributions (SEDs) and integrated bolometric luminosities for M~81 and NGC~4579 \\citep{ho99} and NGC~1097 \\citep{n04,n05}. \\citet{ehc03} presented an SED for Arp~102B, and we adopt their SED, excluding the IR bump and adjusting the specific luminosity in the X-ray band.  \\footnotemark\\footnotetext{ We adopted $\\log[\\nu\\;L_{\\nu}(1\\;{\\rm keV})]=42.331$ for consistency with earlier observations with the {\\it ROSAT} and {\\it Einstein} X-ray observatories \\citep{h97,bier81}.  During the {\\it ASCA} observation \\citep{ehc03} the X-ray flux was found to be a factor of 3 higher than the earlier observations and was attributed to an unusually high state. We retained the 2--10~keV spectral slope measured from the {\\it ASCA} X-ray spectrum.} An SED for 3C~390.3 was complied from various sources in the literature, and the full references are provided in Table~\\ref{results_tab}. There are no SEDs for the remaining objects and we estimated their bolometric luminosities using their measured X-ray luminosities and assuming a value for $L_{2-10\\;{\\rm keV}}$/$L_{\\rm bol}$, as given in Table~\\ref{results_tab}. For IRAS~0236.6--3101, 1E~0450.3--1817, and NGC~4203; we adopted the average $L_{2-10 {\\rm keV}}$/$L_{\\rm bol}$ for the LINERs in this study (Arp~102B, NGC~1097, M~81, and NGC~4579). For Pictor~A and PKS~0921--213 we adopted the radio-loud quasar SED of \\citet{e94} since this is the most likely SED for these objects (had we adopted the LINER-like SED, even the lower limit to the Eddington ratio would be larger than $10^{-2}$). To compute $L_{bol}$, the SEDs were integrated from 1~MHz to 100~keV, using power-law interpolations between data points.  We note that the 2--10~keV flux of an AGN is variable by a factor of a few, which implies an additional uncertainty to the bolometric luminosities estimated from the X-ray luminosity.  The Eddington ratios listed in Table~\\ref{results_tab} span a wide range and there is a general positive trend between bolometric luminosity and Eddington ratio, which we illustrate in Figure~\\ref{lbol_eddrat}. However, we do not consider this trend convincing because the quantities plotted are not independent of each other ($L_{bol}/\\ledd\\propto L_{bol}/M_{BH}$).  The trend is merely a consequence of the fact that the estimated black hole masses span only a factor of 10 (a factor of 2, if 3C~390.3 is excluded), while the bolometirc luminosities span a factor of $10^4$. The existence of such a trend is also called into question by the fact that 3C~390.3 has $L_{bol}$ that is two orders of magnitude larger than IRAS~0236.6--3101 and 1E~0450.4--1817, yet its Eddington ratio is similar.  While many double-peaked emitters may harbor radiatively inefficient accretion flows, the accretion flow in some objects, such as Pictor~A and PKS~0921--213, are probably more similar to those found in Seyfert 1s. A vertically extended structure may still be required in these objects to illuminate the outer accretion disk and may take the form of a spherical corona such as that used by \\citet{csd90}, a beamed corona \\citep{b99a,mbp01}, and/or the base of the jet \\citep*{mff01}.  That the double-peaked emitters in even this small sample has a wide range Eddington ratios is not completely surprising. The double-peaked emitters found in the SDSS by \\citet{s03} are very heterogeneous. Although 12\\% were classified as LINERs, many others had properties which were indistinguishable from the general $z<0.332$ AGN population. We note that our results confirm the {\\it general} findings of \\citet{wl04}, who estimated black hole masses for both the \\citet{eh94} and \\citet{s03} samples through a series of correlations obtained from reverberation mapping of AGNs \\citep{k00}. These authors found that (1) double-peaked emitters are a heterogeneous group with a wide range in Eddington ratios and (2) those objects with higher bolometric luminosities tend to have larger Eddington ratios. However, the black hole masses of individual double-peaked emitters deduced through that method can be very inaccurate, as acknowledged by these authors. Although the black hole masses for 1E~0450.3--1817 and Arp~102B inferred by these authors are consistent with those obtained through the \\msig\\ relationship, the black hole masses for Pictor~A and IRAS~0236.6--3101 were overestimated by an order of magnitude. Thus we caution that for the purposes of modeling and interpreting individual objects, it is necessary to obtain black hole masses through a method which is more well tested and well calibrated in the range of BH masses relevant to our objects, such as the \\msig~relationship."
    },
    "0601/astro-ph0601351_arXiv.txt": {
        "abstract": "{We come back to the Cauchy integral equations occurring in radiative transfer problems posed in finite, plane-parallel media with light scattering taken as mono\\-chromatic and isotropic. Their solution is calculated following the classical scheme where a Cauchy integral equation is reduced to a couple of Fredholm integral equations. It is expressed in terms of two auxiliary functions $\\zeta_+$ and $\\zeta_-$ we introduce in this paper. These functions show remarkable analytical properties in the complex plane. They satisfy a simple algebraic relation which generalizes the factorization relation of semi-infinite media. They are regular in the domain of the Fredholm integral equations they satisfy, and thus can be computed accurately. As an illustration, the $X$- and $Y$-functions are calculated in the whole complex plane, together with the extension in this plane of the so-called Sobouti's functions} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601167_arXiv.txt": {
        "abstract": "The present paper focuses on the much debated Butcher-Oemler effect: the increase with redshift of the fraction of blue galaxies in clusters. Considering a representative cluster sample made of seven group/clusters at $z \\sim 0.35$, we have measured the blue fraction from the cluster core to the cluster outskirts and the field  mainly using wide field CTIO images. This sample represents a random selection of a volume complete x-ray selected cluster sample, selected so that there is no physical connection with the studied quantity (blue fraction), to minimize observational biases. In order to statistically assess the significance of the Butcher--Oemler effect, we introduce the tools of Bayesian inference. Furthermore, we modified the blue fraction definition in order to take into account the reduced age of the universe at higher redshifts, because we should no longer attempt to reject an unphysical universe in which the age of the Universe does depend on redshift, whereas the age of its content does not. We measured the blue fraction from the cluster center to the field and we find that the cluster affects the properties of the galaxies up to two virial radii at $z\\sim0.35$. Data suggest that during the last 3 Gyrs no evolution of the blue fraction, from the cluster core to the field value, is seen beyond the one needed to account for the varying age with redshift of the Universe and of its content. The agreement of the radial profiles of the blue fraction at $z=0$ and $z\\sim0.35$ implies that the pattern infall did not change over the last 3 Gyr, or, at least, its variation has no observational  effect on the studied quantity. ",
        "introduction": "The nature and the time scale of the processes that shape galaxy properties in clusters and groups are still unclear. The presence of a hot intercluster gas observed in X-rays should have a role in shaping some galaxy properties (e.g. Gunn \\& Gott 1972). The window opened by the redshift dependence of the galaxy properties has been used to set constraints on the time scales of the processes (e.g. Butcher \\& Oemler 1978, 1984; Dressler et al. 1997; Stanford, Eisenhardt \\&  Dickinson, 1998; Treu et al. 2003). However, the observational evidence of the environmental effect is still uncertain. For example, the existence of a Butcher--Oemler (BO) effect (Butcher \\& Oemler 1984), i.e. the fact that clusters at higher redshift have a larger fraction of blue galaxies, $f_b$, is still controversial. The controversy is raised by two criticisms concerning measurements and sample.  Andreon, Lobo \\& Iovino (2004; hereafter ALI04) analyse three clusters at $z\\sim0.7$ without finding evidence of a high blue fraction with respect to $z\\sim0$. They also show the drawbacks of the various definitions of $f_b$ adopted in the literature. They conclude then that ``twenty years after the original intuition by Butcher \\& Oemler, we are still in the process of ascertaining the reality of the effect\". The same work put in a different perspective the results of Rakos \\& Shombert  (1995), clarifying the fact that even if all the galaxies in the Universe are passively evolving, the blue fraction will be $f_b\\approx 1$ at $z\\ga0.7$ in the Rakos \\& Shombert (1995) scale. Therefore, the very high fraction they found at high redshift does not require any special mechanism to account for the present day counterparts other than ageing. ALI04 introduce also a first discussion about the difficult task of measuring the error on $f_b$, given the observations, showing that at least some previous works have underestimated errors and, by consequence, overstated the evidence for the BO effect. The role of the inference, the logical step going from the observed data to the true value and its error, has been further elaborated in D'Agostini (2004) in the more general case of unknown individual membership for the galaxies.   Kron (1994) claimed that all the ``high\" redshift clusters known in the early  80's  ($z\\approx 0.3-0.5$) were somewhat extreme in their properties, and this was precisely the reason why they were detected. Andreon \\& Ettori (1999) quantify this issue, and show that many of the clusters compared at different redshifts have also different masses (or X-ray luminosities), in such a way that ``we are comparing unripe apples with ripe oranges in understanding how fruit ripens\" (Andreon \\& Ettori 1999). Together with  Allison-Smith et al. (1993) and Andreon \\& Ettori (1999), ALI04 show that the optical selection of clusters is prone to produce a biased - hence inadequate - sample for studies on evolution since at larger redshifts it naturally favours the inclusion in the sample of clusters with a significant blue fraction. They show that clusters with a blue fraction as the observed ones are over-represented in optical cluster catalogs by a factor two, with respect to identical clusters but without a bursting population.     There is therefore a compelling need to study the properties of galaxies in clusters at intermediate redshift ($z\\approx0.35$), avoiding the bias of an optical selection, by choosing clusters of the same mass as present day studied clusters to avoid an ``apple vs orange\" issue. This is the aim of this paper, where we present a BO--style study of a small but representative sample of 7 clusters, X-ray selected, of low to average mass (velocity dispersion) and at intermediate redshift.   The layout of the paper is the following. In sect. 2 we present optical imaging  and spectral data. In sect. 3 we show that the studied sample is both representative and X--ray selected. We revisit in section 4 the definition of the blue  fraction, in order to account for the reduced age of the universe at higher redshift. Sect. 5 presents some technical details. Results are summarized in Sect. 6, whereas Sect. 7 discusses relevant results published in the literature and some final conclusions. Appendixes present a Bayesian estimate of cluster velocity dispersion, richness, and blue fraction.  We adopt $\\Omega_\\Lambda=0.7$, $\\Omega_m=0.3$ and $H_0=70$ km s$^{-1}$ Mpc$^{-1}$.   \\begin{table} \\caption{The cluster sample} \\begin{tabular}{l r @{.} l r r r r r} \\hline Name & \\multicolumn{2}{c}{z} & $N_{z}$ & $\\sigma_v$ & error & $r_{200}$ \\\\ \\multicolumn{2}{c}{} & & & km/s & km/s & Mpc \\\\ \\hline XLSSC 024  & 0&29  & 11 &  430 & 96 & 1.0 \\\\  % XLSSC 028  & 0&30  &  8 &  376 & 98 & 0.8 \\\\ % XLSSC 009  & 0&33  & 12 &  236 & 52 & 0.5 \\\\ % XLSSC 010  & 0&33  & 11 &  367 & 96 & 0.8 \\\\ % XLSSC 016  & 0&33  &  5 &  915 & 294 & 2.0 \\\\ % XLSSC 006  & 0&43  & 39 &  837 & 106 & 1.7 \\\\ % XLSSC 012  & 0&43  & 12 &  741 & 165 & 1.5 \\\\ % \\hline \\end{tabular} \\hfill \\break { } \\end{table} ",
        "conclusions": "\\subsection{Comparison with previous works}  Comparison of our results with other ones requires to pay attention to the prescriptions adopted to define the blue fraction, to the way the cluster sample is built, and, sometime, to the adopted statistical approach.      \\subsubsection{Evolution of the blue fraction}  The most similar work to ours is the seminal Butcher \\& Oemler paper, from which we modelled our prescriptions. By selecting a small subsample of clusters of richness similar to ours but located in the nearby universe ($z\\sim0.02$, where BO and our prescriptions are identical) they find $f_b$ values within $r_{30}$, the radius that includes 30 per cent of cluster galaxies, in the $0.02$ to $0.19$ range, with a typical error of $\\pm0.03$. This range of values is not significantly lower, considering the various sources of uncertainties, than our central value $f_b=0.24\\pm0.04$, to claim that the two values are different at a high significance level, especially taking into account the fact that the Butcher \\& Oemler errors are sometimes optimistically estimated (Andreon, Lobo \\& Iovino 2004).   de Propris et al.'s (2004) large sample of nearby clusters matches our sample in terms of richness: we find for their sample\\footnote{We thank R. de Propris for giving us their blue fraction within $r_{200}/2$} an average $N_{gal}$ of 30 galaxies and a blue fraction inside $r_{200}/2$ of $0.17$, taking into account that the blue fraction is a binomial deviate (the authors assumed it to be a Gaussian and find $f_b=0.13$). The error due to the sample size is negligible (0.01) because their sample is large. However, the largest source of uncertainty in their work comes from their large photometric errors (they use photographic plates). Such photometric errors induce a bias in the blue fraction that, as discussed in Sect 3.1, is difficult to correct for (Jeffreys 1938, Eddington 1940), and is neglected by the authors.  After accounting for minor differences in the luminosity cuts between de Propris et al. and BO and for the mentioned Malmquist bias, the estimated blue fraction within $r_{200}/2$ in the de Propris et al.'s (2004) sample becomes $\\approx 0.25$, but with an error hard to quantify. Inside $r_{200}/2$ we find  $f_b=0.26\\pm0.04$, which identical to what found in the de Propris et al.'s large nearby sample.  To summarize, our $z\\sim0.35$ sample matches in terms of richness the nearby samples in Butcher \\& Oemler (1984) and, especially, de Propris et al. (2004) and shows equal blue fractions within $r_{30}$ and $r_{200}/2$, i.e. no Butcher-Oemler effect is seen. It is to be noted that our sample has an almost identical size and redshift distribution as the high redshift clusters in the BO sample, and thus that our lack of detection of a BO effect is not due to a smaller or closer sample.  The compared clusters matches in terms of richness, but are constructed using different selection criteria, because the low redshift sample is an optically selected one, while our cluster sample is x-ray selected. As mention in sec. 1, our x-ray selection is chosen to minimize the observational bias on $f_b$, and hence to derive a fair measure of the blue fraction. At low redshift, we are not aware of any reason why $f_b$ should be biased {\\it at a fixed richness} for an optically selected sample, such as the ones of Butcher \\& Oemler (1984) and de Propris  et al. (2004): why clusters of a given richness and rich  in blue galaxies should be over/underrapresented in cluster catalogs of the nearby universe? Therefore, even if the selection criteria used to build the compared cluster samples are different, the comparison of the blue fractions is safe, because both cluster samples provide unbiased values of $f_b$.  There are hints that confirm the constancy of the blue fraction at even larger redshifts. ALI04 show evidence for a {\\it low} blue fraction at $z\\sim0.7$. Recently, Tran et al. (2005) also find a {\\it low} blue fraction ($f_b=0.13$) for a cluster at $z\\sim0.6$, computed inside a cluster portion that, if not rigorously identical to the one prescribed by Butcher--Oemler, does support the {\\it non} existence of a Butcher--Oemler effect.  Both works adopted a non-evolving $\\Delta$. If an evolving $\\Delta$ is used, the derived blue fraction at high redshift would even be lower than claimed, giving further support to our conclusion. ALI04 have also disproved all the reported literature evidence accumulated thus far for the existence of a Butcher--Oemler effect, i.e for a change of the blue fraction inside $r_{30}$.  All the above suggests that the fraction of blue galaxies, computed by separating the galaxies using a population formed by stars whose age increases at the same rate as the universe age increases, does not evolve. Or, if the reader prefers, there is no systematic drift from the blue to the red classes of galaxies as the look--back time evolves, between $z\\sim0$ and $z=0.44$.    \\medskip   A result similar to the one depicted in Fig. 8 is presented in Lewis et al. (2002) based on nearby clusters. Our results are in qualitative agreement with theirs since we also find that the cluster affects the fraction of active galaxies up to the virial radius. Lewis et al. (2002) have studied a {\\it nearby} cluster sample composed of 440 member galaxies inside the virial radius (vs our sample of 320). Figure 8 shows that their fraction of galaxies with star formation rates, normalized to $M^*$, larger than 1 solar mass per year (triangles) nicely compares with our derived fraction of blue galaxies. Our error bars for their points show the expected central 68 per cent credible intervals, only accounting for sampling errors, computed by us from a straightforward application of statistics. In Lewis et al. (2002), the sample is split in classes very similar to ours and BO: in fact, our spectro-photometric Sa has a star formation rate, normalized to $M^*$, equal to their adopted threshold (one solar mass per year per $M^*$ galaxy) if $M^*=8.2 \\ 10^{10} M_{\\odot}$, a value well inside the range of values usually observed (e.g. Blanton et al. 2001; Norberg et al. 2002). I.e. what is called blue by them is also called blue by us, on average. It is not surprising, therefore that integrating the Lewis et al.'s (2002) blue profile within $r_{200}/2$ gives a blue fraction identical to the one observed in the Propris  et al.'s (2004) sample ($0.26$ vs $0.25$), further supporting the similarity of the two classes (blue by colour and blue by star formation rate).   Their cluster sample has an overlapping, but different, range of masses (velocity dispersion) with respect to our sample: our richest clusters have a velocity dispersion typical of the average values of Lewis et al. (2002) clusters. However, their profile is only marginally affected, if at all, by separating clusters in (two) velocity dispersion classes (Lewis et al. 2002). Furthermore, Gomez et al. (2003) indirectly confirm that the radial profile is not too much affected by differences in cluster mass, by studying a sample of nearby clusters having velocity dispersions similar to our sample. Therefore, differences in the way cluster sample are built seems not to affect the derived ``blue\" profile.  The agreement between Lewis et al. (2002) and our blue fraction profiles is almost perfect; however the however the studied clusters are located at quite different look back times:  clusters in Lewis et al. are in the very nearby universe (at $z\\sim0.07$), whereas our clusters have $z\\sim0.35$, implying a $\\sim3$ Gyr time difference for the adopted cosmology. As long as the separation of galaxies in classes by Lewis et al. (2003) and in our work is similar, the agreement of the two radial profiles means that there is no evolution of the blue fraction between $z\\sim0$ and $z\\sim0.35$, from the cluster center to the field value.      \\subsubsection{Disagreements or different ways in interpreting the data?}  Considering a much more luminous X-ray (and therefore massive) sample of clusters at intermediate redshift, Fairley et al. (2002) find an increasing blue fraction as a function of the clustercentric distance, up to 2 $r_{30}$, in good agreement with the results we found over much larger clustercentric distances. A quantitative comparison between the two pieces of work is however impossible, because there are too many uncontrolled variables that are allowed to change between these.  The authors find a blue fraction $f_b\\approx0.2 \\pm0.1$, in agreement with the value we observe in the cluster center $f_b=0.24\\pm0.04$. However, we believe that this apparent agreement  largely arises by chance. First, the authors used a non-evolving template to separate the galaxies in red and blue classes, and find two sets of (different) values for their two sets of available colours. Secondly, they considered higher redshift than we do, by observing clusters with comparable exposure times but with smaller telescopes (2.5m vs 4.0m). In spite of their expected larger errors, they neglect the effect of photometric errors on their blue fraction estimates (sect 3.2). Third, they do not adopt an evolving luminosity limit. And, finally, a comparison of the values derived in the two works requires an extrapolation, because clusters with very different masses (X--ray luminosities) are considered.   Ellingson et al. (2001) performed a study quite different from ours, and adopt a galaxy separation that is the same irrespective of redshift (i.e. of galaxy age), because they decomposed their spectra on non-evolving spectral templates. Their claim for a change in the population gradient is just a restatement of the fact that the blue fraction is higher everywhere in the cluster and in the field because galaxies are bluer when they were younger, i.e. is not informative about processes running in clusters or in the field, but just informative about aging. These authors would observe an evolution of the gradient even if galaxies would be kept isolated from the surrounding environment and the infall in the cluster would be fixed (i.e. no new galaxy falls in the cluster, and galaxies are kept fixed at their observed position): their ``old population\" fraction increases going from high to low redshift  because galaxies become older, and the effect is more marked  at large clustercentric radii than in the centre,  because in the cluster core the  ``old population\" fraction is already near to 1 and cannot take values larger than 1. We, instead, choose to reduce by one the number of parameters, removing the age dependency by using an evolving (Sa) template.    In summary, the referenced analysis do not reveal results in disagreement with our own work, although their interpretation may sometimes be different (or even opposite to ours).   \\subsection{Conclusions}  This paper revises the definition used to separate galaxies in two colour classes in a way that takes in to account the reduced age of the Universe at higher redshift.  It is nowadays uninteresting to know whether the fraction of blue  galaxies changes with redshift in a way that is different from the expectation of a model that we now know is unphysical (that has the same age at all redshifts). If the model is unphysical, there is no need to make observations to rule it out. A stellar population whose age does not change in a Universe whose age instead changes, as it is supposed by using a non-evolving spectral template  (or a fixed $\\Delta$, i.e. the BO prescription), is clearly non-physical. It was useful a long time ago to show that a universe with the same age at all redshifts is rejected by observations. However, nowadays we can attack a more essential question: to know whether galaxies evolve differently from a reference evolution that is physically acceptable. Our measurements of evolution are, therefore, zero-pointed on the evolution of an object whose age increases as required by the current cosmological model. We select a  spectro-photometric Sa to conform to the BO prescription in the local universe.  Effectively, this is a change in perspective: we should no longer attempt to reject an unphysical universe, in which the age of the Universe does depend on redshift, whereas the age of its content does not, but we should study whether the observed differences between the low and high redshift content are in agreement with differences of the Universe age at the considered redshifts.  Furthermore, we have introduced in our specific domain the tools of Bayesian inference (see Appendix), dramatically improving on previous approaches that led some authors to claim that they have observed unphysical values (such as blue fractions outside the $[0,1]$ range or negative star formation rates). Such tools allow us to use all our data without rejecting blue fractions measured at large clustercentric radii, where the signal to noise is low, contrary to previous researchers obliged to discard such data (or claiming that they have observed unphysical values).   The main result of this work is that {\\it we find that the cluster affects the properties of the galaxies up to two virial radii at $z\\sim0.35$}.  We have measured the blue fraction of a {\\it representative} sample of clusters at intermediate redshift. Indeed, our sample is a random sampling of a volume complete X-ray selected cluster sample. The X-ray selection has no cause-effect relationship on the cluster blue fraction, all the remaining parameters being kept fixed, to the best of our knowledge, and, therefore, the studied sample consists of an unbiased one (from the blue fraction point of view). Our statement should not be over-interpreted, however, because we are only sampling a portion of the X-ray parameter space: very X-ray luminous clusters are missing in our sample because they are intrinsically rare, and clusters with fainter X-ray emission than the limiting flux are missing because they lie outside the sampled space.  Studied clusters show a variety of values for the blue fraction, when the fraction is measured within $r_{200}$. The variety is too large to solely be accounted for by errors. At smaller radii, instead, the blue fractions are more homogeneous. Actually, there is some evidence that the blue fraction within $r_{200}$ increases with the cluster velocity dispersion, i.e. with the cluster mass, whereas the increase at smaller radii is much smaller, if present at all. Therefore, intermediate redshift clusters with the largest masses show the largest fractions of star-forming galaxies, when measured within $r_{200}$. However, the evidence is good but not definitely conclusive and still requires an independent confirmation.  The radial dependence of the blue fraction is quite shallow: it smoothly and monotonically increases from the centre to the field. The latter has been determined according to our prescriptions using COMBO-17 data.  The radial dependence (i.e. the blue fraction at every computed clustercentric radius) is equal to the one recently found in a comparable sample of clusters, but in a 3 Gyr older universe, i.e. at $z\\sim0$ (Lewis et al. 2002). The agreement between the two derived profiles (amplitude and shape), our blue fraction within $r_{30}$ and $r_{200}/2$ and the local similar determinations (Butcher \\& Oemler 1984, de Propris et al. 2004), the low blue fractions at high redshift (ALI04, Tran et al. 2005), all of these suggest that there is no colour evolution beyond the one needed to account for the different age of the Universe and of its content. The above is found to hold from the cluster core to the field value. Previous controversial evidence from the literature assumed that the universe becomes older while its content does not, and overstated the significance of the evidence or compared heterogeneously measured blue fractions (as shown in ALI04).       The interpretation of the observed radial trend is complicated by the possible existence of a backsplash population. If the backsplash population represents a negligible fraction of galaxies at a given clustercentric radius, then the large clustercentric distance at which the cluster still produces some effect rules out short--range scale mechanisms. However, the predicted backsplash population is precisely what is needed to explain our observed fraction at $r_{200}$, given the fraction at the cluster center and in the field, and also qualitatively accounts for different kinematics of galaxies having different star formation rates (blue and red spirals) in the Coma cluster. If this is the case, mechanisms efficient in the cluster center only come into play, because galaxies are affected when they reach the cluster core, and are then scattered at large clustercentric radii where they spend a lot of time and are observed.  The possible existence of the backsplash population does not offer us the possibility to draw a final inference about the nature and the time scale of the processes that shape galaxy properties in clusters. The backsplash mechanism is a physical one: it affects the interpretation of measured radial (or density) trends drawn by us and other authors, and forces us to keep in hold their interpretations. However, the infall pattern  turns out not to have changed during the last 3 Gyr, as measured by the identical blue fraction profiles at $z\\sim0$ and $z\\sim0.35$, in spite of apparently contradictory previous claims, based on the use of a fraction definition that has one more (uncontrolled) dependence."
    },
    "0601/astro-ph0601484_arXiv.txt": {
        "abstract": "{ In the past, different works based on numerical simulations have been presented to explain magnetic fields (MFs) in the large scale structure and within galaxy clusters.  In this review, I will summarize the main findings obtained by different authors and - even if many details are still unclear - I will try to construct a consistent picture of our interpretation of large-scale magnetic fields based on numerical effort. I will also sketch how this is related to our understanding of radio emission and summarize some arguments where our theoretical understanding has to be improved to match the observations. ",
        "introduction": "There are several different approaches to study the build-up of MFs in the intergalactic medium. Many papers consider  MHD simulations of cloud-wind interaction (see \\citealt{2000ApJ...543..775G} and references therein) or are simulating the rise of relic radio bubbles (see \\citealt{2005ApJ...624..586J}, \\citealt{2005MNRAS.357..242R} and references herein). Such work is usually focused more on the relevance of the local MFs within these processes, whereas in this review I will concentrate more on the discussion of MHD simulations which aim at understanding the build-up of the cluster MF and its possible origin. So far, firm evidence for the presence of extended MFs has been found only in galaxy clusters. For recent reviews see \\citet{Carilli.Taylor..2002} and \\citet{2004IJMPD..13.1549G}. Recently there are observational claims of a detection of substantial MFs also within the large scale structure (LSS), see contribution by Kronberg. The origin of the MF in galaxy clusters is still under debate. The variety of possible contributors ranges from primordial fields, battery and dynamo fields over all classes of astrophysical objects {which contribute with their ejecta.  The later possibility is supported by the observation of the metal enrichment of the intracluster medium (ICM), which has to be originated as ejecta of astrophysical objects. In addition, the MFs produced by all these contributors will be compressed and amplified by the process of structure formation. The exact amount of this amplification and the resulting MF filling factor will depend on place and time at which the contributer is thought to be most efficient. ",
        "conclusions": "\\label{sec:discussion}  It seems that within the last years a probably consistent picture of MFs arises from numerical works and observations.  Supported by simulations of individual events/environments like shear flows, shock/bubble interactions or turbulence/merging events, a super-adiabatic amplification of MF is predicted.  It is worth to notice that this common finding is obtained by using a variety of different codes, based on different numerical schemes.  Further support for such super-adiabatic amplifications comes from analytical estimates of anisotropic collapse making use of the Zel'dovich approximation.  When applied in fully consistent cosmological simulations, various observational aspects are reproduced (see Figs. \\ref{fig:RMprof} and \\ref{fig:Lx_RM}) and the resulting MF amplification reaches a level sufficient to link models predicting seed MFs at high redshift with the MFs observed in galaxy clusters today. It is important to mention that all simulations show that this effect increases when the resolution is improved, and therefore all the numbers have to be taken as lower limits of possible amplification.  \\begin{figure} \\includegraphics[width=0.49\\textwidth]{dolag_f10.eps} \\caption{{\\small The MF strength as a function of baryonic overdensity. The long-dashed line shows the expectation for a purely adiabatic evolution, the solid line gives the mean field strength at a given overdensity within a cosmological simulation (\\citealt{2005JCAP...01..009D}). While the latter is close to the adiabatic value in underdense regions, there is a significant inductive amplification in clusters due to shear flows and turbulence, subject however to saturation in the cluster cores. At any given density, a large fraction of particles remains close to the adiabatic expectation, as shown by the dotted line, which gives the median of the distribution at each density.}} \\label{fig:b_ampli} \\end{figure}  Note that on the other hand there are significant differences for the predictions of the MF structures coming from different models of seed MFs. In particular there are main differences between the up-scaled, cosmological battery fields (\\citealt{2001ApJ...562..233M,2004PhRvD..70d3007S}) and the MF predicted from high-resolution simulations of galaxy clusters using either AMR (\\citealt{2005ApJ...631L..21B}) or SPH (\\citealt{1999A&A...348..351D,Dolag:2002,2005JCAP...01..009D}). In the latter case it is possible to follow the amplification of seed fields within the turbulent ICM in more detail. A good visual impression can be obtained by comparing the regions filled with high MFs shown in Figs. \\ref{fig:Sigl_I} and \\ref{fig:Brueggen_I}. It is clear that the high MF regions for the battery fields are predicted to be much more extended, leading to a flat profile around the forming structure, whereas for the turbulent amplified MFs the clusters show a much more peaked MF distribution. Part of this difference originates from the physical model behind, as the cosmological shocks are much stronger outside the clusters than inside.  Somewhat more unclear is what the contribution of the different numerical resolutions of these simulations to these discrepancies is. Calibrating such simulations using the MF measured only in the high-density regions of galaxy clusters makes it crucial to perform a more detailed comparison with all available observations. Note that extrapolating the predictions of the simulations into lower density regions, where no strong observational constraints exist, will even amplify the differences in the predictions for the MF structure between the different simulations.  Moreover, one has to keep in mind that depending on the ICM resistivity the MF could be suffering decay, which so far is neglected in all the simulations discussed before.  Furthermore a clear lack of the present simulations is that they do not include the creation of MF by all the feedback processes happening within LSS (like radio bubbles inflated by AGNs, galactic winds, etc.): this might alter the MF prediction if their contribution turns out to be significant. Also all the simulations done so far neglect radiative losses: if included, this would lead to a significant increase of the density in the central part of clusters and thereby to a further MF amplification in these regions.  Finally, there is an increasing number of arguments suggesting that instabilities and turbulence on very small scales can amplify the MFs in a relatively short timescale reaching the observed $\\mu$G level. However it is still unclear how such very small scale fields can be ordered on large (up to hundreds of kpc) scales as observed.  \\begin{figure}[t] \\includegraphics[width=0.49\\textwidth]{dolag_f11.eps} \\caption{{\\small Comparison of RMs from the simulation with observations for Abell clusters, as a function of distance to the closest cluster. Smooth lines represent the median values of $|{\\rm RM}|$ produced by simulated clusters with masses above $5 \\times 10^{14}~{\\rm M}_\\odot$ and $3 \\times 10^{14}~{\\rm M}_\\odot$. The broken line represents the median of combined data taken from the independent samples presented in \\citet{1991ApJ...379...80K} and \\citet{2001ApJ...547L.111C}. We also include data (diamonds) for the three elongated sources observed in A119 (\\citealt{1999A&A...344..472F}), and for the elongated source observed in the Coma cluster (\\citealt{1995A&A...302..680F}).}} \\label{fig:RMprof} \\end{figure}  Concerning Radio haloes and relics the picture is only partially consistent ( for a more detailed discussion of primary and secondary models see \\citealt{2004JKAS...37..493B} and references therein). Radio relics seem to be most likely related to strong shocks produced by major merging events and therefore produced by direct re-acceleration of CRs, the so-called primary models. Although some of the observed features like morphology, polarization and position with respect to the cluster centre can be resonably well reproduced, there might be still some puzzles to solve. On one hand, direct acceleration of CRs in shocks seems to overestimate the abundance and maybe the luminosity of radio relics; on the other hand simulations which illuminate fossil radio plasma can produce reasonable relics only starting from a small range of parameter settings.  A similar situation arises for modelling the central radio emission of galaxy clusters by secondary models. On one hand, the total luminosity seems to be reproduced using reasonable assumptions and also the steep observed correlation between cluster temperature/mass and radio power seems to be reproduced quite well. But such models suffer from two drawbacks. The first one is that in the framework of such models every massive cluster produces a powerful radio halo, but this is not confirmed by observations (see Fig. \\ref{fig:pt}).  The other problems is that the detailed radio properties are not reproduced. Firstly, the profile of radio emission in most cases is too steep, so that these models cannot reproduce the size of the observed radio halos in almost all cases. Secondly, the observed spectral steepening (e.g. \\citealt{2005A&A...440..867G}) also cannot be reproduced. Note that there is also no indication from observations that clusters showing radio emission contain higher MFs than the ones without observable extended, diffuse radio emission. On the contrary, the cluster A2142 has a MF strength similar to the Coma cluster (see Fig. \\ref{fig:Lx_RM}), but the upper limit on its radio emission is at least two orders of magnitude below the value expected from the correlation (see Fig. \\ref{fig:pt}).  Note that both clusters are merging systems characterized by the presence of two central cD galaxies. This indicates that there should be further processes involved or additional conditions to fulfill to produce radio emission. It is worth to notice that recent models, based on turbulent acceleration, seem to overcome this problem (see \\citealt*{2003ApJ...594..732K,2005MNRAS.363.1173B,2005MNRAS.357.1313C} and references therein). Note that \\citet{2005MNRAS.357.1313C} predicted the probability for a galaxy cluster to show giant radio halo to be an increasing function of cluster mass and reproduced the observed fraction of $\\approx 30$\\% for massive galaxy clusters.   There is a big challenge for the next generation of cosmological MHD simulations. Simulations are quite close to having the resolution necessary to properly describe the MF components down to the observed scales, but on the other hand this means to resolve galaxies inside the simulations as well.  Note that following the dynamics of such structures for which the MF might dominate the evolution, is a real challenge within LSS simulations. However, we expect that - if succeeding in overcoming these limitations - the dynamical impact of the MF on regions like the cooling flows at the centres of galaxy clusters will be significant and will contribute to solve these outstanding puzzles.  \\begin{figure}[t] \\includegraphics[width=0.49\\textwidth]{dolag_f12.eps} \\caption{{\\small Correlation of rotation measurements with X-ray surface brightness for different galaxy clusters. The data points are observations collected from the literature. The colors represent the virial cluster temperature taken from literature. This is an updated version of the correlation presented in \\citet{Dolag:2001}.}} \\label{fig:Lx_RM} \\end{figure}"
    },
    "0601/hep-th0601105.txt": {
        "abstract": "We discuss the cosmological evolution of the inflationary gravitational wave background (IGWB) in the Randall-Sundrum single-brane model. In braneworld cosmology, in which three-dimensional space-like hypersurface that we live in is embedded in five-dimensional anti de Sitter (AdS$_5$) spacetime, the evolution of gravitational wave (GW) modes is affected by the non-standard expansion of the universe and the excitation of the Kaluza-Klein modes (KK-modes). These are significant in the high-energy regime of the universe.  We numerically evaluate these two effects by solving the evolution equation for GWs propagating through the AdS$_5$ spacetime. Using a plausible initial condition from inflation, we find that the excitation of KK-modes can be characterized by a simple scaling relation above the critical frequency $f_{\\rm crit}$ determined from the length scale of the fifth dimension $\\ell$. The remarkable point is that this relation generally holds as long as the matter content of the universe is described by the perfect fluid with the equation of state (EOS) $p=w\\rho$ for $0\\leq w\\leq 1$. The resultant scaling relation is translated into the energy spectrum of the IGWB as $\\Omega_{\\rm GW}\\propto f^{(3w-1)/(3w+2)}$ for $f>f_{\\rm crit}$. This indicates that in the radiation dominant case ($w=1/3$), the two high-energy effects accidentally compensate each other and the spectrum becomes almost the same as the one predicted in the four-dimensional theory, i.e., $\\Omega_{\\rm GW}\\propto f^0$. ",
        "introduction": "\\label{sec:introduction} %===========================================================================% %===========================================================================%  Gravitational waves (GWs) are ultimate probes of the untouched region of the universe. Currently, large scale ground-based interferometers (TAMA300\\cite{Ando2005}, LIGO\\cite{Sigg2004}, VIRGO\\cite{Acernese2005}, GEO600\\cite{Grote2005}) are enthusiastically trying to detect the signals emitted from stellar objects with relativistic motion -- supernovae explosions, coalescence of neutron star binaries and so on. Among numerous types of GWs, the gravitational wave background (GWB) may possess much interesting information on the cosmology, though its detection is expected to be so challenging \\cite{Abbott2005GWB}. Especially, the inflationary GWB (IGWB), generated during the inflationary epoch by the quantum fluctuations of the spacetime\\cite{Starobinsky1979, Rubakov1982, Fabbri1983, Abbott1984a, Abbott1984b, Abbott1986, Allen1988, Turner1997, Maggiore}, is thought to be one of the most fundamental predictions of the inflationary scenario \\cite{Sato1981, Guth1981, Linde1982, Albrecht1982}. Since the history of the cosmological expansion is imprinted in the power spectrum of the IGWB, it helps us to understand the extremely early universe if we can detect the signals by the future space-based experiments, such as DECIGO\\cite{DECIGO} and BBO\\cite{BBO, Ungarelli2005}.  As an ultimate cosmological tool, the IGWB may also be useful to probe the presence of extra-dimensional spaces. Recent developments in particle physics suggest that we live in a higher dimensional spacetime. In particular, braneworld scenarios have recently attracted much attention theoretically and observationally (for a review, see \\cite{rev_Maartens}). According to them, we live in a three-dimensional hypersurface (brane) embedded in the higher-dimensional spacetime (bulk). While gravity can propagate in the bulk with the curvature scale $\\ell$, the standard model particles are confined to the three-dimensional brane. In the low-energy regime of the universe ($H\\ell \\ll 1$), four-dimensional general relativity is successfully recovered and the extra-dimensional effects should be fairly small. On the other hand, in the high-energy regime ($H\\ell \\gg 1$), the localization of gravity is not always guaranteed and a significant deviation of the time evolution of GWs from the standard four-dimensional theory is expected. If this scenario is true, the spectrum of the IGWB may be significantly modified by the high-energy effects, which can provide a direct probe of the extra-dimensions.  The goal of this paper is to investigate the high-energy effects on the evolution of the IGWB and quantify the power spectrum. During the inflationary epoch, the wavelength of the GWs exceeds the Hubble horizon scale due to the exponential expansion of the universe and the amplitude of GWs becomes frozen. After the end of inflation, the universe enters the decelerated expansion phase and wavelengths of GWs soon become shorter than the Hubble scale. When the Hubble horizon scale becomes comparable or smaller than the characteristic size of the extra-dimension, the high-energy effects may significantly affect the evolution of GWs. There are two main high-energy effects : i) peculiar cosmological expansion due to the high-energy correction of the Friedmann equation, which enhances the spectrum in the high frequency region and ii) excitation of Kaluza-Klein modes (KK-modes) freely escaping from the brane to the bulk spacetime, which may suppress the amplitude of the GWs on the brane. While the former effect is simply estimated from the expansion rate of the universe, the amount of the latter effect requires the knowledge of the wave propagation in the bulk. Focusing on the Randall-Sundrum single-brane model, in which the three-dimensional hypersurface (brane) is embedded in the five-dimensional anti-de Sitter (AdS$_5$) spacetime\\cite{RS2}, many authors have tried to estimate these effects in an analytic way. However, these analyses were restricted to the idealistic situations \\cite{Gorbunov,Kobayashi,Kobayashi2004} or low-energy cases in the Friedmann universe \\cite{Battye,Easther}, since the analytic study of wave equation is generally intractable due to the complicated form of equation as well as boundary condition. Thus, in this paper, we numerically solve the wave equation and try to estimate the observed IGWB spectrum.  Concerning numerical studies, several authors have used different numerical techniques and coordinate systems to solve the wave equation of GWs \\cite{Hiramatsu1,Hiramatsu2,Ichiki,Kobayashi2005A,Kobayashi2005B,Sanjeev2006}. In our previous studies \\cite{Hiramatsu1, Hiramatsu2}, numerical simulations were carried out in the two types of coordinate systems. One is the Gaussian normal coordinate system in which we found the suppression of the amplitude of the IGWB on the brane. Unfortunately, the coordinate singularity appears in the bulk and this restricts our analyses to a relatively low-energy scale \\cite{Hiramatsu1}. On the other hand, another coordinate system we used is the Poincar\\'e coordinates in which we observed that the two high-energy effects compensate each other and the spectrum became same one as predicted in the four-dimensional theory \\cite{Hiramatsu2}.  In this paper, we first show that the spectrum of the IGWB estimated in the paper \\cite{Hiramatsu2} is robust against several numerical artifacts, such as the dependence of the regulator brane and initial time. Next, we study the dependence of the spectra on the equation of state (EOS) of the universe to clarify how the two high-energy effects change with the expansion rate of the universe.  This paper consists of seven sections. In Sec. 2, we discuss how the spectrum of the IGWB is affected by the presence of the extra-dimensions. In Sec. 3, we introduce the Poincar\\'e coordinate system and derive the evolution equation of GWs. After briefly discussing the details of the numerical simulations  and initial conditions in Sec. 4, we check our numerical scheme by solving a simple case in Sec. 5. In Sec. 6, we present the numerical results in the case that the IGWB re-enters the Hubble horizon during the radiation dominated (RD) universe. Also in Sec. 6, the results in cases with other equation of state (EOS) are shown. Finally, Sec. 7 is devoted to the summary and conclusions.  In this paper, we use units in which $c=\\hbar=1$. $M_\\text{pl}$ represents the four-dimensional Planck mass/energy.  %===========================================================================% %===========================================================================% ",
        "conclusions": "\\label{sec:conclusion} %===========================================================================% %===========================================================================%  We have investigated the power spectrum of the IGWB in the five-dimensional cosmology based on the Randall-Sundrum model. In the braneworld scenario, the two high-energy effects affect the shape of the spectrum above the critical frequency $f_{\\rm crit}$ defined in (\\ref{eq:critical_frequency}). One is the non-standard cosmological expansion on the brane caused by the high-energy correction of the Friedmann equation. The analytical estimate taking account of this effect reveals that the effect makes the spectrum steeply blue [see Eq. (\\ref{eq:O_propto_f_specific_5D})]. By contrast, another important effect is the excitations of KK-modes which escapes from our brane into the five-dimensional bulk, leading to the suppression of the spectrum. In order to quantify these two effects, we solved the wave equation of each Fourier mode of GWs numerically for various EOS parameters $w$.  The systematic survey of numerical simulations with various parameter sets reveals that there may exist the universal scaling law for the KK-mode excitation in the high-energy regime [Eq.(\\ref{eq:universal_relation1})]: %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % \\begin{equation*} \\left|\\frac{h_{\\rm 5D}}{h_{\\rm ref}}\\right| \\propto (1+H_*^2\\ell^2)^{-1/4}. \\end{equation*} % %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% Using the universal scaling law, we constructed the power spectrum of the IGWB in the cases with $w=0$ (MD universe), $w=1/3$ (RD universe) and $w=1$ (stiff matter dominant universe). From the results (\\ref{eq:Omega_w_4D}) and (\\ref{eq:IGWB_spectra}), the frequency dependence of the spectrum in the high-frequency region $f>f_{\\rm crit}$ becomes %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% % \\begin{equation} \\Omega_{\\rm GW} \\propto \\begin{cases} f^{-2}\\;({\\rm 4D})\\;,\\;f^{-1/2}\\;({\\rm 5D}) &{\\rm for}\\;\\; w=0,\\\\ f^{0},\\;\\hspace{3em} \\;f^{0}                &{\\rm for}\\;\\; w=1/3,\\\\ f^{1},\\;\\hspace{3em} \\;f^{2/5}              &{\\rm for}\\;\\; w=1. \\end{cases} \\end{equation} % %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% Particularly, in the RD case, the accidental cancellation of the two high-energy effects occurs, which yields the same spectrum as one predicted in the four-dimensional (4D) theory. This scaling law might be understood in a context of moving mirror problems. The discussion about the analytic derivation of the scaling law is work in progress.  Finally, we briefly comment on the other numerical works using the different schemes. Recently, Seahra solved numerically the wave equation in the null coordinate system based on the Poincar\\'e coordinates using a sophisticated numerical scheme \\cite{Sanjeev2006}. He observed the agreement between his numerical results and our results for the same choice of the EOS parameters with the same initial conditions as ours.\u00a1\u00a1In addition, it is observed that, even if a constant initial condition on the initial null hypersurface is chosen, the amplitude of GWs on the brane is almost identical with our results. Moreover, the numerical calculation based on the quantum theory has been performed by Kobayashi and Tanaka \\cite{Kobayashi2005B}. They reported the same spectrum in the RD case as ours, even if KK modes are taken into account in the initial de Sitter phase. On the other hand, Ichiki and Nakamura have obtained a tilted spectrum $\\Omega_{\\rm GW}\\propto f^{-0.46}$ \\cite{Ichiki}. While their early results have included errors associated with numerical accuracy, the new calculation using the revised code did not converge to the flat spectrum either. Currently, we do not know the reason why the result by Ichiki and Nakamura is different from ours. In order to understand these numerical results well, the analytical study of the scaling relation (\\ref{eq:universal_relation1}) is essential, which is definitely our next task.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%"
    },
    "0601/astro-ph0601447_arXiv.txt": {
        "abstract": "The superb phase resolution and quality of the OGLE data on LMC and SMC Cepheids, together with existing data on Galactic Cepheids, are combined to study the period-colour (PC) and amplitude-colour (AC) relations as a function of pulsation phase. Our results confirm earlier work that the LMC PC relation (at mean light) is more consistent with two lines of differing slopes, separated at a period of 10 days. However, our multi-phase PC relations reveal much new structure which can potentially increase our understanding of Cepheid variables. These multi-phase PC relations provide insight into why the Galactic PC relation is linear but the LMC PC relation is non-linear. This is because the LMC PC relation is shallower for short ($\\log P < 1$) and steeper for long ($\\log P > 1$) period Cepheids than the corresponding Galactic PC relation. Both of the short and long period Cepheids in all three galaxies exhibit the steepest and shallowest slopes at phases around 0.75-0.85, respectively. A consequence is that the PC relation at phase $\\sim0.8$ is highly non-linear. Further, the Galactic and LMC Cepheids with $\\log P > 1$ display a flat slope in the PC plane at phases close to the maximum light. When the LMC period-luminosity (PL) relation is studied as a function of phase, we confirm that it changes with the PC relation. The LMC PL relation in $V$- and $I$-band near the phase of 0.8 provides compelling evidence that this relation is also consistent with two lines of differing slopes joined at a period close to 10 days. ",
        "introduction": "In \\citet[][hereafter Paper I]{kan04}, we investigated the behavior of period-colour (PC) and also amplitude-colour (AC) relations at the $V$-band maximum, mean and minimum light, as the AC relations are closely related to the PC relations via the following equation, derived in \\citet[][hereafter SKM]{sim93} and in Paper I:  \\begin{eqnarray} \\log T_{max} - \\log T_{min} = \\frac{1}{10} (V_{min} - V_{max}). \\end{eqnarray}  \\ni This equation implies that if the PC relation at the maximum (or minimum) light is flat (i.e. the colour is independent of period), then there will be an AC relation at the minimum (or maximum) light \\citep[see SKM, Paper I and][hereafter Paper II]{kan04a}. Thus an examination of AC relations can be used to confirm changes to the PC relation. The major assumption used in equation (1) is that temperature, not radius, fluctuations are the dominant cause of light variations, at least at the optical wavelengths \\citep{cox80}. In this paper, in addition to the maximum, mean and minimum light, we extend the study of PC and AC relations at all other pulsation phases. Furthermore, the extension to multi-phase relations will also expand our understanding of the pulsational behavior of classical Cepheids. The results of our study justify our methods and reveal new properties of Cepheids in the Galaxy and Magellanic Clouds that can be used to constrain the theoretical Cepheid pulsation/evolution models in future studies.  Based on the extensive data and large number of Cepheids recently discovered in the Large Magellanic Cloud (LMC), there is strong evidence that the mean light PC relation for the LMC Cepheids is non-linear: there are two PC relations for the short ($\\log P<1$) and long ($\\log P > 1$) period Cepheids \\citep{tam02,tam02a,kan04,san04,kan06,nge05}\\footnote{In the rest of this paper, this is what we mean by non-linearity, and the long and short period Cepheids refer to those with $\\log P>1.0$ and $\\log P<1.0$, respectively.}, respectively. In contrast, the Galactic or the Small Magellanic Cloud (SMC) Cepheids do not show any non-linearity of the mean light PC relation \\citep{tam03,kan04}. Therefore, our motivation is to construct and study the PC relations for Galactic, LMC and SMC Cepheids at all phases between zero and one, because the PC relation at mean light is the average of the PC relation at all phases (see Paper I). Thus it is important to understand the properties of the PC relation at other phases. A direct consequence of the non-linear LMC PC relation at mean light is that the LMC period-luminosity (PL) relation at mean light is also non-linear\\footnote{See, for example, \\citet{mad91} for the basic physics of the Cepheid PL and PC relations.} \\citep{tam02,kan04,san04,nge05}. Although the LMC PL relation is extensively used in distance scale studies \\citep[see, e.g.,][]{fre01,sah01}, \\citet{nge05w} have discussed the implications of the non-linear mean light LMC PL and PC relations in distance scale applications.  In Section 2 of this paper, we present the data and the numerical methods used in this study. The multi-phase PC relations are studied in Section 3, and the detailed comparison between the Galactic and LMC PC relations is given in Section 3.1. The multi-phase AC relations are briefly presented in Section 4. Since it is also of great interest to investigate the connections between the PL and PC relations for LMC Cepheids, the multi-phase LMC PL relations are presented in Section 5. Finally, the main conclusion of this paper is summarized in Section 6. ",
        "conclusions": "In this paper we have used the excellent phase resolution of OGLE data, combined with existing Galactic data to present multi-phase PC and AC relations for Galactic, LMC and SMC Cepheids. The main conclusions from this study are summarized below.  \\begin{enumerate} \\item By examining the slopes of the PC relation as a function of phase, our results confirm that the Galactic PC(mean) relation is linear and the LMC PC(mean) relation is non-linear in the sense that the LMC PC relation is more consistent with two linear relations, one for short and long period Cepheids respectively.  \\item Both short and long period Cepheids exhibit interesting new features on these multi-phase plots. The slopes of the PC relation for long period Cepheids reach maximum at a phase around $0.7-0.8$, close to the phase of minimum light. In contrast the slope of the PC relation for short period Cepheids displays a minimum (close to zero) around $\\Phi \\sim 0.8$. These are the new observational features that have not be reported in the literature, and can be used to further constrain stellar pulsation/evolution models of Cepheids.  \\item Combining these behaviors of the slopes at $\\Phi \\sim0.8$ for the long and short period Cepheids, the PC relation becomes strongly non-linear at this phase. Further, the Galactic and LMC long period Cepheids display a flat PC relation at phases close to the maximum light (at $\\Phi=0.0$). This also cause both of the Galactic and LMC PC relation to be non-linear at maximum light. There are some indications that the extent of this flatness may be metallicity dependent since the SMC has a non-zero PC slope around maximum light.  \\item The reason why previous workers found a non-linear and linear LMC and Galactic PC relation is because short and long period LMC Cepheids have a shallower and steeper slope respectively than their Galactic counterparts in most of the phases. In addition, the LMC PC relations are actually non-linear at most pulsational phases. The Galactic PC relations are linear at most pulsational phases.  \\item The multi-phase SMC PC relations show some similarities and differences to the Galactic and LMC counterparts. In particular, the non-linear PC relations around $\\Phi \\sim0.8$ is significant for all three galaxies, and this suggests the non-linearity of the PC relation around this phase is a universal feature. Perhaps it may has to do with the interaction between the hydrogen ionization front (HIF) and the photosphere at outer envelope of Cepheid variables, or even a shock+HIF+photosphere interaction. However a detailed investigation is beyond the scope of this paper.  \\item The shapes of the multi-phase AC relations are similar for both long and short period Cepheids, except the curves for short period Cepheids are shifted downward. The multi-phase AC relations also support the validity of equation (1).  \\item The multi-phase LMC PL relations also trace the corresponding multi-phase PC relations (except at maximum light). Evidence of non-linear PL relations at $\\Phi \\sim0.8$ is also compelling. It is well known that near resonance, the shape of Cepheid light curves changes dramatically and it could be that near this period, the secondary maximum associated with the Hertzsprung progression may become the primary maximum. Perhaps this may be responsible for the sharp change seen at $\\Phi \\sim0.8$ in the LMC PC/PL relations. However, even if this is the case, it does not negate the observational fact that there are very compelling changes in the LMC PC/PL relations as a function of phase.  \\item Since the mean light PC and PL relation is an average over the relations at every phase, the behavior around $\\Phi \\sim0.8$ has some impact on the mean light relation which makes the LMC PC/PL relation non-linear but the Galactic and SMC PC/PL relations linear at mean light. Our contention is that this effect is reduced in the Galaxy and SMC because of the combined effect of the changes of the HIF-photosphere interaction and the amplitude of the Cepheids.  \\end{enumerate}   \\begin{figure} \\hbox{\\hspace{0.1cm}\\epsfxsize=7.5cm \\epsfbox{MF920_fig13.eps}} \\caption{Comparison of the $F$-values for the LMC PL and PC relation as a function of phase. The dashed lines represent the 95\\% confidence level (or $p[F]=0.05$), corresponding to $F\\sim3.0$. Hence, the $F$-values above the dashed line indicate that the PL/PC relation is non-linear, while the $F$-values below this line suggest the PL/PC relation is linear.} \\label{figcom_fpl} \\end{figure}"
    },
    "0601/astro-ph0601392_arXiv.txt": {
        "abstract": "If supermassive black holes in centres of galaxies form by merging of black-hole remnants of massive Population III stars, then there should be a few black holes of mass one or two orders of magnitude smaller than that of the central ones, orbiting around the centre of a typical galaxy. These black holes constitute a weak perturbation in the gravitational potential, which can generate wave phenomena in gas within a disc close to the centre of the galaxy. Here we show that a single orbiting black hole generates a three-arm spiral pattern in the central gaseous disc. The density excess in the spiral arms in the disc reaches values of 3-12\\% when the orbiting black hole is about ten times less massive than the central black hole. Therefore the observed density pattern in gas can be used as a signature in detecting the most massive orbiting black holes. ",
        "introduction": "Numerical simulations of structure formation indicate that baryonic cores within dark matter haloes collapse at redshifts $z\\approx20-30$ to sufficiently high densities to be available for primordial star formation (\\eg Abel et al 1998). Because of the lack of metals that could facilitate cooling, the collapse occurs without much fragmentation and it gives rise to the Population III stars, much more massive than stars forming today (Abel, Bryan \\& Norman 2000). Stars as massive as $10^3$ \\solm could form, and they could evolve into massive black holes (MBHs) with little mass loss (Madau \\& Rees 2001). A Population III star with a mass greater than 260 \\solm will not experience a supernova explosion, because the gravity of the star is too strong, and instead it will simply collapse to form a black hole. Islam, Taylor \\& Silk (2003, 2004a) followed in detail hierarchical merging of such a population of primodial black holes, and they found that about $10^3$ MBHs would be present in the galactic halo today. In total, mass comparable or greater than that of the central supermassive black hole (SMBH) should still reside in MBHs orbiting in galactic halos. Most of these MBHs would have masses near the initial seed mass of the remnant Population III star, but there would be a few as massive as $10^6-10^7$ \\solm.  In this paper we investigate whether the most massive MBHs orbiting in galactic halos can be revealed by structures related to waves that they generate in gaseous discs in centres of galaxies. At the nuclear scales, gas dynamics becomes decoupled from dynamics of the stars, and one can study the response the of the former to any imposed gravitational potential. A linear theory applicable here was developed by Goldreich \\& Tremaine (1978, 1979), who showed that density waves can propagate within the gas as a response to a rigidly rotating potential. The propagation of waves depends crucially on the resonances created by the rotating potential, since waves are generated at these resonances. Waves propagating in gas generate spiral structure there, and the linear theory was first invoked by Lindblad \\& J\\\"{o}rs\\\"{a}ter (1981) to explain the spiral structure observed in the centre of a barred galaxy. Maciejewski (2004 a,b) used the linear theory and hydrodynamical modelling to analyse nuclear spirals generated by a wide range of galactic potentials, including ones that have a SMBH at the centre and a bisymmetric bar. These latter papers suggest that asymmetries in the galactic potential too weak to be detected observationally might be sufficient to generate nuclear spirals that should be easily observed.  A single MBH orbiting in the galactic halo gives rise to an asymmetric perturbation in the potential rather than to a bisymmetric one, caused by a bar, that has been studied previously. In this paper we investigate in detail gas dynamics generated by this potential using hydrodynamical modelling. There are similarities here to the interaction between a protoplanet and the protostellar disc (e.g. Lin \\& Papaloizou 1993), although there are also significant differences which we point out.  The structure of the paper is as follows. Section 2 presents the numerical code, the gravitational potential and the parameters used. All of the models are summarized in Table 1, and in Section 3 representative models are analyzed further. An interpretation of the results is given in Section 4. Implications of the models for the detection of MBHs orbiting in galaxies, and other detection methods are discussed in Section 5. We summarize our conclusions in Section 6. ",
        "conclusions": "This paper was based on the innovative idea that orbiting MBHs could produce a signature pattern in the gas that may be used to identify them. The main result of this study is that a three-armed spiral is generated in gas when a single MBH orbits around the galactic centre. This structure is the consequence of resonance-excited waves in gas. Only MBHs at the very end of the upper range of masses permitted by dynamical constraints and evolutionary scenarios, \\ie with masses nearing 10$^7$ \\solm, are capable to generate structure in gas with density contrast high enough to be observed with current techniques. However, given that there are not many techniques that could assure detection of high-mass orbiting MBHs, the signature proposed here might be well worth pursuing."
    },
    "0601/astro-ph0601671_arXiv.txt": {
        "abstract": "Using moderate-resolution Keck spectra, we have examined the velocity profiles of 15 members of cluster Cl0024+1654 at $z=0.4$.  WFPC2 images of the cluster members have been used to determine structural parameters, including disk sizes, orientations, and inclinations.  We compare two methods of optical rotation curve analysis for kinematic measurements.  Both methods take seeing, slit size and orientation, and instrumental effects into account and yield similar rotation velocity measurements.  Four of the galaxies in our sample exhibit unusual kinematic signatures, such as non-circular motions.  Our key result is that the Cl0024 galaxies are marginally underluminous ($0.50 \\pm 0.23$ mag), given their rotation velocities, as compared to the local Tully-Fisher relation.  In this analysis, we assume no slope evolution, and take into account systematic differences between local and distant velocity and luminosity measurements.  Our result is particularly striking considering the Cl0024 members have very strong emission lines, and local galaxies with similar H$\\alpha$ equivalent widths tend to be overluminous on the Tully-Fisher relation.  Cl0024 Tully-Fisher residuals appear to be correlated most strongly with galaxy rotation velocities, indicating a possible change in the slope of the Tully-Fisher relation.  However, we caution that this result may be strongly affected by magnitude selection and by the original slope assumed for the analysis.  Cl0024 residuals also depend weakly on color, emission line strength and extent, and photometric asymmetry.  In a comparison of stellar and gas motions in two Cl0024 members, we find no evidence for counter-rotating stars and gas, an expected signature of mergers. ",
        "introduction": "The Tully-Fisher relation (TFR, Tully \\& Fisher 1977) between disk galaxy luminosities and rotational velocities is an essential diagnostic of galaxy evolution.  By comparing galaxy evolution models to TFR observations at $z \\sim 0$ through $z \\sim 1$, it is possible to probe dark matter halo properties and galaxy star formation histories (e.g., Mo \\& Mao 2000; Navarro \\& Steinmetz 2000; Ferreras \\& Silk 2001; Buchalter, Jimenez \\& Kamionkowski 2001a, 2001b; Mo \\& Mao 2004).  Observations of the distant field galaxy TFR have clearly demonstrated luminosity evolution since $z \\sim 1$, but the amount of evolution is contested.  Forbes et al.\\ (1996), Vogt et al.\\ (1996, 1997) and Rigopoulou et al.\\ (2002) find $\\le$1 mag brightening between $z \\sim 0$ and $z \\sim 1$, whereas Rix et al.\\ (1997) find evidence for 1.5 mag brightening at $z \\sim 0.25$.  In agreement with the latter work, Simard \\& Pritchet (1998) and Mall\\'{e}n-Ornelas et al.\\ (1999) find evidence for 1.5--2 mag brightening at $z \\sim 0.25 - 0.6$.  These two studies, however, also hint that there may be a change in TFR slope with redshift.  Recent studies using larger samples of field galaxies ($\\geq$60 compared to sample sizes of $\\la$20 in the studies mentioned above) by Ziegler et al.\\ (2002) and B\\\"{o}hm et al.\\ (2004) have provided further evidence for evolution in TFR slope, possibly reconciling previous findings.  More specifically, their work indicates that the mass-to-light ratios of the least massive galaxies evolve most strongly.  These authors attribute the differing results of previous studies to selection effects and small sample sizes.  Their results are consistent with kinematic studies of blue compact galaxies at $z \\le 1$ (Koo et al.\\ 1995, 1997; Phillips et al.\\ 1997; Guzm\\'{a}n et al.\\ 1996, 1997, 1998).  The latter studies demonstrate that distant compact galaxies are overluminous as compared to the local TFR, in many cases an order of magnitude less massive than expected for their luminosities.  While the work done to date is promising, there are still ways in which distant Tully-Fisher studies can be improved.  For instance, selection effects are thought to be responsible for the differing results of some of the studies mentioned above.  However, while effects such as magnitude incompleteness have been explored in local studies (e.g. Willick 1994, Giovanelli et al.\\ 1997), little work has yet been done toward quantifying and correcting for such effects on distant samples.  Furthermore, the velocity measurement techniques used in local and distant Tully-Fisher studies are often very different and may account for some of the discrepancies between previous results.  Authors of distant Tully-Fisher studies also tend to compare their samples to the relatively small local sample of Pierce \\& Tully (1992) to draw evolutionary conclusions.  While this is a good starting point, larger and better-calibrated local samples now exist for comparison (e.g., Tully \\& Pierce 2000; Kannappan, Fabricant \\& Franx 2002).  We also note that while significant recent progress has been made investigating the evolution of the Tully-Fisher relation in the field, TFR evolution in the cluster environment is only beginning to be explored.  Many physical processes are expected to affect galaxies in this dense environment, including tidal encounters, interactions with the intracluster medium, and group infall.  All of these processes are predicted to affect cluster galaxy disks and gas, and in particular their mass-to-light ratios, thus affecting the TFR.  As a specific example, high-velocity tidal encounters between cluster galaxies and with the gravitational potential of the cluster itself (``harassment\") have been modeled by Moore et al.\\ (1996, 1998, 1999).  These encounters are expected to increase star formation rates while stripping galaxies of a significant fraction of their mass (Fujita 1998), thus decreasing their mass-to-light ratios.  Gnedin (2003) predicts that lower-density galaxies in clusters are more likely to be disrupted and lose a greater percentage of mass in tidal encounters than large spirals.   The mass-to-light ratios of low- and high-mass galaxies in clusters may therefore evolve differently.  In distant galaxy clusters, we have the opportunity to observe such evolutionary processes in action.  Distant cluster Tully-Fisher studies will therefore allow us to test model predictions and quantify evolutionary effects.  Few studies of the Tully-Fisher relation in distant clusters currently exist.  While studies of $z \\sim 0.1$ clusters have been conducted with significant sample sizes ($>$50 galaxies; e.g., Rubin et al.\\ 1999, Dale \\& Uson 2003), kinematics of few $z>0.1$ cluster disks have been measured.  In a pioneering study, Vogt et al.\\ (1993) measured the rotation velocities of two cluster galaxies at $z=0.20$ and $z=0.38$.  Franx (1993) also demonstrated rotation in an ``E+A\" galaxy in Abell 665 at $z=0.18$.  All three objects were found to be $\\leq 1$ mag brighter than local disk galaxies.  More recently, Milvang-Jensen et al.\\ (2003) analyzed the rotation curves of eight members of cluster MS1054-03 at $z=0.83$ and found that they were $\\sim$1.5 mag more luminous than local field galaxies.  Bamford et al.\\ (2005) have expanded this study with a larger sample size of 22 members of three clusters at $0.3 < z < 0.9$.  They have found that the cluster members are somewhat ($\\sim$0.5 mag) more luminous than field galaxies at the same redshifts.  On the other hand, Ziegler et al.\\ (2003; see J\\\"{a}ger et al.\\ 2004 for details on analysis techniques), presented rotation curves of 13 members of three clusters at $0.3 < z < 0.5$ and found no evidence for luminosity evolution of cluster galaxies with respect to the field.  Other than these noteable exceptions, the majority of kinematic studies in distant clusters have focused on the fundamental plane of early-type galaxies.  These works indicate that the amount of luminosity evolution in cluster early-types is consistent with passive evolution out to $z = 0.83$ (van Dokkum \\& Franx 1996; van Dokkum et al.\\ 1998; Kelson et al.\\ 1997, 2000a,b; J$\\o$rgensen et al.\\ 1999) and possibly out to $z = 1.27$ (van Dokkum \\& Stanford 2003).  As described above, environmentally-dependent physical processes are expected to affect cluster disks more strongly than early-types.  The distant cluster TFR may provide clues as to which processes are dominant.  As a first step, we have chosen to study the TFR in rich cluster Cl0024+1654 at $z=0.4$, where evolutionary processes are very likely to be apparent.   In general, morphological, luminosity, and color evolution have already been seen in cluster galaxies since $z \\simeq 0.5$ (e.g., Dressler et al.\\ 1997, J$\\o$rgensen et al.\\ 1999, and Butcher \\& Oemler 1984, respectively).  The redshift of this cluster alone places it at an epoch where we would expect to observe evolutionary signatures.  Cl0024 was one of the two clusters originally studied by Butcher \\& Oemler (1978) and thus has long been known to have a high fraction of blue, presumably star-forming members.  Furthermore, evidence has been found for evolutionary processes such as mergers (Lavery, Pierce, \\& McClure 1992) and recent subcluster infall (Czoske et al.\\ 2002).  The structure of this paper runs as follows: in the next section, we discuss our observations of Cl0024+1654 and our sample selection for this study.  Note that our sample of 15 Cl0024 members is one of the largest samples of rotation curves presented thus far for galaxies within a single $z > 0.2$ cluster.  We detail our basic spectral and image reductions, two methods of rotation curve analysis, and we discuss the differences between local and distant galaxy velocity measurements in \\S 3.  In \\S 4, we present the Cl0024 Tully-Fisher relation, including a comparison to three local samples.  In \\S 5, we discuss evolutionary indications from correlations between Tully-Fisher residuals and other galaxy properties.  We also take a look at the effect of magnitude incompleteness on our analysis.  We present our conclusions in \\S 6.  Throughout this paper we assume $H_{0}=70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{M} = 0.3$ and $\\Omega_{\\Lambda} = 0.7$.  Using this cosmology, one arcsecond corresponds to 5.332 kpc.  We present all photometry in the Vega system. ",
        "conclusions": "In this paper, we have measured the rotation velocities of 15 members of galaxy cluster Cl0024 at $z=0.4$, and we have presented one of the first Tully-Fisher analyses for a distant cluster.  Our key points are the following:  \\begin{enumerate} \\item We have measured Cl0024 member velocities via two different rotation curve analysis methods: GAUSS2D, which closely resembles the method used by Vogt et al.\\ (1996, 1997), and GELFIT2D, which is similar to the method used by Simard \\& Pritchet (1998, 1999).  Both methods yield similar rotation velocity measurements, within the errors. \\item We find that four of the galaxies in our sample have unusual kinematic properties or emission distributions.  One of these galaxies is a likely AGN, another has a knotty ring of star formation.  The remaining two have odd morphologies and are poorly fit by our velocity models, implying possible non-circular motions and/or recent interactions.  Similar fractions ($\\sim25\\%$) of kinematically ``anomalous\" galaxies have been found in the distant field studies of Rix et al.\\ (1997) and Simard \\& Pritchet (1998).  In the local Virgo cluster, Rubin, Waterman, \\& Kenney (1999) find that as many as 50\\% of spirals exhibit unual kinematic distributions. \\item The Cl0024 Tully-Fisher relation is consistent with no luminosity evolution, or perhaps a marginal {\\it underluminosity}, as compared to the local sample of Kannappan, Fabricant, \\& Franx (2002).  In our analysis of the Cl0024 TFR, we have demonstrated the importance of accounting for differences between distant and local galaxy velocity and luminosity measurement techniques.  Furthermore, the choice of local comparison sample influences the interpretation of Tully-Fisher evolution.  Combined, these two effects change our Cl0024 TFR results by $\\sim$1 mag. \\item We have probed the dependence of Cl0024 members' Tully-Fisher residuals upon other galaxy properties, in the hopes of uncovering evidence of the evolutionary processes that may be affecting the galaxies.  The strongest relationship we find is between TF residuals and rotation velocities in the cluster, suggesting Tully-Fisher slope evolution.  We caution that magnitude selection and the choice of Tully-Fisher slope used in the residual analysis may strongly affect this result, however. \\item Evidence for galaxy interactions as an evolutionary driver in Cl0024 is unclear from the weak correlations we find between TF residuals and photometric residuals and asymmetries.  No apparent relationship exists between TF residuals and emission line residuals and asymmetries, but this may be due to the relatively low resolution of our kinematic data.  Furthermore, this result may be affected by small number statistics. \\item We find no evidence for counter-rotation between stars (as measured from absorption) and gas (as measured from emission) in two Cl0024 members for which we have high enough S/N data to make such measurements.  One of these galaxies exhibits characteristics of an AGN; the other appears to be morphologically disturbed. \\end{enumerate}  This study is intended as a pilot investigation of the Tully-Fisher relation in distant clusters.  With our small sample size and focus on a single cluster, we cannot make sweeping statements about the TFR and TF residuals as indicators of evolutionary processes in clusters.  However, our results are intended as a baseline for comparison to future studies made with larger samples of distant cluster galaxies.  In the future, observations with integral field units (e.g., Pracy et al.\\ 2005) will likely be essential to understanding the interplay between galaxy internal kinematics and evolutionary drivers in both the cluster and field environments.  Such observations will provide two-dimensional measurements linking star formation rates, metallicities, shock diagnostics, and kinematics, allowing observers to determine direct links between these diagnostics and evolutionary predictions."
    },
    "0601/astro-ph0601501_arXiv.txt": {
        "abstract": "Nine 20$\\,M_\\odot$ models  were computed with metallicities ranging from solar, through $Z=10^{-5}$ ([Fe/H]$\\sim$-3.1) down to $Z=10^{-8}$ ([Fe/H]$\\sim$-6.1) and with initial rotational velocities between 0 and 600\\,km\\,s$^{-1}$ to study the impact of initial metallicity and rotational velocity \\cite{H05}. The very large amounts of $^{14}$N observed ($\\sim$0.03$\\,M_\\odot$) are only produced at $Z=10^{-8}$ (PopII 1/2). The strong dependence of the $^{14}$N yields on rotation and other parameters like the initial mass and metallicity may explain the large scatter in the observations of $^{14}$N abundance. The metallicity trends are best reproduced by the models with $\\upsilon_{ini}/\\upsilon_c \\sim 0.75$, which is slightly above the mean observed value for OB solar metallicity stars. Indeed, in the model with $\\upsilon_{ini}$=600\\,km\\,s$^{-1}$ at $Z=10^{-8}$, the $^{16}$O yield is reduced due to strong mixing. This allows in particular to reproduce the upturn for C/O and a slightly decreasing [C/Fe], which are observed below [Fe/H]$\\sim$ -3. ",
        "introduction": "Precise measurements of abundances of low metallicity stars have recently been obtained by \\cite{FS5}, \\cite{FS6} and \\cite{IR04}. These provide new constraints for the stellar evolution models (see \\cite{CMB05}, \\cite{Fr04} and \\cite{Pr04}). The most striking constraint is the need for primary $^{14}$N production in very low metallicity massive stars. Rotation helps producing large amounts of primary $^{14}$N through mixing of newly synthesised carbon and oxygen during helium burning (see \\cite{MEM05} and contributions by these authors in this volume). Other constraints are an upturn of the C/O ratio with a [C/Fe] about constant or slightly decreasing (with increasing metallicity) at very low metallicities, which requires an increase (with increasing metallicity) of oxygen yields below [Fe/H]$\\sim$ -3. This seems hard to reproduce with rotating models were mixing usually increases the size of the helium burning core and therefore increases the oxygen yields. ",
        "conclusions": ""
    },
    "0601/hep-ph0601108_arXiv.txt": {
        "abstract": "We present fully non-linear simulations of a self-interacting scalar field in the early universe undergoing tachyonic preheating. We find that density perturbations on sub-horizon scales which are amplified by tachyonic instability maintain long range correlations even during the succeeding parametric resonance, in contrast to the standard models of preheating dominated by parametric resonance. As a result the final spectrum exhibits memory and is not universal in shape. We find that throughout the subsequent era of parametric resonance the equation of state of the universe is almost dust-like, hence the Jeans wavelength is much smaller than the horizon scale. If our 2D simulations are accurate reflections of the situation in 3D, then there are wide regions of parameter space ruled out by over-production of black holes. It is likely however that realistic parameter values, consistent with COBE/WMAP normalisation, are safetly outside this black hole over-production region. ",
        "introduction": "Over the past decade, the theory of reheating after inflation has been developed and it was recognized that reheating can be accompanied by an era of preheating, where fluctuations of scalar fields are greatly amplified either by parametric resonance \\cite{Dolgov:1990,Traschen:1990sw,Fugisaki:1996,Kofman:1994rk,Kofman:1997yn} or by tachyonic instability \\cite{Green:1997,Felder:2000hj,Felder:2001kt}. In preheating, energy transfer from the inflaton field to another field occurs in bursts where non-perturbative effects are important for the evolution of scalar fields.  A common feature among the various preheating models is that preheating leaves an imprint on the subsequent evolution of the universe, by amplifying the density perturbation especially on sub-horizon scales at that time by many orders of magnitude. If the resultant density perturbations at horizon crossing becomes of order unity, black holes can be formed by gravitational collapse \\cite{Carr:1974nx}. Since the energy density of produced BHs increases proportional to the scale factor relative to the total energy density of the universe in the subsequent radiation dominated era, BHs will dominate density of the universe at late time even if a tiny fraction of energy is converted into BHs at preheating. There are strong constraints on the abundance of BHs for a wide range of mass scales \\cite{kohri} and they can be a powerful means for constraining inflationary models \\cite{Yokoyama:1999xi}.   In the past, it was pointed out by several authors that BHs can be over-produced during preheating for two-field models \\cite{agreen, bassett,finelli}. In these models, parametric resonance causes the rapid amplification of field fluctuations. Then, non-perturbative effects such as rescattering may be crucial for calculating the abundance of produced BHs. Indeed, it was shown via lattice simulations that BHs are unlikely to be over-produced during preheating \\cite{Suyama:2004mz}. Accurate simulations are required since a small error in the estimate of backreaction time which ends the resonance causes a large error in the amplitude of density perturbation at the end of preheating because field fluctuations grow exponentially during preheating.  In this paper, we use lattice simulations to examine the evolution of density perturbations during tachyonic preheating and the succeeding parametric resonance, mainly focusing on the production of black holes. To evaluate the amplitude of density perturbations taking the non-linear interactions into account, we performed two and three-dimensional lattice simulations using the C++ code LATTICEEASY developed by Felder and Tkachev \\cite{latticeeasy}.  The equations of motion of the minimally coupled scalar field that we solve are \\begin{eqnarray} && {\\ddot \\phi}+3 H {\\dot \\phi} -a^{-2}\\triangle {\\phi} + V'(\\phi)=0, \\\\ && H^2 = \\frac{8\\pi}{3 M^2_{pl}} \\langle \\rho \\rangle, \\end{eqnarray} where $ \\rho $,  $ a $ and $H$ are, respectively, the total energy density, the scale factor and the Hubble parameter, and $ \\langle \\cdots \\rangle $ denotes spatial average. To estimate the variance of density perturbations at the horizon scale ($k=aH$), we simply use the power spectrum ${\\cal P}_{\\delta} (aH)$. Here ${\\cal P}_{\\delta} (k)$ is defined by \\begin{equation} \\langle {\\delta}_k {\\delta}_{-k'} \\rangle = \\frac{2 {\\pi}^2}{k^3} {\\cal P}_{\\delta} (k) \\delta ( \\vec{k}-\\vec{k'}), \\label{pbh5} \\end{equation} and $\\delta_k$ is the Fourier component of the density contrast $ \\delta \\equiv \\delta \\rho/\\langle \\rho \\rangle $. ",
        "conclusions": "In summary, we studied non-linear evolution of density perturbations on sub-horizon scales by two and three dimensional fully nonlinear lattice simulations. We adopted the Coleman-Weinberg potential as a model of tachyonic instability, since this model is expected to capture typical features in more extended preheating models.  We found that the density perturbations below the horizon scale are selectively amplified by tachyonic instability and thereafter maintain their amplitude, while fluctuations of the inflaton field decrease in time if the following parametric resonance dominates the perturbations. If the non-linear interactions become significant during the tachyonic phase, parametric resonance is less efficient and long range correlations of inflaton fluctuations do not decay.  We also found that there is no enhancement of perturbations due to the non-linear interactions and the terminal value of the spectrum ${\\cal P}_{\\delta}$ on subhorizon scales is bounded from above by the linear estimate $\\delta N^2\\sim \\lambda$. Hence if we extrapolate this result to the realistic $\\lambda \\sim 10^{-10}$ determined by temperature fluctuations of the cosmic microwave background, we see that BH over-production is unlikely to occur for realistic $\\lambda$ with $\\delta N \\sim 10^{-5}$.    We also found that the resulting power spectrum of density perturbations strongly depends on the initial conditions when preheating is dominated by the tachyonic instability. The latter point is in contrast to the cases dominated by parametric resonance, in which the spectrum tends to obey a universal power law ${\\cal P}_\\delta \\propto k^D $ ($ D $ is the spatial dimension) independently of the initial shape of the spectrum.   Finally we briefly comment on the role of metric perturbations which are not taken into account in our simulations. It is necessary to include metric perturbations for the evolution of long wavelength perturbations. In linear perturbation theory, the evolution equations for inflaton fluctuations on flat slicing is given by \\cite{Sasaki:1995aw} \\begin{equation} {\\ddot {\\delta \\phi}}+3H {\\dot {\\delta \\phi}} +\\frac{k^2}{a^2}{\\delta \\phi}+V'' \\delta \\phi=a^{-3} {\\left( \\frac{a^3}{H} {\\dot{\\phi}}^2 \\right)}^{\\cdot} \\delta \\phi. \\label{dis1} \\end{equation} The right hand side of eq.~(\\ref{dis1}) represents the contributions from metric perturbations. We solved eq.~(\\ref{dis1}) numerically and compared the resulting power spectrum with the one without metric perturbations. We found little difference between the two power spectra, which shows that the neglecting the metric perturbation is a good approximation within linear perturbation theory.  There is a possibility that non-linearity may enhance the perturbations when metric perturbations are taken into account but this possibility is beyond the scope of the current work.   T.S. thanks Misao Sasaki and Takashi Nakamura for useful comments. This work is supported in part by Grant-in-Aid for Scientific Research, Nos. 14047212 and 16740141, and by that for the 21st Century COE \"Center for Diversity and Universality in Physics\" at Kyoto university, both from the Ministry of Education, Culture, Sports, Science and Technology of Japan."
    },
    "0601/nucl-th0601086_arXiv.txt": {
        "abstract": "This review describes the properties of hadronic phases of dense matter in compact stars. The theory is developed within the  method of real-time Green's functions and is applied to study of baryonic matter at and above the saturation density. The non-relativistic and covariant theories based on continuum Green's functions and  the $T$-matrix and related approximations to the self-energies are reviewed. The effects of symmetry energy, onset of hyperons and meson condensation on the properties of stellar configurations are demonstrated on specific examples. Neutrino  interactions with baryonic matter are introduced within a  kinetic theory. We concentrate on the classification, analysis and first principle derivation of neutrino radiation processes from unpaired and superfluid hadronic phases. We then demonstrate how neutrino radiation rates from various microscopic processes affect the macroscopic cooling  of neutron stars and how the observed  $X$-ray fluxes from pulsars constrain the properties of dense hadronic matter. ",
        "introduction": "\\label{sec:INTRO}  Neutron stars are born in a gravitational collapse of luminous stars whose core mass exceeds the Chandrasekhar limit for a self-gravitating body supported by degeneracy pressure of electron gas~\\cite{CHANDRA}. Being the densest observable bodies in our universe they open a window on the physics of matter under extreme conditions of high densities, pressures and strong electromagnetic and gravitational fields. Most of the known pulsars are isolated objects which emit radio-waves at frequencies $10^8-10^{10}$ Hz, which are pulsed at the rotation frequency of the star. Young objects, like the pulsar in the Crab nebula, are also observed through $X$-rays that are emitted from their surface as the star radiates away its thermal energy. Relativistic magnetospheres of young pulsars emit detectable non-thermal optical and gamma radiation; they could be sources of high-energy elementary particles. Neutron stars in the binaries are powered by the energy of matter accreted from companion star.  The radio observations of pulsars stretch back to 1967 when the first pulsar was discovered. Since then, the observational pulsar astronomy has been extremely important to the fundamental physics and astrophysics, which is evidenced by two Nobel awards, one for the discovery of pulsars (Hewish 1974)~\\cite{HEWISH}, the other for the discovery of the first neutron star -- neutron star binary pulsar (Hulse and Taylor, 1993)~\\cite{HULSE_TAYLOR} whose orbital decay confirmed the gravitational radiation in full agreement with Einstein's General Theory of Relativity. The measurements of neutron star (NS) masses in binaries provide one of the most stringent constraint on properties of the superdense matter.  The timing observations of the millisecond pulsars set an upper limit on the angular momentum that can be accommodated by a stable NS  - a limit that can potentially constrains the  properties of dense matter. Noise and rotational anomalies  that are superimposed on the otherwise highly stable rotation of these objects provide a clue to the superfluid interiors of these stars.  The orbiting $X$-ray satellites allow astronomers to explore the properties of the NS's surface, in particular its composition through the spectrum of thermal radiation and the transients like thermonuclear burning of accreted material on the surface of an accreting NS. The currently operating Chandra $X$-ray satellite,  Newton $X$-ray Multi Mirror Mission (XMM),  Rossi $X$-ray timing explorer (RXTE) continue to provide new insight into compact objects, their environment and thermal histories. Confronting the theoretical models of NS cooling with the $X$-ray observations constrains the properties of dense hadronic matter, its elementary particle content and its condensed matter aspects such as superfluidity and superconductivity.  The currently operating gravitational wave observatories VIRGO, LIGO, GEO and TAMA are expected to detect gravity waves from various compact objects, the NS-NS binaries being one of the most important targets of their search. Higher sensitivity will be achieved by the future space-based observatory LISA, currently under construction. The global oscillations of isolated NS and NS in binaries with compact objects (in particular with NSs and black holes) are believed to be important sources of gravity waves, which will have the potential to shed light on the internal structure and composition of a NS.  The theory of neutron stars has its roots in the 1930s when it was realized that self-gravitating matter can support itself against gravitational collapse by the degeneracy pressure of fermions (electrons in the case of white dwarfs and neutrons and heavier baryons in the case of neutron stars). The underlying mechanism is the Pauli exclusion principle. Thus, unlike the ordinary stars, which are stabilized by their thermal pressure, neutron stars owe their very existence to the quantum nature of matter. The theory of neutron stars has been rapidly developing during the past four decades since the discovery of pulsars. The progress in this field was driven by different factors: the studies of elementary particles and their strong and weak interactions at terrestrial accelerators and the parallel developments in the fundamental theory of matter deeply affected our current understanding of neutron stars. The concepts of condensed matter, such as superfluidity and superconductivity, play a fundamental role in the dynamical manifestation of pulsars, their cooling and transport properties. Another factor is the steady increase of computational capabilities.  This review gives a survey of the many-body theory of dense matter in NS. It concentrates on the uniform phases and hadronic degrees of freedom which cover the density range $\\rho_0 \\le \\rho \\le 3\\rho_0$, where $\\rho_0$ is the nuclear saturation density. The theory is developed within the framework of  continuum Green's function technique at finite temperature and density~\\cite{SCHWINGER61,KADANOFF62,KELDYSH64,ABRIKOSOV75}. Such an approach allows for a certain degree of coherence of the presentation. Subsection~\\ref{other_CHAP2} at the end of Chapter~2 gives a brief summary of  methods and results omitted in the discussion.  It is virtually impossible to cover all the aspects of the NS theory even when restricting to a certain domain of densities and degrees of freedom.  The topics included in this review are not surprisingly aligned with the research interests of the author. It should be noted that there are very good textbooks~\\cite{SHAPIRO_TEUKOLSKY,ZELDOVICH_NOVIKOV}, monographs~\\cite{GLEN_BOOK,WEBER_BOOK,SAAKYAN} and recent lecture notes~\\cite{BGS} that cover the basics of compact star theory much more completely than it is done in this review. Furthermore there are several recent comprehensive reviews that cover more specialized topics of NS theory such as the role of strangeness in compact stars~\\cite{WEBER_REVIEW}, cooling of neutron stars~\\cite{PAGE_REVIEW,PETHICK}, neutrino propagation~\\cite{VOLKAS}, phases of quark matter at high densities~\\cite{RISCHKE,ALFORD,RAJAGOPAL} and the equations of state of hadronic matter~\\cite{BALDO_BOOK,LATTIMER_PRA,HH}.  The remainder of this introduction gives a brief overview of phases of dense matter in neutron stars. Chapter 2 starts with an introduction to the finite temperature non-equilibrium Green's functions theory. Further we give a detailed account of the $T$-matrix theory in the background medium and discuss the precursor effects in the vicinity of critical temperature of superfluid phase transition. Extensions of the $T$-matrix theory to describe the superfluid phases and three-body correlations are presented. A virial equation of state is derived, which includes the second and the third virial coefficients for degenerate Fermi-systems. The formal part of this Chapter includes further a discussion of the Brueckner-Bethe-Goldstone theory of nuclear matter, the covariant version of the $T$-matrix theory (known as the Dirac-Brueckner theory). We further discuss isospin asymmetric nuclear matter, onset of hyperons and meson condensates, and the conditions of charge neutrality and $\\beta$-equilibrium that are maintained in compact stars. We close the chapter by a discussion of stellar models constructed from representative equations of state.  Chapter~3 discusses the weak interactions in dense matter within the covariant extension of the real-time Green's functions formalism described in Chapter~2. We  discuss the derivation of neutrino emissivities from quantum transport equations for neutrinos and the dominant processes that contribute to the neutrino cooling rates of hadronic matter. Special attention is paid to processes that occur in the superfluid phases of neutron stars.  Chapter~4 is a short overview of cooling simulations of compact stars. It contains several example simulations with an emphasis on the processes that were discussed in Chapter~3.   \\subsection{\\it A brief overview of neutron star structure}  \\begin{figure}[tb] \\begin{center} \\epsfig{figure=ppnp_fig1.eps,width=13.0cm,angle=0} \\begin{minipage}[t]{16.5 cm} \\caption{ Schematic cross section of a $M = 1.4 M_{\\odot}$ mass neutron star. The different stable phases and their constituents (using the standard notations)  are shown as a function of the radius in the upper half of the diagram and as a function of the logarithm of the density $\\rho$ in the lower half. Note that the both scales are strongly expanded in the low-density/large-radius domain. For details see the text. } \\label{fig:NS_CROSS} \\end{minipage} \\end{center} \\end{figure}  We now adopt a top to bottom approach and review briefly the sequence(s) of the phases of matter in neutron stars as the density is increased. A schematic picture of the interior of a $M=1.4 M_{\\odot}$ mass neutron star is shown in Fig.~\\ref{fig:NS_CROSS}. The low-density region of the star is a highly compressed fully ionized matter at densities about $\\rho = 10^6$ g cm$^3$ composed of electrons and ions of $^{56}$Fe. This phase could be covered by several cm thick `blanket' material composed of H, He, and other light elements, their ions and/or molecules. The composition of the surface material is an important ingredient of the photon spectrum of the radiation, which is used to infer the surface photon luminosities of NS. Charge neutrality and equilibrium with respect to  weak processes imply that the matter becomes neutron rich as the density is increased. In the density range $10^7\\le \\rho \\le 10^{11}$ g cm$^{-3}$ a typical sequence of nuclei that are stable in the ground state is $^{62}$Ni, $^{86}$Kr, $^{84}$Se, $^{82}$Ge, $^{80}$Zn, $^{124}$Mo, $^{122}$Zr, $^{120}$Sr and their neutron rich isotopes. The matter is solid below the melting temperature $T_m\\sim 10^7$ K, the electron wave functions are periodic Bloch states, and the elementary excitations are the electron quasiparticles and phonons. The lattice may also contain impurities, i.~e. nuclei with mass numbers different from the predicted stable nucleus, as the time-scales for relaxation to the absolute ground state via weak interactions could be very large. The transport properties of the highly compressed solid (CS) are fundamental to the understanding of the way the thermal energy is transported from the core to the surface and the way the magnetic fields evolve in time.  Above the density $\\rho \\simeq 4\\times 10^{11}$ g cm$^{-3}$ not all the neutrons can be bound into clusters, and those which are free to form a continuum of states fill a Fermi-surface characterized by a positive chemical potential. Thus, the ``inner crust\" is a CS featuring a neutron fluid. The sequence of the nuclei in the inner crust are neutron rich isotopes of Zr and Sn, with the number of protons $Z=40$ and $50$ respectively and mass number in the range $100 \\le A \\le 1500 $. Below the critical temperature $T_c \\sim 10^{10}$~K, which corresponds to 1~MeV (1~MeV $= 1.602\\times 10^{10}$ K) neutrons in the continuum undergo a phase transition to the superfluid state. At low temperatures, the electron quasiparticles and the lattice phonons are the relevant degrees of freedom which control the thermal and magnetic properties of the matter. At non-zero temperatures neutron excitations out of condensate can play an important role in mass transport and weak neutral current processes.  At about half of the nuclear saturation density, $\\rho_0 = 2.8\\times 10^{14}$ g  cm$^{-3}$ the clusters merge into continuum leaving behind a uniform fluid of neutrons, protons and electrons. The uniform neutron ($n$), proton ($p$) and electron ($e$) and possibly muon ($\\mu$) phase extends up to densities of a few $\\rho_0$; the $p$, $e$ and $\\mu$ abundances are  in the range 5-10$\\%$. The many-body theory which determines the energy density of matter in this density range (as well as at higher densities) is crucial for the structure of the neutron stars, since most of the mass of the star resides above the nuclear saturation density. The neutrons and protons condense in superfluid and superconducting states below critical temperatures of the order $10^9$~K. Because of their low density the protons pair in the relative $^1S_0$ state; neutron Fermi-energies lie in the energy range where attractive interaction between neutrons is in the $^3P_2-^3F_2$ tensor spin-triplet channel. The relevant quasiparticle excitations of the $npe\\mu$-phase are the electrons and muons at low temperatures; at moderate temperatures the neutron and proton excitations out of the condensate can be important.  The actual state of matter above the densities $2-3\\times \\rho_0$ is unknown. The various possible phases are shown in Fig~\\ref{fig:NS_CROSS}. At a given density the largest energy scale for charge neutral and charged particles are the Fermi-energies of neutrons and electrons, respectively. Once these scales become of the order of the rest mass of strangeness carrying heavy baryons, the $\\Sigma^{\\pm,0}$,  $\\Lambda$, $\\Xi^{\\pm,0}$ hyperons nucleate in matter. Their abundances are again controlled by the equilibrium with respect to weak interactions and charge neutrality. Since the densities reached in the center of a massive NS are about 10$\\times\\rho_0$, it is likely that the critical deconfinment density at which the baryons lose their identity and disintegrate into  up ($u$), down ($d$), and possibly strange ($s$) quarks is reached inside massive compact stars. The critical density for the deconfinement transition cannot be calculated reliably since it lies in a range where quantum chromodynamics (QCD) is non-perturbative. NS phenomenology is potentially useful for testing the conjecture of high-density cold quark matter in compact stars.  At densities $\\rho \\ge 2\\rho_0$ Bose-Einstein condensates (BEC) of pions ($\\pi$), kaons ($K$), and heavier mesons can arise under favorable assumptions about the effective meson-nucleon and nucleon-nucleon interactions in matter. For example, the pion condensation arises because of the instability of the particle-hole nucleonic excitations in the medium with quantum numbers of pions. This instability depends on the details of the nuclear interaction in the particle-hole channel and is uncertain. There are very distinct signatures associated with the pion and kaon BEC in the physics of NS featuring such a condensate, which includes softening of the equation of state and fast neutrino cooling. ",
        "conclusions": "This review covered a number of topics that are relevant to the theory of compact stars, in particular the physics of hadronic matter with baryonic degrees of freedom and weak interaction processes involving hadrons. It should be clear from the presentation that despite the enormous progress achieved during the past decade, the theory is far from being completed and a large number of exciting topics remain to be studied in the future.  Compact stars continue to pose an enormous intellectual challenge to the physics and astrophysics communities. Rapid progress during the last four decades since the discovery of the first pulsar is an evidence of the vitality of this field. Current and planned observational programs continue the exploration of compact objects in the electromagnetic spectrum; studies of compact stars through gravity waves may become possible in the near future. All this provides an excellent basis for future theoretical studies of compact stars."
    },
    "0601/astro-ph0601117_arXiv.txt": {
        "abstract": "% The Local Interstellar Medium (LISM) is a unique environment that presents an opportunity to study general interstellar phenomena in great detail and in three dimensions.  In particular, high resolution optical and ultraviolet spectroscopy have proven to be powerful tools for addressing fundamental questions concerning the physical conditions and three-dimensional (3D) morphology of this local material.  After reviewing our current understanding of the structure of gas in the solar neighborhood, I will discuss the influence that the LISM can have on stellar and planetary systems, including LISM dust deposition onto planetary atmospheres and the modulation of galactic cosmic rays through the astrosphere --- the balancing interface between the outward pressure of the magnetized stellar wind and the inward pressure of the surrounding interstellar medium.  On Earth, galactic cosmic rays may play a role as contributors to ozone layer chemistry, planetary electrical discharge frequency, biological mutation rates, and climate.  Since the LISM shares the same volume as practically all known extrasolar planets, the prototypical debris disks systems, and nearby low-mass star-formation sites, it will be important to understand the structures of the LISM and how they may influence planetary atmospheres. ",
        "introduction": "The interstellar medium (ISM) is a critical component of galactic structure.  Its role in the lifecycle of stars, mediating the transition from stellar death to stellar birth, evokes a sense of a ``galactic ecology'' \\citep{burton04}.  The ISM provides a platform for the recycling of stellar material, by transferring and mixing the remnants of stellar nucleosynthesis and creating environments conducive for the creation of future generations of stars and planets. It also transfers energy and momentum, absorbing flows from supernovae blasts and strong winds from young stars and coupling these peculiar motions with galactic rotation and turbulence.  Ultimately, when a dense interstellar cloud collapses, it is the conservation of momentum from the parent cloud that leads to the formation of a protostellar disk from which stars and planets are formed.  The local interstellar medium (LISM) is the interstellar material that resides in close ($\\lesssim$100\\,pc) proximity to the Sun.  For a discussion on more distant ISM structures, see McClure-Griffiths, in this volume.  Proximity is a special characteristic that drives much of the interest in the LISM.  First, proximity provides an opportunity to observe general ISM phenomena in great detail, and in three dimensions.  ISM structures and processes are repeated almost {\\it ad infinitum} in our own galaxy \\citep{dickey90}, and beyond in other galaxies \\citep{mccray87}, even at high redshift \\citep{rauch99}. Knowledge of general ISM phenomena in our local corner of the galaxy, discussed in \\S2, can be applied to more distant and difficult to observe parts of the universe.  Second, proximity implies an interconnectedness.  The relationship between stars and their surrounding interstellar environment will be discussed in \\S3, with particular attention paid to the interaction of the Sun with the LISM.  In \\S4, the consequences of the relationship between stars and the LISM on planetary atmospheres are discussed, and the LISM-Earth, or more generally, the ISM-planet connection is explored in more detail.  This manuscript should not be considered a comprehensive review of the subject of the LISM, but an individual, and biased, thread through a rich research area.  I certainly will not be able to explore many LISM topics to the level they deserve, nor will I be able to highlight all the work done by the large number of researchers in this area. Hopefully, this short review will introduce you to some new ideas, and the references provided can escort you to even more work that was not specifically mentioned in this manuscript. ",
        "conclusions": "The LISM is a unique conduit that connects large scale galactic and extragalactic structures with planetary atmospheres.  Although every observation of an astronomical object outside our solar system (and even some ``within'' our solar system) peers through the LISM, we do not, as of yet, have a detailed three-dimensional global understanding of the morphological or physical properties of our own interstellar environment.  The challenges of observing the properties of the LISM are slowly being met with larger databases of high resolution spectra taken from the FUV to the optical, along sightlines toward nearby stars.  Due to the limited number of transitions that are sensitive to the low column densities of LISM material, it is critical to have access to high resolution spectrographs from the FUV to the optical. With the recent loss of the Space Telescope Imaging Spectrograph (STIS) on {\\it HST}, no high resolution astronomical spectrograph ($R\\,\\equiv\\,\\lambda/\\Delta\\lambda\\,\\geq\\,$50,000) is currently operating in the ultraviolet.  A new high resolution UV instrument is needed in order to preserve and expand our observational capabilities in the UV.  Without such a facility, among other losses, we will no longer have the ability to observe gas in the LISM, and risk being ignorant of our most immediate interstellar surroundings.  The LISM interacts with stellar and planetary systems and as we explore the subtle astrobiological parameters that control the degrees of habitability of planets, the role of the LISM may turn out to be significant.  Since the LISM, and the ISM in general, is an important part of a grand ``galactic ecology,'' research on the LISM touches many different areas of astrophysics, often at a profound level.  Carl Sagan once said that ``we are made of star-stuff'' \\citep{sagan80}. Created in stars, our ``stuff'' was mixed and transported from across the galaxy by the ISM and, eventually, it is likely that this ``stuff'' will return there to be recycled yet again."
    },
    "0601/astro-ph0601321_arXiv.txt": {
        "abstract": "We present a multi-wavelength study ({\\em BVRI} band photometry and \\hi{} line interferometry) of nine late-type galaxies selected from the HIPASS {\\it Bright Galaxy Catalog} on the basis of apparently high \\hi{} mass-to-light ratios (3\\mls{} $<$ \\mlr{} $<$ 27\\mls{}).  We found that most of the original estimates for \\mlr{} based on available photographic magnitudes in the literature were too high, and conclude that genuine high \\hi{} mass-to-light ratio ($>$5\\mls{}) galaxies are rare in the Local Universe.  Extreme high \\mlr{} galaxies like \\esoq{} appear to have formed only the minimum number of stars necessary to maintain the stability of their \\hi{} disks, and could possibly be used to constrain galaxy formation models.  They may to have been forming stars at a low, constant rate over their lifetimes.  The best examples all have highly extended \\hi{} disks, are spatially isolated, and have normal baryonic content for their total masses but are deficent in stars.  This suggests that high \\mlr{} galaxies are not lacking the baryons to create stars, but are underluminous as they lack either the internal or external stimulation for more extensive star formation. ",
        "introduction": "\\label{sec:intro}  The number of low mass dark matter halos predicted by models of a CDM dominated Universe far exceeds the quantity of observed dwarf galaxies, typically by several orders of magnitude \\citep*[see][]{kau93,moo99,kly99}.  In this context we consider a galaxy to be a dark matter halo that contains baryons.  Consequently the slope of dark matter mass functions generally rises much more steeply than observed galaxy luminosity functions (\\citealp*{tre02,hil03,bla03}; but see also \\citealp{bla04}).  While there are some physical processes that could help narrow this discrepancy \\citep*{kly99,sha04}, the theoretical low mass halo frequency is not reduced enough to reconcile them with observations.  So we are left with the conclusions that either the current most favored cosmological models significantly over-estimate the number of low mass dark matter halos present in the Local Universe, or the observations have failed to find the vast majority of low mass galaxies to date.  If the latter were true, then it is important to look at why dwarf galaxies could be missed and how we could detect them.  Two reasons why galaxies might not have been found yet in optical surveys are that they could exhibit low stellar densities ($<$ 1\\Msun{}\\,pc$^{-2}$) or most of their baryons are in invisible form.  In the extreme case, they do not contain any baryonic matter at all and are in fact ``empty'' dark matter halos.  These possibilities could be dark matter halos in which the star formation from accreted gas has been halted, suppressed, never began, or there were simply no baryons to form stars to begin with.  Such objects lacking stars might be considered as being ``dark galaxies,'' which could be easily missed in surveys biased towards optical wavelengths \\citep{dis76}.  However, dark galaxies have yet to be found in the blind \\hi{} surveys such as HIPASS \\citep{kor04,doy05}.  If they do exist, large numbers of low mass dark galaxies could naturally steepen the mass functions from observations, which are mostly derived from optical or near-infrared galaxy luminosity functions without consideration of other baryonic matter.  Recent large scale observations of neutral hydrogen gas (\\hi{}) are now allowing mass functions to be derived based on non-stellar properties.  \\citet{zwa03} produced one of the most extensive \\hi{} mass functions based on a catalogue of the 1000 \\hi{}-brightest galaxies in the Southern hemisphere and found a similar low mass end slope to other observational studies, again in contradiction with $\\Lambda$CDM models.  Galaxies that have failed to convert most of their primordial gas into stars, and thus retained a high proportion of \\hi{}, may well provide a partial solution.  While not entirely ``dark'' these galaxies are hard to detect optically, but may be detectable through 21cm line observations.  The \\hi{} mass-to-light ratio (\\mlr{}, which compares the \\hi{} mass to the {\\em B} band luminosity) of these objects could be significantly higher than the typical ratios measured for late-type galaxies, so that they would be in a lower mass bin in a luminosity function than they would be for a baryonic mass function.  If high \\mlr{} galaxies existed in significant numbers then they could help correct the discrepancy in two ways, by including more previously unknown galaxies, and by shifting known galaxies to higher mass bins than they would be placed in with purely optical results.  An example of an extreme \\hi{} mass-to-light ratio object, which could be described as a ``dim'' galaxy, is the nearby dwarf irregular \\esoq{} with \\mlr{} = $22 \\pm 4$\\mls{} \\citep*[][ hereafter \\pI{}]{war04}.  This faint low surface brightness dwarf irregular was found to be spatially isolated (1.7~Mpc from the nearest neighbor), with a low current star formation rate ($\\la 2.5 \\times 10^{-3}$\\Msun{}\\,yr$^{-1}$).  It has an extended regularly rotating \\hi{} disk, which can be traced out to over six times the Holmberg radius of the optical galaxy, making it one of the most extended \\hi{} envelopes relative to the optical extent.  \\pI{} included an analysis of the \\hi{} gas surface density of \\esoq{}.  The azimuthally averaged surface density at all radii was below the critical gas surface density needed for large scale star formation as defined by the \\citet{too64} stability criteria \\citep{ken89,mar01}.  It was proposed in \\citet{ver02} that a large fraction of low mass halos may form \\citeauthor{too64} stable gas disks and become ``dark'' galaxies, possibly 95\\% of objects with halo masses of $\\la 10^{10}$\\Msun{}.  If so \\esoq{} would be just the tip of the iceberg.  If we can find more galaxies similar to \\esoq{}, where the gas density after gravitational collapse is too low for efficient star formation, it may go some way to explaining the discrepancy between the dark matter halo mass function and the observed galaxy luminosity function.  To continue our study of the stellar and gas properties of galaxies with high \\mlr{} we have selected a sample of nine galaxies from the BGC in the approximate range 3\\mls{} $<$ \\mlr{} $<$ 27\\mls{}. In this paper, \\S~\\ref{sec:sample} looks at what was previously known about the sample galaxies. \\S~\\ref{sec:obs} summarizes our 21cm and optical observations.  \\S\\S~\\ref{sec:radio} and \\ref{sec:optp} present the results of the \\hi{} line observations and optical photometry, respectively.  \\S~\\ref{sec:dis-indi} compares the properties of individual galaxies.  \\S~\\ref{sec:dis} contains the discussion of these results and the possible implications, while \\S~\\ref{sec:conc} gives our conclusions. ",
        "conclusions": "\\label{sec:conc}  We obtained accurate optical CCD apparent magnitudes and \\hi{} flux densities for nine late type dwarf galaxies and recalculated their \\hi{} mass-to-light ratios.  The new \\mlr{} values are: \\begin{itemize} \\item $22 \\pm 4$\\mls{} for \\esoq{}, \\item $4.8 \\pm 1.1$\\mls{} for ESO\\,572-G009, \\item $4.5 \\pm 0.9$\\mls{} for ESO\\,428-G033, \\item $3.0 \\pm 0.3$\\mls{} for \\mcg{}, \\item $2.8 \\pm 0.6$\\mls{} for ESO\\,473-G024, \\item $2.6 \\pm 0.2$\\mls{} for IC\\,4212, \\item $1.6 \\pm 0.2$\\mls{} for ESO\\,348-G009, \\item $1.5 \\pm 0.1$\\mls{} for ESO\\,121-G020, and \\item $1.2 \\pm 0.1$\\mls{} for ESO\\,505-G007. \\end{itemize} Many of these \\hi{} mass-to-light ratios are significantly below the initial estimates, due to inaccurate magnitude estimates in the literature. This strongly emphasises the importance of having accurate observations in both the \\hi{} line and the optical.  Based on the new \\hi{} mass-to-light ratio distribution we conclude that genuine ``dim'' galaxies with high ratios (\\mlr{}$>$5\\mls{}) are rare in the local Universe.  A previously uncatalogued companion galaxy to ESO\\,121-G020 was found at a projected distance of 3\\arcmin{}.  \\atg{} has an \\hi{} mass of $\\sim10^{7}$\\Msun{} and \\mlr{} of $2.2 \\pm 0.3$\\mls{}.  This was the only such companion detected, and is well within the beam of the Multibeam instrument used by the HIPASS survey.  Despite our low resolution \\hi{} observations we were able to separate the galaxy ESO\\,505-G007 from its close neighbor ESO\\,505-G008 and determined \\hi{} flux densities of $21 \\pm 3$ and $8 \\pm 3$\\jks{}, respectively.  The best examples of high \\mlr{} dwarf galaxies in the literature all have highly extended \\hi{} disks, are spatially isolated and have normal baryonic content for their dynamical masses.  The galaxies are not lacking the baryons to create stars, but are underluminous as they lack either the internal or external stimulation for further star formation.  Future examination of element abundances within star formation sites of high \\mlr{} galaxies may reveal important clues about their star formation history.  Recent observations \\citep{ski03b,ken01} support the idea that high \\mlr{} galaxies may have been forming stars at a low, constant rate over their lifetimes as proposed in \\pI{}.  There may be a minimum quantity of stars a galaxy will form that depends on the initial baryonic mass, which is supported by the theoretical work of \\citet{tay05}.  If this is true then maybe high \\hi{} mass-to-light ratio galaxies have over their lifetimes only formed the minimum number of stars necessary to maintain the stability of their \\hi{} gas disk."
    },
    "0601/astro-ph0601103_arXiv.txt": {
        "abstract": "Flat radio spectra with large brightness temperatures at the core of AGN and X-ray binaries are usually interpreted as the partially self-absorbed bases of jet flows emitting synchrotron radiation. Here we extend previous models of jets propagating at large angles to our line of sight to self-consistently include the effects of energy losses of the relativistic electrons due to the synchrotron process itself and the adiabatic expansion of the jet flow. We also take into account energy gains through self-absorption. Two model classes are presented. The ballistic jet flows, with the jet material travelling along straight trajectories, and adiabatic jets. Despite the energy losses, both scenarios can result in flat emission spectra, however, the adiabatic jets require a specific geometry. No re-acceleration process along the jet is needed for the electrons. We apply the models to observational data of the X-ray binary Cygnus X-1. Both models can be made consistent with the observations. The resulting ballistic jet is extremely narrow with a jet opening angle of only 5\". Its energy transport rate is small compared to the time-averaged jet power and therefore suggests the presence of non-radiating protons in the jet flow. The adiabatic jets require a strong departure from energy equipartition between the magnetic field and the relativistic electrons. These models also imply a jet power two orders of magnitude higher than the Eddington limiting luminosity of a 10\\,M$_{\\odot}$ black hole. The models put strong constraints on the physical conditions in the jet flows on scales well below achievable resolution limits. ",
        "introduction": "The centres or cores of many AGN show a flat radio spectrum in the sense that for the flux density as a function of frequency $\\nu$ we observe $F_{\\nu} \\propto \\nu^{\\alpha}$ with $\\alpha \\sim 0$. The high surface brightness temperature associated with these spectra suggests a synchrotron origin of the emission. Observations with high spatial resolution reveals that the flat spectrum arises in the base of jet flows which continue to much larger scales \\citep[for a review see][]{tc91}. Similar flat or inverted ($\\alpha > 0$) radio spectra are also observed in X-ray binaries in the low-hard state \\citep[e.g.][]{rf01}. If optically thin, the flat synchrotron spectrum would imply a power-law energy distribution of the radiating relativistic electrons with a slope of unity. Such a distribution is very unlikely to arise for the usually assumed mechanism for the acceleration of the electrons at shock fronts \\citep[e.g.][]{ab78}.  A magnetised plasma containing very energetic electrons with a power-law energy distribution will produce a power-law spectrum at high frequencies. The slope of the spectrum, typically $\\alpha < 0$, is determined by the slope of the energy distribution. However, below a critical frequency the radiating electrons will re-absorb some of the photons. In this self-absorbed, optically thick regime the spectrum has a power-law slope of $5/2$, independent of the slope of the electron energy distribution \\citep[e.g.][]{rl79}. The spectrum of a uniform, self-absorbed synchrotron source therefore shows a pronounced peak. \\citet{bk79} pointed out that in a jet the plasma conditions are changing along the flow and therefore the peaks of the self-absorbed spectra of different parts of the jet can occur at different frequencies. If the plasma conditions change such that the spectra peak at the same level, then the overall spectrum, observed with a spatial resolution insufficient to resolve the individual parts of the jet, will be flat. Their model has become the standard tool for interpreting observations of flat radio spectra from jetted sources.  In the \\citet{bk79} model the jet is assumed to have a conical geometry, i.e. the velocity with which the jet is expanding sideways, is constant. The bulk velocity of the jet material along the jet axis is also assumed to be constant. The magnetic field is assumed to be directed perpendicular to the jet axis and `frozen' into the jet plasma. Adiabatic losses of the electrons are mentioned by the authors, but are assumed to be replenished by an unknown, continuous re-acceleration process along the entire jet. The same assumption is made for radiative energy losses associated with the emission of synchrotron radiation. The subsequent model of \\citet{am80} includes a simplified treatment of energy losses of the electrons due to adiabatic expansion and radiative processes. It also allows for more confined jets, i.e. the jets are not necessarily conical. With these assumptions, the model is unable to produce a flat emission spectrum. A similar model was developed by \\citet{hj88}. They consider adiabatic, but not radiative, energy losses of the electrons. The jet geometry is again conical, but they also investigate a more confined jet. The model can predict flat spectra, but \\citet{hj88} point out that these may only arise under special circumstances, particularly in the case of confined jets. The model of \\citet{gm98} includes a detailed treatment of the energy losses of the relativistic electrons, but it concentrates only on the optically thin part of the spectrum of jets propagating close to the line of sight for which numerical solutions are presented. The perhaps most comprehensive study of jet emission models is that of \\citet{sr82} which includes the effects of energy losses on the electron population, but neglects the effects of self-absorption on the electron energy spectrum.  In this paper we extend the previous models by including adiabatic and radiative energy losses and gains (due to absorption) for the electrons as well as investigating various possibilities for the evolution of the magnetic field. We consider two distinct cases: The ballistic and the adiabatic jet models. In the ballistic case the jet material follows straight trajectories and does not behave like a fluid, because individual fluid elements do not interact with each other. In many ways this model is similar to the \\citet{bk79} model, but we show that because of self-absorption effects we do not need to invoke a re-acceleration process to achieve flat emission spectra. in the adiabatic jet model the relativistic electrons suffer from adiabatic energy losses as well as radiative losses. Again we show that the models can produce flat spectra without re-acceleration of the electrons, but only for a very specific jet geometry. The emphasis of our treatment is on the construction of analytical models and so we concentrate on jets propagating at large angles to the line of sight, i.e. the viewing angle is larger than the inverse of the Lorentz factor of the jet flow.  In Section \\ref{conical} we briefly discuss the basic properties of our jets in terms of their geometry, the evolution of the magnetic field and that of the relativistic electrons. We present the first fully analytical solution of the equations governing partially self-absorbed synchrotron emission from a jet in Section \\ref{radiation}. Section \\ref{simple} summarises the model results for the case without radiative energy losses as studied in many previous models. In Section \\ref{cutoff} we develop the formalism for including radiative energy losses in the model and the resulting spectra are discussed in Section \\ref{losses}. We apply the model to the data obtained for Cygnus X-1 in Section \\ref{obs} and derive the properties of this jet. Finally, we summarise our conclusions in Section \\ref{conc}. ",
        "conclusions": "\\label{conc}  We construct a model for the synchrotron emission of partially self-absorbed jets. The model does not invoke a re-acceleration process for the relativistic electrons. All electrons are accelerated only once at the lower physical limit of the jet, $l_{\\rm min}$. It is not necessary to postulate an unknown re-acceleration mechanism as was suggested in previous work \\citep{bk79}, because synchrotron self-absorption counteracts excessive energy losses.  Two classes of models are considered. The ballistic jets have a conical geometry and their contents do not suffer adiabatic energy losses. This situation may arise when jets are initially highly overpressured with respect to their environments. They expand unimpeded and random thermal energy is converted into bulk kinetic energy, but not dissipated to any external medium. At the end of this very rapid expansion the mean free path of the jet material exceeds the physical dimensions of the jet itself and follows ballistic trajectories.  The other class of models considered here are adiabatic jets confined by either the pressure of the gas surrounding them, surface magnetic fields or both. Their geometrical shape is dictated by the details of the confinement mechanism which is not the subject of this paper. The jet material dissipates energy to the external gas during its adiabatic expansion.  Both classes of models can predict flat emission spectra if energy losses of individual electrons are neglected. The ballistic jet with a magnetic field perpendicular to the jet axis produces a flat spectrum without further assumptions. The adiabatic jet models require a specific jet geometry to allow for flat emission spectra.  Both model classes can predict flat emission spectra, even when taking energy losses of the electrons and the magnetic field into account. Synchrotron self-absorption prevents radiative energy losses below a critical Lorentz factor $\\gamma _{\\rm thick}$. In the ballistic case $\\gamma _{\\rm thick}$ is constant along the jet flow and because adiabatic energy losses are absent, the energy distribution of the relativistic electrons remains stationary. The emission properties of the jet in this case are essentially identical to the model of \\citet{bk79}. For the adiabatic jet the with a perpendicular magnetic field electrons with Lorentz factors at and below $\\gamma _{\\rm thick}$ continue to lose energy because of the sideways expansion of the jet. However, for a geometrical shape given by $r \\propto x^{1/5}$ the resulting spectrum is again flat. Other configurations of the magnetic field can also lead to flat emission spectra. The spectral slope is very sensitive to the jet shape. For example, slightly changing the jet shape to $r \\propto x^{1/4}$ results in a spectrum with $F_{\\nu} \\propto \\nu^{0.38}$.  We show an application of the model to observations of a resolved jet in the X-ray binary Cygnus X-1 \\citep{ssf01}. As input we use the flux density of the jet in the flat part of the spectrum, the physical extent of the resolved jet and an assumed break of the flat spectrum to optically thin conditions in the near-IR. Both model classes can be made consistent with the observational constraints. However, in doing so the adiabatic jet models require a significant departure from energy equipartition between the magnetic field and the relativistic particles. The associated reduced radiative efficiency of the jet plasma implies extremely high energy transport rates for the jet of around $10^{34}$\\,W. This jet power exceeds the Eddington limiting luminosity of a 10\\,M$_{\\odot}$ black hole by two orders of magnitude.  The ballistic jet is consistent with current observations and requires energy transport rates well below the time-averaged jet power \\citep{gfk05}. This result holds even if the kinetic energy of one non-radiating proton per relativistic electron is taken into account. Unless Cygnus X-1 ceases to produce a jet for long periods of time, then the ballistic model requires a proton-electron jet plasma to explain the large accumulated energy in the region where the jet interacts with the surrounding gas.  The jet opening angle of 5\" required by the ballistic model is very small. The velocity perpendicular to the jet axis established in the jet acceleration region must be very small for this model to work. If the ballistic jet results from a rapid initial expansion of a highly overpressured jet, then the requirements are even more severe. This result places stringent limits on the acceleration process.  Both model classes can probe the conditions in jets on scales unresolvable with current telescopes. The most important measurement in this respect is the high frequency cut-off of the flat spectrum which is formed closest to the jet acceleration region. In order to resolve the remaining uncertainties with the model, the most helpful measurements would be to resolve the jet along its axis in quasi-simultaneous observations at two different frequencies. As mentioned above, even a factor 2 between the observing frequencies employed would help to constrain the geometrical shape of the jet and thereby help to decide which of the model classes is compatible with observations.  Obviously from the work presented here we cannot rule out the possibility of continued re-acceleration of the radiating electrons in the jet. However, such an energetisation process is not necessary to produce flat spectra from self-absorbed, synchrotron emitting jets."
    },
    "0601/astro-ph0601273_arXiv.txt": {
        "abstract": "Two disks of young stars have recently been discovered in the Galactic Center. The disks are rotating in the gravitational field of the central black hole at radii $r\\sim 0.1-0.3$~pc and thus open a new opportunity to measure the central mass. We find that the observed motion of stars in the clockwise disk implies $M=(4.3\\pm 0.5)\\times 10^6M_{\\odot}$ for the fiducial distance to the Galactic Center $R_0=8$~kpc and derive the scaling of $M$ with $R_0$. As a tool for our estimate we use orbital roulette, a recently developed method.  The method reconstructs the three-dimensional orbits of the disk stars and checks the randomness of their orbital phases. We also estimate the three-dimensional positions and orbital eccentricities of the clockwise-disk stars. ",
        "introduction": "A supermassive black hole is believed to reside in the dynamical center of the Galaxy (Eckart \\& Genzel 1996; Genzel et al. 1997; Ghez et al. 1998; Sch\\\"odel et al. 2002; Genzel et al.~2000; Ghez et al.~2000; Sch\\\"odel et al.~2003; Ghez et al.~2003). It is associated with the compact radio source \\Sgr and dominates the gravitational potential at distances $r< 1$~pc from the center. The central mass has been estimated with a few different methods. In particular, the statistical analysis of sky-projected positions and velocities of stars in the central 0.5~pc gave estimates that range above $3\\times 10^6M_\\odot$ (Genzel et al. 2000; Sch\\\"odel et al.~2003). The systematic error of these estimates, however, is difficult to quantify as they depend on the assumed three-dimensional statistical model of the cluster.  Independent estimates of the black hole mass were derived from observations of a few stars (S-stars) that move on very tight orbits around \\Sgr, at radii $r\\sim 0.01$~pc, and have orbital periods as short as a few decades. Fractions of S-star orbits have been mapped out over the past decade (Sch\\\"odel et al.~2002; Sch\\\"odel et al.~2003; Ghez et al.~2003; Ghez et al.~2005). By checking which value of the central mass provides an acceptable fit to the orbital data, Eisenhauer et al. (2005) found $M=(4.1\\pm 0.4)\\times 10^6M_\\odot$, and Ghez et al. (2005) found $M=(3.7\\pm 0.2)\\times 10^6M_\\odot$ for the fiducial distance to the Galactic Center $R_0=8$~kpc. Same star orbits were used to derive $R_0=7.62\\pm 0.32$~kpc (Eisenhauer et al. 2005), in agreement with previous methods (e.g. Reid et al. 1993).  In this paper, we report a new independent estimate of the central mass that uses the young massive stars at $r\\sim 0.1-0.3$~pc as test particles. A remarkable property of these stars has been discovered recently: they form a disk-like population (Levin \\& Beloborodov 2003; Genzel et al. 2003). Two stellar disks apparently coexist at these radii. One of them rotates clockwise on the sky and the other one --- counter-clockwise. New observations of Paumard et al. (2006) have increased the number of known disk members by a factor of three and show clearly the two-disk structure of the young population. The clockwise disk is especially well pronounced. It has about 30 members and a small thickness $H/r \\approx 0.1$.  In this paper, we investigate the orbital motions of stars in the clockwise disk and infer the central gravitating mass. As a tool of our mass estimate we use orbital roulette, a recently developed statistical method (Beloborodov \\& Levin 2004, hereafter BL04). This method is generally more precise than the estimators based on the virial theorem and provides a statistically well-defined way to quantify the uncertainty of the mass estimate. ",
        "conclusions": "The reconstructed orbits of stars in the clockwise disk depend on the assumed central mass $M$. In particular, the star orbital phases depend on $M$. We have used this dependence to estimate $M$: the true $M$ must give time phases consistent with a uniform distribution. This is the first practical application of orbital roulette, a new method of estimating the gravitational potential from instantaneous motions of test bodies (BL04).  The obtained estimate $M=(4.3^{+0.5}_{-0.4})\\times 10^6M_{\\odot}$ is consistent with the recent estimates based on the maps of orbits of S stars in the very vicinity of the center ($r<0.01$~pc). The independent estimate presented here uses different stars at $r=0.1-0.3$~pc and is complementary to the small-scale estimates. It measures the total mass within $r\\sim 0.1$~pc which may include stars and dark matter.  The approximate equality of the masses within the central 0.1~pc and 0.01~pc implies an upper bound on the mass between 0.01 and 0.1~pc: $\\Delta M \\simlt 0.5\\times 10^6M_\\odot$. This constrains a possible cusp of the stellar population around the cenral black hole as well as the dark-matter content of the central region. Further observations may provide a better constraint or resolve $\\Delta M$.  The data shows an inner cutoff of the disk population at a radius $r\\sim 0.05$~pc (see Fig.~1 and Paumard et al. 2006). The observed peak of the population at $r\\sim 0.1$~pc supports the hypothesis that the young stars were born in a gravitationally unstable accretion disk (LB03).  The found geometrical thickness of the clockwise stellar population $H/r=0.09\\pm 0.03$ (eq.~\\ref{eq:H}) does not imply that it was born with this $H/r$. The parent accretion disk was likely thinner, and its stellar remnant have thickened with time. The age of the stellar population, $t_{age}=(6\\pm 2)$ million years, should be compared with the timescale of diffusion of orbital parameters. The latter may be written as \\be t_{diff}\\sim 0.1 \\frac{v^3}{G^2m_*\\rho_*\\ln(N_*/2)}, \\ee where $m_*$ is a characteristic mass of the objects that dominate fluctuations of the gravitational potential, $\\rho_*$ is their mass density averaged over the considered region $r\\sim 0.3$~pc, and $N_*\\sim (4\\pi/3)r^3\\rho_*/m_*$ is the total number of these objects. In a uniform stellar core of total mass $\\sim 3\\times 10^5M_\\odot$, the fluctuations would be dominated by low-mass stars, whose total number is $N_*\\sim 10^6$. Substituting the observed characteristic velocity, $v\\sim 300$~km/s, one would find $t_{diff}\\sim 10^{10}$~yr$\\gg t_{age}$.  However, the Galactic Center is not uniform. The two stellar disks and stellar clusters are observed, which significantly perturb the gravitational potential. These perturbations should strongly reduce $t_{diff}$. For example, if 1/30 of the core mass resides in $N_*\\sim 10$ objects with $m_*\\sim 10^3M_\\odot$ then $t_{diff}\\sim 10^8$~yr. The corresponding acquired thickness of the stellar disk is $H/r\\sim t_{age}/t_{diff}\\sim 0.1$ in agreement with the observed value. A similar estimate for $t_{diff}$ may be derived from the interaction between the two young disks (Nayakshin et al. 2005).  The eccentricities of disk stars must also diffuse with time. If the stars were born on circular orbits then $e\\sim 0.1$ are expected to develop as the disk thickens to $H/r\\sim 0.1$. Table~1 shows that a significant fraction of the young stars in the clockwise disk have relatively large eccentricities. The data are at best marginally consistent with the hypothesis that the stars used to have circular orbits 6 million years ago, and we leave a detailed analysis to a future work. Note that the stars may have been born in a gaseous disk that was not completely circularized: the circularization time was longer than the timescale for star formation. Then the initial eccentricities $e$ of the stars are non-zero. Besides, the interaction of stars with the disk might have influenced $e$ before the disk was gone.  An alternative scenario assumes that the stars were originally bound to an inspiralling intermediate-mass black hole (IMBH; Hansen \\& Milosavljevic 2003), so that the observed disk plane is associated with the orbital plane of the IMBH. In this scenario, the inspiralling young cluster was stripped by the tidal field of \\Sgr and left a disk of stars with a small dispersion of orbital planes. Calculations of Levin, Wu, \\& Thommes (2005) show that the stripped stars can generally have eccentric orbits. However, the inspiralling cluster scenario appears to be in conflict with the observed stellar distribution in the clockwise disk (see also Paumard et al. 2006 for discussion)."
    },
    "0601/astro-ph0601045_arXiv.txt": {
        "abstract": "{We present the identification of the optical counterparts to the low-mass X-ray binaries 1A~1246$-$588 and 4U~1812$-$12. We determine the X-ray position of 1A~1246$-$588 from \\emph{ROSAT}/PSPC observations and find within the error circle a blue star with $V=19.45$, $B-V=0.22$ and $R-I=0.22$ which we identify as the counterpart. Within the \\emph{Chandra} error circle of 4U~1812$-$12, a single star is present which appears blue with respect to the stars in the vicinity. It has $R=22.15$, $R-I=1.53$. Distance estimates for both systems indicate that the optical counterparts are intrinsically faint, suggesting that they are ultra-compact X-ray binaries. These identifications would increase the number of candidate ultra-compact X-ray binaries from 2 to 4, whereas orbital periods are measured for only 7 systems in the Galactic disk. ",
        "introduction": "The canonical low-mass X-ray binary consists of a neutron star or black hole and a low-mass main-sequence or (sub)giant donor star, and has an orbital period longer than one hour, up to several hundred days.  Recently, it has been found that the class of ultra-compact low-mass X-ray binaries makes up about half (5 out of 12) of the low-mass X-ray binaries in globular clusters (e.g., review by \\citealt{vl04}), whereas a growing number of such systems is also discovered in the Galactic disk (7 with measured periods, and 2 candidates, see e.g., \\citealt{jc03,njmk04,wc04}). In addition, observations with the Wide Field Cameras (WFCs) of \\emph{BeppoSAX} have found a new class of low-mass X-ray binaries, bursters with (very) low persistent X-ray emission \\citep{cbn+01}. The distribution of the sources in this class is more concentrated towards the Galactic center than that of the canonical low-mass X-ray binaries \\citep{cvz+02}.  To elucidate the evolutionary status and history of these systems, observations at longer wavelengths, in particular optical/infrared, are crucial. Such observations may reveal the orbital period, directly or indirectly (e.g., \\citealt{vpm94}), or provide information about the donor and its chemical composition (e.g.\\ \\citealt{njmk04}). The first step to such studies is the optical identification of the X-ray binary. In this letter we report two new optical identifications, which may be ultra-compact X-ray binaries. ",
        "conclusions": "\\label{sec:discussion} We have identified the optical companions to the low-mass X-ray binaries 1A~1246$-$588 and 4U~1812$-$12 based on their positional coincidence with the X-ray position and their colours. The counterpart to the first has $V=19.45$, $B-V=0.22$, while that of 4U~1812$-$12 is somewhat fainter at $R=22.15$, $R-I=1.53$.  Due to its position somewhat out of the Galactic plane, the hydrogen absorption column $N_\\mathrm{H}$ towards 1A~1246$-$588 is moderate and suggests $A_V=1.7$ \\citep{ps95}. This is smaller than the maximum absorption in this line-of-sight, which is predicted to reach $A_V=1.9$ around $d=7$\\,kpc by the model of \\citet{dcl03}. This limit constrains the absolute magnitude of the companion of 1A~1246$-$588 to $M_V\\ga3.5$. Though this reasoning assumes that both the model by \\citet{ps95} and \\citet{dcl03} are correct, this distance is in agreement with estimates from \\emph{BeppoSAX}/WFC and RXTE/ASM observations of photospheric radius expansion bursts of 1A~1246$-$588, which suggest a distance of 5\\,kpc (in't Zand et al.\\ in prep.).  For 4U~1812$-$12, the photospheric radius expansion bursts that have been observed by \\citet{cbn+00} provide an estimate on the distance to this LMXB.  Assuming an Eddington peak luminosity of $L_\\mathrm{X}=3.8\\times10^{38}$\\,erg\\,s$^{-1}$ for the accretion of helium-rich material \\citep{khz+03}, the distance is estimated at 4.6\\,kpc.  If hydrogen-rich material is accreted instead, the Eddington luminosity of $\\sim\\!2\\times10^{38}$\\,erg\\,s$^{-1}$ reduces the distance to 3.4\\,kpc (see also \\citealt{jn04}). Furthermore, the value of $N_\\mathrm{H}$ derived from the X-ray absorption suggests $A_V=6.4$ \\citep{ps95} and, using the relative extinction coefficients of \\citet{sfd98}, $A_R=5.2$. As such, the optical companion of 4U~1812$-$12 has an absolute $R$-band magnitude in the range of 3.6--4.2. If we assume that the counterpart has the same intrinsic colours as the counterpart of 1A~1246$-$588, which has $(V-R)_0=-0.2$, this would translate to $M_V=3.4$--4.0.  These absolute magnitudes place both systems amongst the intrinsically fainter of the LMXBs known. According to \\citet{vpm94}, this suggests that these systems are ultra-compact X-ray binaries (UCXBs; having an orbital period below an hour). Here, the faintness of the counterpart is due to the reprocession of X-rays in a physically small accretion disk. However, these systems remain candidate UCXBs until the orbital period is determined, i.e.\\ either through optical/IR or X-ray observations.  If these systems indeed turn out to have ultra-compact orbits, it is interesting to note that for the observed X-ray luminosities of $L_\\mathrm{X}\\approx0.9\\times10^{36}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$ (1--10\\,keV) for 4U~1812$-$12 and $L_\\mathrm{X}\\la10^{36}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$ (0.1--2.4\\,keV) for 1A~1246$-$588, these systems satisfy the notion presented by \\citet{zcm05}; that LMXBs with persistent luminosities with $L_\\mathrm{X}\\la10^{36}$\\,erg\\,cm$^{-2}$\\,s$^{-1}$ may be ultra-compact X-ray binaries."
    },
    "0601/astro-ph0601209_arXiv.txt": {
        "abstract": "{We have investigated the problem of the distribution of both masses and orbital radii of planets resulting from the gas-accretion, gas-capture model.  First we followed the evolution of gas and solids from the moment where all solids are in the form of small grains to the stage when most of them are in the form of planetesimals for a set of different initial masses and sizes of protoplanetary disks. Based on that we performed Monte Carlo calculations describing the formation of giant places at different locations. We included the effects of type II migration and growth of the mass of the planet after the gap opened. We discuss how these effects influence the final distribution of giant planets. We show that when the giant planets are not able to migrate or grow in mass after the gap opens, their distribution is mainly determined by the properties of the gaseous disk. However, with those two effects included, reproducing the parameters of the gaseous disks from the distribution of planets becomes difficult. We also checked the roles of both the material of which the solids consist and the mass of the central star. The main result is that, in disks around less massive stars, giant planets at the given location tend to be less massive. At the same time, the giant planets with the given mass tend to form closer to the less massive stars.}   \\begin{document} ",
        "introduction": "Current radial velocity surveys have led to the discovery of over 150 extrasolar planets around main-sequence stars \\citep[see][]{marcy05,mayor04}. The large majority of these planets  are probably gas giant planets, as their masses are above $100 M_\\oplus$. Such a collection provides a good data set for comparing predictions of theories of planet formation \\citep[e.g.][]{ida04a,ida04b,alibert05,kac05}.  The standard model for the formation of giant planets is the core-accretion, gas capture-model. The numerical calculations of \\citet{pollack96} show, that the formation of a giant planet in this model can be divided into three phases. During the first one, the solid core is formed by the collisional accumulation of planetesimals. The second phase begins when the core reaches the mass of a few Earth masses and starts to accrete a significant amount of gas. During this phase the envelope stays in quasi - static and thermal equilibrium, as the energy radiated by the envelope is compensated for by the energy released by accreted planetesimals. As during this phase the protoplanet accretes gas with a higher rate than solids, the mass of the envelope finally becomes equal to the mass of the core. At this moment phase 3 begins, during which the planet rapidly grows in mass by runaway accretion of gas. The final mass and location of a giant planet are determined by its gravitational interaction with its environment. As it grows in mass it induces spiral waves in the gaseous disk. This leads to the transfer of angular momentum resulting in the inward migration of the planet and possibly in the gap opening \\citep{lin86,lin96,ward97}. This last phenomenon strongly reduces the further growth of the planet.   The main problem with that scenario is related to the timescale required to form a giant planet in it. In general it is the same order of magnitude as the lifetime of the protoplanetary disk, and it is not necessarily certain if the giant planet is able to form before dispersion of the disk.  Close to the star, the formation time of a giant planet in the gas capture model is determined by phase 2, while at larger distances ($\\gtrsim 10 \\mr{AU}$) the lengths of phases 1 and 2 become comparable. The lengths of these two phases depend on the initial surface density of the planetesimal swarm in the given location. The larger the density, the faster the core grows and reaches higher mass at the end of phase 1. With the higher mass of the core, the length of phase 2 also diminishes. In general, at every distance from the star there is a critical value of the surface density of planetesimals which enables formation of a giant planet before dispersion of the protoplanetary disk; for a more detailed discussion see \\citet{kac_mass}.  However, the density of the protoplanetary swarm is not in a simple relation with the density of gas in the disk from which it emerges. While the gaseous component evolves in a viscous way, the evolution of the solid component is governed by processes of gas drag, coagulation, sedimentation, and evaporation \\citep{weiden93}.  A significant redistribution of solid material takes place in the result, and in the inner parts of the disk its surface density can be significantly enhanced compared to the initial value \\citep{sv97,weiden03}. Consequently, analysis of the possible masses and orbits of giant planets resulting from the core accretion scenario should also include the global evolution of solids in protoplanetary disks. A simple model of this evolution was proposed by \\citet{kac1} and further extended by \\citet{kac2} and \\citet{kac05} to include subsequent formation of giant planets.   In this paper we extend our models to also include the effects of migration and the gap opening by the planet. It enables us for the first time to characterize the distribution of planetary masses and final locations resulting from the core accretion model, not only including  the growth of the planet but also the preceding evolution of solids, which determines the surface density of the planetesimal swarm. In Sect. \\ref{s:methods} we explain our approach to the evolution of the protoplanetary disk and planet formation. The results are presented in Sect. \\ref{s:results} and discussed in Sect. \\ref{s:concl}. ",
        "conclusions": "\\label{s:concl} We have presented the results of simulations describing the formation of giant planets. Our model include all important phases of the core-accretion scenario, beginning from the formation of the planetesimal swarm from smaller solids, until the gap opening in the disk and subsequent planet migration. We presented the distribution of $a-M_\\mr{p}$ for planets resulting from our models for a set of different initial masses and sizes of the protoplanetary disk and for the place of planet formation. However, our results cannot be compared in a simple way to data for extrasolar planets, because of the lack of knowledge of the distribution of the initial parameters of protoplanetary disks that occur in nature. Instead our aim was to characterize the relationship between the final orbital radii and masses of planets that are in general possible to obtain and to check how the resulting $a-M_\\mr{p}$ distribution changes when including different physical processes such as migration and gas accretion by the planet after gap opening. We also investigated whether the distribution of giant planets is mainly determined by the parameters of the gaseous disk or the distribution of the planetesimal swarm.  We have shown that, if we do not include the effects of migration nor accretion of gas by the planet after gap opening, the $a-M_\\mr{p}$ distribution for giant planets is mainly determined by the parameters of the gaseous disk. In that case one is able to read the information about the dependence of the scale height of the gaseous disk on the distance from the star from the distribution of giant planets. The minimal and maximal distance from the star gives information about the surface density of planetesimal swarm at those places.   By including migration in our models, reconstructing the the properties of the gaseous disk from the distribution of giant planets becomes more difficult. The distance that the  planet can move inward depends not only on the gaseous environment but also on the time between gap opening  and dispersion of the disk. But the first of these two moments is determined by the surface density of planetesimal swarm. This, on the other hand, is not related in a simple way to the local parameters of the gaseous disk. Our simulations show that the simple reconstruction of the scale height of the gaseous disk is possible only from distribution of giant planets at greater distances than $\\sim 8\\,\\mr{AU}$, as planets at such large orbits did not significantly move away from their original places.  If we include in our models  that the planet can still grow in mass after gap opening, it is even more difficult, if at all possible, to derive conclusions about the gaseous disk (its temperature, scale height, etc.) based on the distribution of giant planets.  This is because the rate of this additional accretion onto the planet depends mainly on the accretion rate in the disk, while the time frame in which this accretion can occur is determined by the surface density of the planetesimal swarm.  The effect of the additional growth of the planet can even lead  to reversion of the correlation between the masses of the planets and semimajor axes of their orbits.  We also checked how the  mass of the central star influences the distribution of planets. The main result is that in disks around less massive stars, giant planets at the given location tend to be less massive. At the same time, the giant planets with a given mass tend to form closer to the less massive stars.  Our models can be seen as an extension of previous calculations performed by \\citet{ida04a} and \\citet{alibert05}. In comparison, the surface density of the planetesimal swarms in our models are computed self-consistently. This additional factor makes it much more difficult to predict the  structure of the gaseous disk from the distribution of giant planets much more difficult. Nevertheless it does seem necessary. Whenever we performed the similar calculations as described above, but with the assumption that the surface density of the planetesimal swarm is always a constant fraction of the surface density of gas, we were not able to perceive any giant planets. This seems  contradict  the results of those two papers. A possible reason for this different result may be the difference in the underlying models of the gaseous disk. However, we think that the main factors are the different times for the dispersion of the gaseous disk (3 Myr vs. 10 Myr) and the different sizes of solids at time 0 (small grains vs. planetesimal sizes). Our models show that the time needed to reach planetesimal sizes can be as long 1~Myr and is not negligible.  Our results also seem to contradic observations that show that more massive planets tend to be located farther away from the host star. This correlation is reproduced better by the calculations of \\citet{ida04a}, who conclude that maximal mass of the planets after including type II migration does not strongly depend on the size of the orbit. On the other hand, in the results from \\citet{alibert05} the correlation is identical to ours, namely the more massive planets tend to be closer to the star. This difference can have two sources. First, \\citet{ida04a} neglect the accretion of gas onto the planet, after it opens a gap in the disk. The second reason, which we think is more important, can again be the difference in the structure of the gaseous disk.  In our calculations ratio $(H/r)$ is mainly a decreasing function of the radius in the inner parts of the disk, while it is an increasing function in the disk model of \\citet{ida04a}. Because this ratio determines the mass of the planet which in turn is able to open a gap in the disk, then it can play a big role in determining the properties of the set of giant planets."
    },
    "0601/gr-qc0601078_arXiv.txt": {
        "abstract": "Brane-world models, where observers are restricted to a brane in a higher dimensional spacetime, offer a novel perspective on cosmology. I discuss some approaches to cosmology in extra dimensions and some interesting aspects of gravity and cosmology in brane-world models. ",
        "introduction": "A century after Einstein first proposed his theory of relativity, it has become a cornerstone of the physical sciences. Four dimensional spacetime provides the setting for describing physical processes and in particular provides the dynamical framework for cosmological models of our expanding Universe.  It was the general theory of relativity, proposed by Einstein in 1915, that for the first time provided equations with which to describe the dynamics of spacetime. Einstein's equation \\begin{equation} \\label{EE} G_{AB} + \\Lambda g_{AB} = 8\\pi G_N T_{AB} \\,, \\end{equation} relates the intrinsic curvature, $G_{AB}$, of the metric, $g_{AB}$, to the local energy-momentum, $T_{AB}$, while allowing for the possibility of a non-zero cosmological constant, $\\Lambda$. Consistency with Newtonian gravity in the weak-field, slow-motion limit is ensured by the appearance of Newton's constant, $G_N$, in the constant of proportionality.  In much of modern cosmology Einstein's tensor equation (\\ref{EE}) conveniently reduces to the Friedmann constraint equation \\begin{equation} \\label{F} 3 \\left( H^2 + \\frac{K}{a^2} \\right) = \\Lambda + 8\\pi G_N \\rho \\,, \\end{equation} which relates the Hubble expansion, $H$, and spatial curvature $K/a^2$, of a homogeneous and isotropic Friedmann-Robertson-Walker (FRW) spacetime to the local energy density $\\rho$.  Homogeneous and isotropic expansion has been used to build up a remarkably successful model for the evolution of our Universe starting with a hot Big Bang at a finite time in our past. This model has been tested not only by qualitative features such as the evolution of galaxy populations and the existence of a cosmic microwave background (CMB) radiation, but also quantitatively tested by comparing models of primordial nucleosynthesis with abundance of light elements.  One only needs to consider linear perturbations about a homogeneous and isotropic metric to build up a coherent picture of the formation of structure in our Universe. Small fluctuations, about one part in a hundred thousand, are observed in the temperature of the microwave background radiation and indicate the existence of small primordial perturbations in the distribution of matter and radiation in the early universe when the CMB last scattered, about 300,000 years after the Big Bang. These primordial density fluctuations provide the seeds around which the observed large-scale structure of our Universe can form simply by gravitational instability, in a cosmological model with appropriate contributions from radiation, baryonic matter, as well as cold dark matter and some form of dark energy, that behaves very much like Einstein's cosmological constant today. A wealth of observational data now enables cosmologists to put this basic picture to the test and attempt to measure parameters such as the density of different forms of matter, the nature of the primordial perturbations, and Einstein's gravitational laws.  At the same time fundamental questions remain unanswered. Why are there 3 large spatial dimensions (not 5 or 15)? why is the value of the cosmological constant so small? and what really happens at the initial Big Bang which represents a singular point at the start of our cosmological evolution? I cannot answer these questions in this talk, but I can show how brane-world models offer some novel and interesting perspectives on these issues. I should emphasize that this is a personal view and not intended to be a systematic review of all aspects of brane-world cosmology. For a more comprehensive review see~\\cite{Roy}. ",
        "conclusions": "In summary, the brane-world has offered a novel perspective on many of the unsolved problems in cosmology including the number of spatial dimensions, the initial singularity and the cosmological constant problem.  There remain many unsolved problems even in the simplest case of a co-dimension one brane. There is no known solution for an isolated black hole on the brane. Indeed it has been conjectured that there is no static black hole solution on the brane \\cite{Emparan}. Only a few special cases are known for anisotropic brane cosmologies \\cite{Fabbri}. Even linear perturbations about FRW branes are restricted to special cases or numerical solutions as the wave equation in the bulk is not in general separable in a Gaussian normal coordinate system.  Some of these problems are techinical difficulties simply due to the extra complication of an inhomogeneous extra dimension, but some problems are also more fundamental. If we are seeking to describe a single brane in an infinite extra dimension that the brane alone does not form a closed system and we must specify initial data in the bulk to determine the subsequent evolution.  A century after Einstein introduced physicists to four-dimensional spacetime, brane-worlds offer a big new higher-dimensional playground for relativists to explore. Cosmological models provide one area to study. If we can understand more about gravity in higher-dimensions then cosmology should also provide an opportunity to test these exotic ideas against observational data.    \\begin{theacknowledgments} I am grateful to the organisers for their hospitality in Oviedo and I thank my collaborators for the many collaborations upon which the work described in this article is based. This work is supported in part by PPARC grant PPA/G/S/2002/00576. \\end{theacknowledgments}"
    },
    "0601/astro-ph0601665_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:intro}  Surface brightness profiles of elliptical galaxies are normally well described by either a de~Vaucouleurs law \\citep{devaucouleurs48} or a S\\'{e}rsic law \\citep{sersic68}. However, the situation can be more complicated when we consider a galaxy embedded in a cluster. In this case, the outer parts of the profile may be modified due to tidal interactions with neighbouring galaxies, giving a `truncated' profile. A different case is often encountered in elliptical galaxies that are the brightest cluster member (BCM) and are located at the cluster centre, many of which have an extended envelope or halo \\citep{schombert87,gonzalez05}. These galaxies are known as `central dominant' (cD) galaxies. The extended envelope is observed as a deviation of the profile from a de~Vaucouleurs or S\\'{e}rsic law at large radii, and appears to correspond to light from stars orbiting in the cluster potential and not bound to the central galaxy \\citep*{bahcall77, gonzalez05}. Simulations indicate that the stars in the envelope may, on average, be older than the stars in the galaxy \\citep{murante04}.  NGC\\,4696 is the BCM of the Centaurus cluster. It is an elliptical galaxy located near the centre of the cluster. It is sometimes described as a cD galaxy, although no evidence has been found up to now of an extended halo.  Previous measurements of the light profile of NGC\\,4696 at optical wavelengths include those of \\citet{schombert87} and \\citet{jerjen97}, both obtained from photographic plates. \\citet{jerjen97} measured the light profile out to a radius of $\\approx$80\\,arcsec ($\\approx16\\,h^{-1}_{70}\\,\\mathrm{kpc}$), while the profile obtained by \\citet{schombert87} reaches much further, out to a radius of $\\approx$9\\,arcmin ($\\approx110\\,h^{-1}_{70}\\,\\mathrm{kpc}$). Neither observation showed signs of an extended envelope. The profile from \\citet{schombert87} followed an $r^{\\frac{1}{4}}$ law (de~Vaucouleurs profile) out to a radius of $\\approx$5\\,arcmin, and a truncation or `distention' of the profile was found for larger radii.  Measurements of the NIR profile were obtained by the Two Micron All Sky Survey (2MASS source 2MASX~J12484927-4118399 in the All-Sky Extended Source Catalog; http://www.ipac.caltech.edu/2mass/). Light profiles were obtained in three bands (\\textit{J}, \\textit{H} and \\textit{K}) for radii up to $\\approx$3\\,arcmin. This corresponds to surface brightnesses $\\mu_{H} \\leq 24$\\,mag\\,arcsec$^{-2}$ in the \\textit{H} band.  The Intra-Cluster Medium (ICM) of the Centaurus cluster has been studied using X-ray emission by the \\textit{ROSAT} and \\textit{ASCA} satellites. \\citet{allen94}, using data from \\textit{ROSAT}, found a correlation between the structure---ellipticity and centroid---of the X-ray emission and optical observations \\citep{sparks89} of the central galaxy NGC\\,4696. \\citet{ikebe99} used data from both satellites, and described the X-ray emitting gas in the central region ($r \\leq 3\\,\\mathrm{arcmin}$) according to a two-phase model. They identified the cool phase with the inter-stellar medium (ISM) of NGC\\,4696, and the hot phase with the ICM of the Centaurus cluster. These relationships between the central galaxy and the ICM could be an indication of the presence of an extended halo.  The previous optical and NIR work pose a quandary in that the form of the profiles measured range from de Vaucouleurs to exponential.  Thus the aim of the present work was to observe the light profile of NGC\\,4696 in the near-infrared (NIR) out to large radii, study the functional form of the profile (de Vaucouleurs or S\\'{e}rsic law), measure the parameters of the profile, and seek evidence for either a truncation due to tidal interactions or an extended envelope following the cluster potential.  The main advantage of the NIR  over optical studies is that the light is dominated by old stars, and thus the NIR profile will better trace the bulk of the stellar mass of the galaxy.  Furthermore the extinction due to dust is smaller in the NIR than at optical wavelengths.  Another motivation for this work is the application to studying radial colour distributions, through coupling with optical data, to understand the distribution of stellar populations throughout BCMs.  Through comparisons with simulations (e.g.\\ \\citealt{som05}), such studies would help to constrain BCM and cluster formation theories and the enrichment of the intra-cluster medium (e.g.\\ \\citealt{lin04}).   We describe our observational data in Section~\\ref{sec:observ} and explain the reductions and calibrations in Section~\\ref{sec:red} and the method used to fit the profile models in Section~\\ref{sec:an}. Section~\\ref{sec:results} presents the results and Section~\\ref{sec:concl} summarises our conclusions.  Throughout the paper we use a Hubble constant of $H_{0}=70\\,h_{70}$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$. We have used a redshift of $z=0.00987\\pm0.00005$ (\\citealt{devaucouleurs91}, which assumes no peculiar velocity) which equates to a distance of $42.3\\pm0.2\\,h_{70}^{-1}$\\,Mpc; the angular scale at this distance is $0.205\\pm0.001$\\,kpc\\,arcsec$^{-1}$.  Distances derived from surface brightness fluctuations, taking into account the peculiar motions of the Cen30 and Cen45 components of the Centaurus cluster, and the Hydra cluster, give a distance of $42.5 \\pm 3.2$ Mpc, consistent with the value used here (\\citealt{mie05}). ",
        "conclusions": ""
    },
    "0601/astro-ph0601386_arXiv.txt": {
        "abstract": "Using the NIRSPEC spectrograph at Keck II, we have obtained $H$ and $K$-band echelle spectra for a young ($\\sim10-15$ Myr), luminous ($M_V\\sim-13.2$) super-star cluster in the nearby spiral galaxy NGC~6946.  From spectral synthesis and equivalent width measurements we obtain for the first time accurate abundances and abundance patterns in an extragalactic super-star cluster.  We find [Fe/H]$=-0.45\\pm0.08$~dex, an average $\\alpha$-enhancement of $\\approx+0.22\\pm0.1$~dex, and a relatively low  $\\rm ^{12}C/^{13}C\\approx 8\\pm2$ isotopic ratio. We also measure a velocity dispersion of $\\approx$9.1~km/s, in agreement with previous estimates. We conclude that integrated high-dispersion spectroscopy of massive star clusters is a promising alternative to other methods for abundance analysis in extragalactic young stellar populations. ",
        "introduction": "Star clusters have a long history as important tools for addressing a wide range of questions in astronomy. With few exceptions, they are ``simple stellar populations'' (SSPs), i.e.\\ they are composed of stars born in a single burst and sharing the same chemical composition (at least to first order) and age. They have played an important role as test labs for models of stellar evolution \\citep[e.g.][]{mm81,ren88,chi92}. As confidence has grown in our ability to model their integrated properties, so has their importance as tracers of stellar populations in galaxies that are too distant for individual stars to be resolved.  The latter point is perhaps best illustrated by considering, for the moment, the old \\emph{globular clusters} (GCs) which are ubiquitous in all major galaxies.  GCs typically contain large numbers of stars ($10^4 - 10^6$), so most phases of stellar evolution are well sampled and stochastic effects therefore minimized. Although there are still unsolved problems (e.g.\\ concerning horizontal branch morphology), SSP models can now provide a reasonably realistic account of integrated GC observables such as broad-band colours and absorption line strengths as a function of age and metallicity.  Thanks to the availability of efficient spectrographs on 8-10 m class telescopes, spectroscopy of GCs in galaxies well beyond the Local Group is now routinely feasible, and has been utilized in many studies to probe the stellar populations in early-type galaxies and constrain their ages and metallicities \\citep[e.g.][and references therein]{bs06}. Much of this work has relied on measurements of absorption line features at relatively low ($\\sim10$ \\AA) spectral resolution. In the optical, the Lick/IDS system of absorption line indices has found wide-spread use \\citep{bur84,tra98}, while indices for the near-infrared have been defined by \\citet{ba87} and \\citet{vaz03}.  Ages are typically estimated from Balmer line strengths (mainly H$\\delta$, H$\\gamma$, and H$\\beta$), with other indices being more sensitive to metallicity. However, detailed abundances of individual elements typically cannot be reliably measured at this resolution. In addition, the Lick/IDS system is not well tailored for studying young stellar populations, partly due to limitations in the empirical libraries used in the construction of SSP models, but also because the index definitions are designed primarily for studies of old stellar populations.  Until now, the main source of information about the chemical composition of \\emph{young} stellar populations beyond the Local Group has been measurements of emission lines in H{\\sc ii} regions. However, these provide access to only a limited set of elements (mainly O, N, S) and often have to rely on empirical calibrations of line strengths vs.\\ metallicity. The auroral lines (e.g.\\ [O{\\sc iii}] 4363\\AA) are usually too faint to be measured directly, thus prohibiting a determination of the nebular temperature \\citep{sta01}. Furthermore, H{\\sc ii} regions only offer a snapshot of the \\emph{present-day} chemistry and do not provide any information about the past history.  Spectroscopy of (massive) star clusters has the potential to provide information on the entire star formation histories of galaxies.  For masses in the $10^5$ M$_\\odot$ -- $10^6$ M$_\\odot$ range, star clusters typically have velocity dispersions of about 5--10 km/s, allowing studies of their integrated properties at spectral resolutions up to $\\lambda/\\Delta\\lambda \\approx$ 30,000 or more.  This is sufficient to constrain individual element abundances. However, extending this type of analysis to young clusters comes with its own set of difficulties.  Compared to studies of old GCs, some of the main challenges are: 1) models for the massive stars found in clusters with ages of a few $\\times10^7$ years are much less certain than those of the low-mass stars found in old GCs \\citep{mas03} 2) the optical spectra are generally dominated by the relatively featureless spectra of hot stars, diluting absorption features and making them harder to measure, 3) for all but the most massive clusters, the integrated spectra may be dominated by a few massive stars, causing unpredictable stochastic fluctuations in integrated properties \\citep[e.g.][]{lm00}.  In this letter we aim to take a first step towards deriving abundances for \\emph{young} extragalactic star clusters from their integrated light, by modelling the near-infrared spectrum of a luminous ($M_V=-13.2$) young ($\\sim10$ Myr) star cluster in the nearby spiral galaxy NGC~6946. This object was first identified as a star cluster by \\citet{lr99} and was labelled NGC6946-1447 by them, but is located within a larger stellar complex first noted by \\citet{hod67}.  Both $UBVI$ broad-band colours and spectroscopic observations of Balmer and He I absorption lines yield consistent age estimates in the range 10--15 Myr \\citep{laretal01,efre02}. From the luminosity and estimated age of the cluster, SSP models yielded a mass of $\\sim0.8\\times10^6$ M$_\\odot$ for a Salpeter-like IMF extending down to 0.1 M$_\\odot$, or $\\sim0.55\\times10^6$ M$_\\odot$ if the IMF is log-normal below 0.4 M$_\\odot$ \\citep{laretal01}.  A dynamical mass estimate, based on high-dispersion optical spectroscopy and HST imaging, yielded a somewhat higher mass of $\\sim1.7\\times10^6$ M$_\\odot$. Such massive clusters are sometimes called ``super star clusters'' (SSCs). Using isochrones from \\citet{gir00}, we estimate that the cluster contains about 130 red supergiants. A detailed description of the cluster and surrounding stellar complex is given in \\citet{laretal01,lar02}.  No line emission is observed from the cluster itself but \\citet{efre02} estimated an oxygen abundance of $12 + \\log(O/H) = 8.95 \\pm 0.2$ from long-slit spectroscopy of nearby H{\\sc ii} regions. \\citet{br92} measured O abundances for H{\\sc ii} regions distributed throughout the disk of NGC~6946 (using narrow-band imaging of O, N and H emission lines) and derived an oxygen abundance gradient of $\\Delta \\log (O/H) / \\Delta R = -0.089 \\pm 0.003 \\, {\\rm dex}\\,  {\\rm kpc}^{-1}$ and a central value of $12 + \\log (O/H)$ = 9.37. At the projected galactocentric distance of NGC6946-1447 (4.8 kpc), this corresponds to $12 + \\log (O/H) = 8.94 \\pm 0.01$.  \\citet{kob98} quote an O abundance of $12 + \\log (O/H) = 9.13$ at a radius of 3 kpc, or $12 + \\log (O/H) = 8.97$ at 4.8 kpc if we adopt the abundance gradient from \\citet{br92}. Thus, all available measurements consistently give $12+\\log(O/H)$ between 8.94 and 8.97, although the above O abundances are all based on measurements of the collisionally excited O lines vs.\\ Balmer line ratios and may therefore be subject to systematic uncertainties at the 0.1-0.2 dex level. ",
        "conclusions": "Based on near-infrared $H$ and $K$-band spectroscopy, we have carried out a detailed abundance analysis of a young massive star cluster in the nearby spiral galaxy NGC~6946. In this exploratory work, we have derived abundances of several individual key elements, including Fe, Ca, Si, Mg and Ti. We find a sub-solar Fe abundance ([Fe/H] = $-0.45\\pm0.08$), while the $\\alpha$-element to Fe abundance ratios are all enhanced with respect to the Solar values with a mean [$\\alpha$/Fe] = $0.22\\pm0.11$. The O abundance derived from the cluster spectrum is [O/H]$ = -0.17\\pm0.09$, about 0.3 dex lower than the value based on H{\\sc ii} regions at the same galactocentric distance. We find a $^{12}$C/$^{13}$C ratio of $\\approx8\\pm2$, similar to or slightly lower than typically observed in red supergiants in Galactic open clusters \\citep{luck94,gw00} and in the (more metal-poor) cluster NGC~330 in the SMC \\citep{gw99}.  Standard stellar models \\citep{schal92} predict a surface $^{12}$C/$^{13}$C ratio of $\\sim17$ for a 15 M$_\\odot$, $\\sim12$ Myr-old (Solar metallicity) red supergiant, about a factor of two higher than the value derived here.  Super-solar [$\\alpha$/Fe] ratios are usually interpreted as signatures of rapid, bursty star formation, with the gas mainly enriched by Type II supernovae with short-lived, massive progenitor stars \\citep{mac97}. Stellar populations formed over timescales of several Gyrs or more are expected to show solar-like $\\alpha$/Fe-element abundance ratios, as observed in the Milky Way thin disk.  In this context, it is worth noting that the complex hosting NGC6946-1447 may qualify as a `localized starburst' \\citep{efre04}.  Timing is a critical issue, however: in order to produce super-solar [$\\alpha$/Fe] ratios, the starburst must have \\emph{preceded} the formation of the cluster itself by several Myrs. The delay must have been long enough for Type II SNe to produce significant amounts of $\\alpha$-elements, and enough time must then have elapsed to allow mixing of the pre-existing gas with $\\alpha$-enhanced ejecta before the cluster formed. Interestingly, reconstruction of the field star colour-magnitude diagram has provided some evidence for star formation in the complex at least 10--15 Myr prior to the formation of the cluster \\citep{lar02}.  Alternatively, the current global SFR in NGC~6946 may be sufficiently elevated above the past average to cause a general enrichment of the ISM with $\\alpha$-elements. Addressing these questions more quantitatively would require a detailed modeling of the chemical enrichment history and a knowledge of the past star formation history which is currently not available.  This study represents one of the first cases of a detailed abundance analysis of a star-forming galaxy beyond the Local Group.  While many uncertainties remain, we suggest that observations of extragalactic young star clusters hold great potential for constraining the chemical enrichment histories of their parent galaxies."
    },
    "0601/astro-ph0601179_arXiv.txt": {
        "abstract": "We have tested for luminosity, colour and morphology dependence of the degree of filamentarity in seven nearly two dimensional strips from the Sloan Digital Sky Survey Data Release Four (SDSS DR4). The analysis is carried out at various  levels of coarse graining allowing us to address different length-scales. We find that the brighter galaxies have a less  filamentary distribution than the fainter ones at all levels of coarse graining. The distribution of red galaxies and ellipticals shows  a higher degree  of filamentarity compared to blue galaxies and spirals respectively at low levels of coarse graining. The behaviour is reversed at higher levels of coarse graining. We propose a  picture where the ellipticals are densely distributed in the vicinity of the nodes where the filaments intersect while the spirals are sparsely distributed along the entire extent of the filaments. Our findings indicate that the regions with an excess of ellipticals  are larger than galaxy clusters, protruding  into the filaments. We have also compared the predictions of a semi-analytic model of galaxy formation (the Millennium Run galaxy catalogue) against our results for the SDSS. We find  the two to be in agreement for the $M^{*}$ galaxies and for the red galaxies, while the model fails to correctly predict the filamentarity of the brighter galaxies and the blue galaxies. ",
        "introduction": "One of the most significant aim of all large redshift surveys is to determine the spatial distribution of galaxies in the Universe. Over the last few decades redshift surveys have revealed the large scale structures in the Universe in it's full glory (e.g. , CfA , \\citealt{gel}; LCRS, \\citealt{shect}; 2dFGRS, \\citealt{colles} and SDSS, \\citealt{stout}, \\citealt{abaz1}, \\citealt{abaz2}). It is found that the vast majority of galaxies preferentially reside in an intricate network of interconnected filaments and walls surrounded by voids. The filaments have a dimension of many tens of {\\rm Mpc} in length (\\citealt{bharad2}) and 1-2 $\\,{\\rm Mpc}$ in thickness (\\citealt{rat}, \\citealt{doro2}). In fact the Local Group in which the Milky Way resides, is believed to lie along an extended filament that originates near the Ursa Major and includes both the IC342 and M81 groups and connects with another primary filament which originates from the Virgo cluster and includes the Sculpture group (\\citealt{peeb}).  The complex filamentary network of galaxies, the so called 'Cosmic web' is possibly the most striking visible feature in the large scale structure of the Universe. Quantifying the Cosmic web and tracing it's origin has remained one of the central issues in cosmology.  The analysis of filamentary patterns in the galaxy distribution has a long history dating back to a few papers in the mid-eighties by \\citet{zel}, \\citet{shand1} and \\citet{einas1}. More recently, extensive studies of the LCRS reveal a rich network of interconnected filaments surrounding voids (eg. \\citealt{shand2} ; \\citealt{bharad1} ; \\citealt{mul} ; \\citealt{doro1} ; \\citealt{trac} ; \\citealt{doro2} ; \\citealt{bharad2}). \\citet{bharad3} have compared the filamentarity in the LCRS galaxy distribution with $\\Lambda$CDM dark matter N-body simulations to  show that the two are consistent provided  a mild galaxy bias is included. \\citet{sheth2} has used a technique SURFGEN (\\citealt{sheth1}) to study the geometry, topology and morphology of the superclusters in mock SDSS catalogues. \\citet{basil} and \\citet{koloko} have studied the super-cluster void network in the PSCz and the Abell/ACO cluster catalogue respectively, finding filamentarity to be the dominant feature. \\citet{pimb} have studied the intercluster galaxy filaments in the 2dFGRS.  The Sloan Digital Sky Survey (SDSS) \\citep{york} is currently the largest galaxy redshift survey. In a recent paper \\citep{pandey} we have analyzed the filamentarity in the two equatorial strips of the First Data Release of the SDSS (SDSS DR1).  These strips are nearly two dimensional (2-D). We have projected the data onto a plane and analyzed the resulting 2-D galaxy distribution.  The large volume and dense sampling of the SDSS allows us to construct volume limited subsamples extending over lengthscales which are substantially larger than possible with earlier surveys.  We find evidence for connectivity and filamentarity in excess of that of a random point distribution, indicating the existence of an interconnected network of filaments. The filamentarity is found to be statistically significant up to length-scales $80 \\, h^{-1} {\\rm Mpc}$ and not beyond \\citep{pandey}. Further, we show that the degree of filamentarity exhibits a luminosity dependence with the brighter galaxies having a more concentrated and less filamentary distribution.  It is now quite well accepted that galaxies with different physical properties are differently distributed in space.  Studies over several decades have established that ellipticals and spirals are not distributed in the same way. The ellipticals are found predominantly inside regular, rich clusters whereas the field galaxies are mostly spirals. This effect is referred to as ``morphological segregation'' which has been studied extensively in the literature (eg. \\citealt{hubble}, \\citealt{zwicky}, \\citealt{davis2}, \\citealt{dress} , \\citealt{guzo}, \\citealt{zevi}, \\citealt{goto}). Further, ellipticals exhibit a stronger clustering as compared to spirals (\\citealt{zevi}). The analysis of the two-point correlation function of galaxies with different colours shows the red galaxies, which are mainly ellipticals with old stellar populations, to have a stronger clustering as compared to the blue galaxies which are mostly spirals (e.g. \\citealt{will}, \\citealt{brown}, \\citealt{zehavi}). Studies of the topology using the genus statistics show that red galaxies exhibit a shift towards a meatball topology (\\citealt{hoyle1}, SDSS EDR, \\citealt{park1}, SDSS DR3) implying that they prefer to inhabit the high density environments. It is found that luminous galaxies exhibit a stronger clustering than their fainter counterparts (e.g. \\citealt{hamil}, \\citealt{davis1}, \\citealt{white}, \\citealt{park}, \\citealt{love}, \\citealt{guzo}, \\citealt{beno}, \\citealt{nor}, \\citealt{zehavi}). The difference increase markedly above the characteristic magnitude $M*$ (\\citealt{nor}, 2dFGRS, \\citealt{zehavi}, SDSS DR2) of the Schecter luminosity function (\\citealt{sek}). The detailed luminosity dependence has been difficult to establish because of the limited dynamic range of even the largest redshift surveys.  \\citet{einas2} have studied the luminosity distribution of galaxies in high and low density regions of the SDSS to show that brighter galaxies are preferentially distributed in the high density environments. \\citet{goto} have studied the morphology density relation in the SDSS(EDR) and found that this relation is less noticeable in the sparsest regions indicating the need for a denser environment for fostering the physical mechanisms responsible for the galaxy morphological changes. \\citet{hog1} and \\citet{blan1} find a strong environment dependence for both the colour and luminosity for the SDSS galaxies.  The dependence of clustering on galaxy properties like the luminosity, colour and morphology provides very important inputs for theories of galaxy formation. It is a natural prediction of hierarchical structure formation that the rarer objects which correspond to peaks in the density field should be more strongly clustered than the population itself (\\citealt{kais}, \\citealt{davis3}, \\citealt{white1}). A possible interpretation of the observations is that the more luminous galaxies are hosted in higher density peaks as compared to the fainter galaxies. In this scenario the properties of a galaxy are largely decided by the initial conditions at the location where the galaxy is formed. This idea is implemented in the halo model (\\citealt{jingmo}, \\citealt{benson}, \\citealt{seljak}, \\citealt{pkok}, \\citealt{ma}, \\citealt{scocci2}, \\citealt{white1}, \\citealt{berlind}, \\citealt{scran}, \\citealt{yang1}, \\citealt{yang2}) where it is postulated that all galaxies lie in virialised halos and the number and type of galaxies in a halo is entirely determined by the halo mass. \\citet{cooray} provide a detailed review of the halo model. It is well known that galaxy-galaxy interactions and environmental effects are also important in determining the physical properties of a galaxy. For example, ram pressure stripping in galaxy clusters is a possible mechanism for the transformation of a spiral galaxy into an elliptical. Galaxy-galaxy interactions could also provide a mechanism for morphological transformation.  In the classical picture galaxies evolve in isolation and morphology is determined at birth, while in the hierarchical model galaxies have no fixed morphology and they can evolve depending on the nature of their mergers. A study of the spatial distribution of different kind of galaxies holds the potential to probe the factors responsible in determining galaxy luminosity, colour and morphology.  Much of the research on the luminosity, morphology and colour dependence of the galaxy distribution has focused either on the morphology segregation in clusters or has studied the behaviour of the two point correlation as a function of these galaxy properties. Filaments are the largest known statistically significant coherent features. They have been shown to be statistically significant to scales as large as $80 \\, h^{-1} \\, {\\rm Mpc}$. In this paper we  study  the luminosity, colour and morphology dependence of filaments  in the Sloan Digital Sky Survey Data Release Four (SDSS DR4). This allows us to study the distribution of different types of galaxies on the largest length-scales possible. Such a study will enable us to estimate  the length-scale over which the factors responsible for the galaxy segregation operate. While the two-point correlation function completely characterizes the statistical properties of a Gaussian random field, it is well known that the galaxy distribution is significantly non-Gaussian. In fact the presence of large-scale coherent features like filaments is a clear indication of non-Gaussianity. There have been studies of how the three point correlation function depends on galaxy properties \\citep{kayo,jingbo, gazta}. While the first two papers do not find any statistically significant effects, \\citet{gazta} find statistically significant colour and luminosity dependence. The present work studies how the largest known non-Gaussian features, the filaments, depend on galaxy properties.  This paper presents progress on many counts compared to our earlier work \\citep{pandey} based on the SDSS DR1. The earlier work was restricted to single volume limited subsamples from each of the two equatorial strips. The absolute magnitude range was divided into two equal halves, and these were analyzed to test for a luminosity dependence in the filamentarity. For a statistically significant conclusion it is necessary to also estimate the cosmic variance (sample to sample variation) of the filamentarity which was done by bootstrap resampling. We found that the difference in filamentarity between the two magnitude bins was in excess of the fluctuations expected from cosmic variance, thereby establishing the statistical significance.  The SDSS DR4 covers a large contiguous region in the Northern Galactic Cap with a few gaps still to be covered.  The SDSS has been carried out in overlapping strips of $5^{\\circ}$ width oriented along great circles. Since the detection of coherent large-scale features requires a completely sampled region with no gaps, we have extracted seven non-overlapping strips of width $2^{\\circ}$ each. The volume corresponding to each strip is nearly two dimensional and we have collapsed the thickness and analysed the resulting 2-D galaxy distribution. The increased number of strips provide a better estimate of the cosmic variance leading to more robust conclusions. The analysis has been carried out on five overlapping bins in absolute magnitude. We have extracted separate volume limited subsamples corresponding to each magnitude bin and compared the filamentarity across them.  The earlier analysis has been extended here to also consider the colour and morphology dependence, which has been tested in each of the volume limited subsamples.  Here we would like to mention that the statistics that we are using to quantify the filamentarity is not absolute in the sense that it is sensitive to the volume and galaxy number density of the sample. It is thus only meaningful to compare different galaxy samples provided they have the same volume (identical shape and size) and number density. This is an important consideration when constructing the galaxy samples for which we test the luminosity, colour and morphology dependence.  A brief outline of our paper follows. Section 2 describes the data and Section 3 the method of analysis. Our results for the SDSS data are presented in Section 4, in Section 5 we compare our results with the predictions of a semi analytic model of galaxy formation and finally we present our conclusions in Section 6.  It may be noted that  we have used a $\\Lambda$CDM cosmological model with $\\Omega_{m0}=0.3$,  $\\Omega_{\\Lambda0}=0.7$ and $h=1$ throughout. ",
        "conclusions": "Filaments are the largest known statistically significant coherent features visible in the galaxy distribution. We test if the degree of filamentarity depends on the galaxy properties.   We find evidence for statistically significant luminosity, colour and morphology dependence in the average filamentarity $F_2$ studied as a function of the filling factor $FF$.  Comparing the average filamentarity $F_2$ of different luminosity bins (Figure \\ref{fig:7}), we find that the fainter galaxies have a more filamentary distribution as compared to the brighter ones.  These differences exist at nearly all values of $FF$.  The drop in filamentarity with increasing luminosity is found to be particularly more pronounced in the two brightest luminosity bins which we have considered, both of which are brighter than $M^*$.  Comparing the thickness and length (Figure \\ref{fig:9}), we find that at a fixed value of $FF$ structures traced by the bright galaxies are thicker and shorter than those traced by the faint ones. While the differences in the length are statistically significant only at large $FF$, the differences in the thickness are statistically significant at all values of the filling factor. With increasing coarse-graining there is a percolation transition at $FF \\sim 0.5$ beyond which there is an interconnected network of filaments encircling voids. After the onset of percolation the length per void gives an estimate of the circumference of the filaments encircling the voids and dividing this by $\\pi$ gives an estimate of the average void diameter. We find that the faint galaxy distribution is more porous (larger number of holes or voids) and has smaller voids with a typical void diameter $\\sim 80 \\pm 62 \\, h^{-1} \\, {\\rm Mpc}$ as compared to the bright galaxy distribution which has a typical void diameter $\\sim 165 \\pm 140 \\, h^{-1} \\, {\\rm Mpc}$.  Note that these estimates of the void diameter could be somewhat flawed because the estimated void diameters exceed the radial extent ($96 \\, h^{-1}\\, {\\rm Mpc}$) of the samples which we have used to test luminosity dependence (Figure \\ref{fig:2}), but they serve the purpose of demonstrating the luminosity dependence of the void size. Our findings indicate that galaxies of different luminosity are not uniformly distributed along the cosmic web. The brighter galaxies are preferentially distributed in more compact and thicker regions with large voids in between whereas the fainter galaxies inhabit thin, elongated regions with more numerous voids of smaller diameter.  We note that studies using the projected two-point correlation function (eg. \\citealt{nor1}; \\citealt{zehavi}) show an increase in correlation amplitude with increasing luminosity, the effect being pronounced beyond $L^*$. These earlier results are consistent with our findings but they do not have any information about the shapes of the clustering patterns. Recent studies (eg. \\citealt{coil}; \\citealt{polo}) also show evidence for luminosity dependence at higher redshifts ($z \\simeq 1$).  The results for the colour and morphology dependence of the average filamentarity are somewhat different from those for the luminosity dependence. We find that the red galaxies and the ellipticals have a higher filamentarity as compared to the blue galaxies and spirals respectively at low filling factors (Figure \\ref{fig:9}). There is a cross-over at $FF \\sim 0.15$ for the colour dependence and $FF \\sim 0.25$ for the morphology dependence. The blue galaxies and the spirals have a higher filamentarity as compared to the red galaxies and the ellipticals respectively at large filling factors. Considering the length and the thickness (Figure \\ref{fig:9}), we find that the length shows a statistically significant colour and morphology dependence only at large $FF$ (beyond percolation) where the distribution of blue galaxies and the spirals has a greater length as compared to red galaxies and ellipticals respectively. The red galaxies and ellipticals have a thicker distribution, the differences being statistically significant at all $FF$. Considering the voids, we find that the distribution of red galaxies has fewer voids which are larger (diameter $\\sim 127 \\pm 46 \\, h^{-1} \\, {\\rm Mpc}$) whereas the distribution of blue galaxies has more voids which are smaller (diameter $\\sim 98 \\pm 17 \\, h^{-1}\\, {\\rm Mpc}$). The ellipticals and spirals show similar differences with void diameters $\\sim 115 \\pm 50 \\, h^{-1}\\, {\\rm Mpc}$ and $\\sim 90 \\pm 21 \\, h^{-1}\\, {\\rm Mpc}$ respectively. These estimates of the void diameter are quite reliable as the sample size (bin 4) is larger than the void diameter (Figure \\ref{fig:3} and \\ref{fig:4}).  We note that our estimates are somewhat larger compared to results using other methods. For example, \\citet{hag} have found that the voids have a scale of around  $40 \\, h^{-1}\\, {\\rm Mpc}$ in IRAS 1.2 Jy redshift survey.  A similar conclusion is reached by \\citet{ul} from their analysis of the Northern Local Void. A recent analysis of voids in the 2dFGRS by \\citet{hoyle} find that voids have typical diameters of $\\sim 30 \\, h^{-1}\\, {\\rm Mpc}$. Most of these  estimates  use the area or volume to determine  the void diameter whereas we use the void perimeter. Substructures in the void perimeter is possibly the dominant reason why we get a larger estimate of the diameter. Further, the large variance in the void diameter seems to indicate that they have a large spread in sizes.  The colour and morphology are very strongly correlated galaxy properties, the red galaxies being predominantly ellipticals and the blue ones spirals. It is well known that elliptical are found primarily in dense groups and clusters \\citep{dress} whereas the spirals are distributed in the field. Our finding that at large length-scales the structures traced by the ellipticals are thicker, and less filamentary as compared to the spirals is consistent with a picture where the entire galaxy population is distributed along an interconnected network of filaments encircling voids, the Cosmic Web. The ellipticals preferentially inhabit the dense groups and clusters which can be identified with the nodes of the Cosmic Web, the places where filaments intersect. The spirals, on the other hand, are sparsely distributed along the filaments. The colour and morphology are also known to be correlated with the luminosity. The more luminous galaxies are predominantly ellitpicals and the fainter ones spirals.  The fact that at small filling factors (which we may associate with small length-scales) the eliipticals have a more filamentary distribution as compared to spirals, whereas the opposite is found at large length-scales is quite intriguing. We discuss below a possible interpretation of this finding.  The higher filamentarity of the ellipticals at low $FF$ seem to indicate that in addition to residing in the clusters at the intersection of filaments, the ellipticals also extend a little along the filaments which originate from the clusters.  The spirals on the other hand, are distributed along the entire extent of the filaments.  This is shown schematically in Figure \\ref{fig:10}.  In this picture the ellipticals have a more compact distribution, and as a consequence they connect up to form filamentary clusters at the initial stages of coarse graining. Since these structures are localized near the nodes of the filaments and they do not extend along the entire length of the filaments, their filamentarity increases slowly during the later stages of coarse graining. The spirals, on the other hand, are sparsely distributed along the entire length of the filaments. They connect up only at a later stage of coarse graining to define the entire filamentary network of the Cosmic Web.  This is a possible explanation why at low filling factors the ellipticals have a higher filamentarity than the spirals. Further, it also indicates that with increasing filling factor (later stages of coarse graining) the filamentarity of the spirals should grow faster than that of the ellipticals as seen in (Figure \\ref{fig:7}).  Finally we note that the qualitative picture presented here is a plausible model which is consistent with the quantitative findings of this paper. We do not claim that it is the only possible explanation, and it is presented here more in the spirit of a hypothesis rather than a conclusion.  Further work is required to establish or refute this picture.  As already mentioned several times, it should be noted, that the curves showing $F2$ as a function of $FF$ are not absolute. They depend on the geometry of the volume (both shape and size) and the galaxy number density. It is only meaningful to compare different  galaxy samples for which these  quantities are the same. Once these are fixed, we can attribute changes in the filamentarity to  factors like luminosity, colour or morphology which we are testing for.  \\begin{figure} \\rotatebox{0}{\\scalebox{.9}{\\includegraphics{interpret.ps}}} \\caption{This shows our picture for the relative distribution of ellipticals and spirals along the Cosmic Web. The ellipticals densely   inhabit the nodes and are distributed within  the encircled regions  shown in the figure. The spirals are sparsely distributed along the larger volume of the entire filaments. } \\label{fig:10} \\end{figure}   Observational findings indicate a close connection between the local density and galaxy type \\citep{dress} with an increase in the  elliptical and S0 population and a  decrease in the spirals in galaxy clusters. Our findings are consistent with this, and it indicates that there is an excess of ellipticals on scales larger than the clusters extending to some extent into the filaments. In the Zel'dovich pancake scenario filaments form at the intersection of sheets and clusters form at the nodes of the filaments. Structure formation occurs in a hierarchical fashion with sheets forming first, matter flows along the sheets into the filaments which in turn drain into the clusters. The higher galaxy density and a denser hot gas environment are possibly the factors  leading to a higher elliptical fraction in the vicinity of the nodes.   The relation between our findings and the various theories for galaxy formation is an important issue. We have compared our findings against the Millennium Run \\citep{springel} which is one of the largest simulated galaxy catalogues that  incorporates the theory of galaxy formation in a semi-analytic fashion. The semi-analytic model has a number of parameters which have been tuned to match various observed properties of the local galaxy distribution like the luminosity functions and the Tully-Fisher relation. In this paper we have compared the filamentarity of the galaxy distribution in the semi-analytic model against the actual data from the SDSS DR4. Studying this in different luminosity bins we find that the two are consistent for  $L^{*}$ galaxies, while the brighter galaxies  show significant differences. For the brighter galaxies,  at large filling factors   the semi-analytic model predicts a substantially higher filamentarity as compared to the actual data, whereas it underpredicts the same at small $FF$. Though the semi-analytic model correctly predicts the number density of very luminous galaxies  as reflected in the luminosity function, it fails to correctly reproduce the geometry of their spatial  distribution. For these galaxies the model does not predict a sufficiently filamentary distribution at small length-scales whereas there is an excess filamentarity predicted at large scales.  Comparing the filamentarity predicted by the model against the actual data for galaxies of different colours , we find that the two are consistent for  red galaxies whereas the distribution of the blue galaxies are quite different. For these galaxies the model  markedly underpredicts the filamentarity at small length-scales or filling factors. We speculate that the discrepancy may be a consequence of two possible ingredients of the model, the first being a radio mode feedback from AGNs in massive halos introduced to stop cooling flows. This effectively quenches  star formation in the galaxies in these halos thereby transforming them into red galaxies. The  second possibility  is the prescription (plane parallel slab model) of \\citet{kofman5} to include the effects of dust when calculating the galaxy luminosities and colours. It may be noted  that  \\citet{springel} have  found discrepancies in the colour dependence of the two point correlation function. They find that the differences in the correlation amplitudes of the red and blue galaxies predicted by the model are in excess of that seen in the 2dFGRS and SDSS.  In conclusion we note that the filamentary pattern seen in the galaxy distribution exhibits statistically significant luminosity, colour and morphology dependence. We have speculated on a  geometrical picture which can explain some features of the colour and morphology dependence of the filamentarity. Explaining the origin of the observed luminosity, colour and morphology dependence in terms of the theories of galaxy formation is an important issue which needs to be addressed in the future."
    },
    "0601/astro-ph0601453_arXiv.txt": {
        "abstract": "We propose a new parametrization of the deceleration parameter to study its time-variation behavior. The advantage of parameterizing the deceleration parameter is that we do not need to assume any underlying theory of gravity. By fitting the model to the 157 gold sample supernova Ia data, we find strong evidence that the Universe is currently accelerating and it accelerated in the past. By fitting the model to the 115 nearby and Supernova Legacy Survey supernova Ia data, the evidence that the Universe is currently accelerating is weak, although there is still a strong evidence that the Universe once accelerated in the past. The results obtained from the 157 gold sample supernova Ia data and those from the 115 supernova Ia data are not directly comparable because the two different data sets measure the luminosity distance up to different redshifts.  We  then use the Friedmann equation and a dark energy parametrization to discuss the same problem. When we fit the model to the supernova Ia data alone, we find weak evidence that the Universe is accelerating and the current matter density is higher than that measured from other experiments. After we add the Sloan Digital Sky Survey data to constrain the dark energy model, we find that the behavior of the deceleration parameter is almost the same as that obtained from parameterizing the deceleration parameter. ",
        "introduction": "Astronomical observations suggest the existence of dark energy which has negative pressure and contributes about 2/3 of the critical density to the total matter density of the Universe within the framework of Einstein's general relativity \\cite{sp99,sdss,wmap}. While a wide variety of dynamical dark energy models were proposed in the literature \\cite{DE}, there are also model-independent studies on the nature of dark energy by using the observational data. In particular, one usually parameterizes dark energy density or the equation of state parameter $w(z)$ of dark energy \\cite{par1,par2,par3,par4,par5,gong05}. For the parametrization of $w(z)$, we need to determine $\\Omega_{m0}$ in addition to the parameters in $w(z)$. Furthermore, one needs to assume the validity of general relativity in all these studies.  Since supernova (SN) Ia data measures the luminosity distance redshift relationship, it provides a purely kinematic record of the expansion history of the Universe. It is possible to probe the evolution of the Hubble parameter or the deceleration parameter by using SN Ia data without assuming the nature and evolution of the dark energy \\cite{TS02,daly,riess}. For instance, by parameterizing the deceleration parameter $q(z) = q_{0} + q_{1} z$, Riess {\\em et al} \\cite{riess} were able to study the kinematics of the universe. However, it was soon realized that such a parametrization cannot re-produce the behavior of the cosmological constant \\cite{Virey05}. An alternative parametrization is a piecewise constant acceleration with two distinct epochs \\cite{TS02}, for which it was found that SN Ia data favors recent acceleration and past deceleration. Recently Shapiro and Turner (ST) applied several simple parametrization of the deceleration parameter $q(z)$ to the 157 SN Ia data \\cite{riess} to study the property of $q(z)$ \\cite{sturner}. They found that {\\em there is little or no evidence that the Universe is presently accelerating, and that  there is very strong evidence that the Universe once accelerated}. The advantage of parameterizing $q(z)$ is that the conclusion does not depend on any particular gravitational theory. The disadvantage is that it will not give much direct information on the cause of an accelerated Universe.  The conclusion arrived in \\cite{sturner} was based on a simple three epoch model of $q(z)$, in which the function $q(z)$ is not smooth. Since the current SN Ia data is still sparse, the division of the data to three different redshift bins may not be a good representation of the data. Therefore, the conclusion based on this technique needs to be further studied. Following ST, we propose a simple smooth function of $q(z)$ which we believe is more realistic and then apply the observational data to get the behavior of the deceleration parameter. After we are sure that the Universe experienced acceleration, we assume that the acceleration is due to the presence of dark energy and use a simple dark energy parametrization to study the property of dark energy in the framework of general relativity.  Specifically, the paper is organized as follows. In section II, we study the property of $q(z)$ by fitting the parametrization $q(z)=1/2+(q_1 z+q_2)/(1+z)^2$ to the 157 SN Ia data and the 115 nearby SN Ia and the Supernova Legacy Survey (SNLS) SN Ia data compiled in \\cite{astier}. In section III, we apply the parametrization $w(z)=w_0+w_1 z/(1+z)^2$ to study the property of $q(z)$ and the nature of dark energy. The Sloan Digital Sky Survey (SDSS) \\cite{sdss} and the Wilkinson Microwave Anisotropy Probe (WMAP) data \\cite{wmap} are also combined with the SN Ia data in our analysis. In section IV, we conclude the paper with some discussion. ",
        "conclusions": "By fitting the parametrization $q(z)=1/2+(q_1 z+q_2)/(1+z)^2$ to the 157 gold sample SN Ia data, we find that $q(z)<0$ for $0\\le z \\alt 0.2$ within the uncertainty of $3\\sigma$ level. Recall that $q(z)<0$ means that the Universe is in the acceleration phase at redshift $z$. We also get $z_t=0.95^{+3.25}_{-0.58}$ at the $1\\sigma$ level. If we fit the parametrization to the 115 nearby and SNLS SN Ia data, we find that $q_0>0$ at the $3\\sigma$ level and strong evidence that the Universe once accelerated. We also find that $z_t=0.36^{+0.24}_{-0.08}$ at the $1\\sigma$ level. Note that these conclusions do not depend on any particular theory of gravity.  To study the role that general relativity might play, in this paper we also considered the Friedmann equation together with  the dark energy parametrization $w(z)=w_0+w_1 z/(1+z)^2$. After applying them to the 157 gold sample SN Ia data or the 115 SN Ia data, we find that $q_0>0$ at the $1\\sigma$ level and the best fitting $\\Omega_{m0}\\sim 0.4$, which is higher than that determined by other observations. In the parametrization $q(z)$, we have two parameters $q_1$ and $q_2$. When we use the dark energy parametrization, there are three parameters $\\Omega_{m0}$, $w_0$ and $w_1$. With the addition of one more parameter apart from the assumption that the acceleration is due to dark energy in the framework of general relativity, we expect that the constraints will be loose. Therefore the dark energy parametrization is then fitted to the combined SN Ia and SDSS data because the combined data will break the degeneracies among the three parameters. If we fit the dark energy model to the combined 157 gold sample SN Ia and SDSS data, we find that $q(z)<0$ for $0\\le z \\alt 0.2$ contains over 99.7\\% of the probability. We also get $z_t=0.30^{+0.23}_{-0.06}$ at the $1\\sigma$ level. If we fit the dark energy model to the combined 115 SN Ia and SDSS data, we find that $q_0>0$ at the $2\\sigma$ level and strong evidence that the Universe once accelerated. The transition redshift is found to be $z_t=0.77^{+0.08}_{-0.36}$ at the $1\\sigma$ level. Whether we use the parametrization $q(z)$ or the dark energy parametrization $w(z)$ to fit the SN Ia data, the evidence that the the Universe once accelerated is very strong. These results confirm the conclusion that there is strong evidence that the Universe once accelerated obtained in \\cite{sturner}. Our results also suggest that $z_t\\agt 0.2$. Although the upper bound on $z_t$ is too large by using the parametrization $q(z)$, $z_t$ is more tightly constrained to be $z_t\\alt 1.0$ by using the dark energy parametrization $w(z)$.  The situation about $q_0$ is more subtle. It depends on the model and the data. If we fit the models to the 157 gold sample SN Ia data, we find strong evidence for $q_0<0$. But if the models are fitted to the 115 SN Ia data, the conclusion is different. The discrepancy is caused by the difference of the redshift range the data probed. The 157 gold sample SN Ia data probes much deeper in the redshift ($z\\sim 1.7$) than the 115 SNLS data which stops at $z\\sim 1$. For the dark energy model, the evidence for current deceleration looks promising if the model is fitted to the SN Ia data alone. The fitting also gives larger value for $\\Omega_{m0}$. When the SN Ia data is combined with the SDSS data for the dark energy model, we get reasonable value for $\\Omega_{m0}$ and the sign of $q_0$ is uncertain. The uncertainty of the question whether the Universe in accelerating now is directly related to the value of $w_0$. If $w_0\\lesssim -0.6$, then the Universe is accelerating currently. If $w_0<-1$, then the dark energy behaves like a phantom and it is not explained by the quintessence field. In fact, different models mean different external priors, so it shows that whether the Universe is currently accelerating depends on external prior.  In conclusion, we confirm that there is strong evidence that the Universe once accelerated."
    },
    "0601/astro-ph0601122_arXiv.txt": {
        "abstract": "Oxygen abundances in the spiral galaxies expected to be richest in oxygen are estimated. The new abundance determinations are based on the recently discovered ff -- relation between auroral and nebular oxygen line fluxes in high-metallicity H\\,{\\sc ii} regions. We find that the maximum gas-phase oxygen abundance in the central regions of spiral galaxies is 12+log(O/H) $\\sim$ 8.75. This value is significantly lower (by a factor $\\ga$ 5) than the previously accepted value. The central oxygen abundance in the Milky Way is similar to that in other large spirals. ",
        "introduction": "The oxygen abundance in spiral galaxies has been  the subject of much discussion for a long time.  Which spiral galaxy is the oxygen-richest one? And how high is the oxygen abundance in the oxygen-richest galaxies? Due to the presence of radial abundance gradients in the disks of spiral galaxies, the maximum oxygen abundance occurs at their centers. The notation (O/H)$_{\\rm C}$ = 12+log(O/H)$_{R=0}$ will be used thereafter to denote the central oxygen abundance.  The  chemical composition of various samples of spiral galaxies has been discussed in a number of papers \\citep[e.g.][]{vila92,zkh,garnettetal97,vanzeeetal98,garnett02}. According to those articles, the oxygen-richest galaxies are: NGC~5194 (M~51) with (O/H)$_{\\rm C}$ = 9.54 \\citep{vila92}, NGC~3351 with (O/H)$_{\\rm C}$ = 9.41 \\citep{zkh}, NGC~3184 with (O/H)$_{\\rm C}$ = 9.50 \\citep{vanzeeetal98}, NGC~6744 with (O/H)$_{\\rm C}$ = 9.64 \\citep{garnettetal97}. Different versions of the one-dimensional empirical method, proposed first by \\citet{pageletal79} a quarter of a century ago, have been used for oxygen abundance determinations in those papers. \\citet{lcal,hcal,m101,vybor} has shown that the oxygen abundances in galaxies determined with the one-dimensional empirical calibrations are significantly overestimated at the high-metallicity end ( 12 + log O/H $>$ 8.25). This for two reasons. First, the then-existing  calibrating points at the high-metallicity end were very few and not reliable. Second, the physical conditions in H\\,{\\sc ii} regions cannot be taken into account accurately in one-dimensional calibrations. \\citet{pilvilcon04} have used instead a two-dimensional parametric empirical calibration to derive oxygen abundances for a sample of spiral galaxies which includes NGC~3184, NGC~3351, NGC~5194, and NGC~6744. They found generally lower oxygen abundances with  (O/H)$_{\\rm C}$ $\\sim$ 9.0.  Here we consider anew the problem of the maximum oxygen abundance in spiral galaxies by attempting to derive more accurate oxygen abundances. These can be derived via the classical T$_{\\rm e}$ -- method, T$_{\\rm e}$ being the electron temperature of the H\\,{\\sc ii} region. Measurements of the auroral lines, such as [OIII]$\\lambda$4363, are necessary to determine T$_{\\rm e}$. Unfortunately, they are very faint and often drop below the detectability level in the spectra of high-metallicity H\\,{\\sc ii} regions. \\citet{ff} has advocated that the faint auroral line flux can be computed from the fluxes in the strong nebular lines via the ff -- relation. Then, using the obtained flux in the auroral line, accurate oxygen abundances can be derived using the classical T$_{\\rm e}$ -- method. We will estimate (O/H)$_{\\rm C}$ in the spiral galaxies reported to be the oxygen-richest ones, NGC~3184, NGC~3351, NGC~5194, and NGC~6744. For comparison, we will also consider  NGC~2903, NGC~2997, NGC~5236 as well the Milky Way Galaxy.   We describe the method used to determine oxygen abundances in H\\,{\\sc ii} regions in Section 2. The oxygen abundances in the spiral galaxies NGC~3184, NGC~3351, NGC~5194, NGC~6744, NGC~2903, NGC~2997, NGC~5236 and the Milky Way Galaxy are determined in Section 3. We discuss the reliability of the derived abundances and summarize our conclusions in Section 4.   For the line fluxes, we will be using the following notations throughout the paper: R$_2$ = $I_{[OII] \\lambda 3727+ \\lambda 3729} /I_{H\\beta }$, R$_3$ = $I_{[OIII] \\lambda 4959+ \\lambda 5007} /I_{H\\beta }$, R = $I_{[OIII] \\lambda 4363} /I_{H\\beta }$, R$_{23}$ = R$_2$ + R$_3$. With these definitions, the excitation parameter P can then be expressed as: P = R$_3$/(R$_2$+R$_3$). ",
        "conclusions": "The (O/H)$_{\\rm ff}$ abundances determined here rely on the validity of the classic T$_e$ -- method. This has been questioned for the high-metallicity regime in a number of investigations by comparison with H\\,{\\sc ii} region photoionization models. One should keep in mind however that the existing numerical models of H\\,{\\sc ii} regions are far from being perfect \\citep{stasinska04}, and, hence, the statement that H\\,{\\sc ii} region models provide more realistic abundances as compared to the T$_e$ -- method should not be taken for granted \\citep{vybor}. Recently \\citet{stasinska05} has examined the biases involved in T$_{\\rm e}$-based abundance determinations at high metallicities. She has found that, as long as the metallicity is low, the derived (O/H)$_{\\rm T_e}$ value is very close to the real one. Discrepancies appear around 12+log(O/H) = 8.6 - 8.7, and may become very large as the metallicity increases \\citep[Fig. 1a in][]{stasinska05}. The derived (O/H)$_{\\rm T_e}$ values are smaller than the real ones, sometimes by enormous factors. This means that, if such metal-rich H\\,{\\sc ii} regions exist, the classic T$_e$ -- method will always lead to systematically lower derived oxygen abundances.   Can the effect discussed by Stasi\\'{n}ska be responsible for the low (O/H)$_{\\rm T_e}$  abundances obtained here? Let us assume that it is the case, and that the true central oxygen abundances in the spiral galaxies are higher than 12 + log(O/H) = 9.0. What radial distributions of (O/H)$_{\\rm T_e}$ abundances in the disk of spiral galaxies would one expect in this case? Starting from the periphery of a galaxy, the (O/H)$_{\\rm T_e}$  abundances would increase with decreasing galactocentric distance (following the true abundances) until the (O/H)$_{\\rm T_e}$  abundance reaches the value of 12+log(O/H) $\\sim$ 8.7. After that the (O/H)$_{\\rm T_e}$ abundances would decrease with decreasing galactocentric distance because of the Stasi\\'{n}ska \\citep[Fig. 1a in][]{stasinska05} effect, although the true abundances continue to increase with decreasing galactocentric distance. Thus, the radial distribution of (O/H)$_{\\rm T_e}$  abundances should show a bow-shaped curve with a maximum value of 12+log(O/H) $\\sim$ 8.7 at some galactocentric distance.  Do the actual data (Fig. 5--7) show such a behavior? The derived radial distributions of (O/H)$_{\\rm T_e}$  abundances in the disks of spiral galaxies do not show an appreciable curvature, and the (O/H)$_{\\rm T_e}$  abundances increase more or less monotonically with decreasing galactocentric distance. Only the H\\,{\\sc ii} regions in the very central parts of the spiral galaxies NGC~3351 (the middle panel of Fig.~\\ref{figure:spmax}) and NGC~2997 (the middle panel of Fig.~\\ref{figure:splow}) may show (O/H)$_{\\rm T_e}$ abundances which are slightly lower than expected from the general radial abundance trends, but the effect is very slight and may not be real if errors of 0.1 dex (Fig. 3) in the calibration and in the line intensity measurements are taken into account. To settle the question of whether the abundances in the metal-richest H\\,{\\sc ii} regions are affected by the Stasi\\'{n}ska's effect, many more accurate measurements of H\\,{\\sc ii} regions (including those in the very central parts) in the most oxygen-rich spiral galaxies are necessary.  Thus, while our results do not rule out the possible existence of the Stasi\\'{n}ska's effect, they suggest that the great majority of H\\,{\\sc ii} regions in galaxies are in a metallicity range where this effect is not important. Indeed, the Stasi\\'{n}ska's effect becomes important only at oxygen abundances higher than 12 + log(O/H) $\\sim$ 8.7. We have found that this \"critical\" value is reached only in the central part of the most oxygen-rich galaxies. We conclude that, in the light of our present results, H\\,{\\sc ii} regions with 12+log(O/H) $>$ 8.7 are rare, if they exist at all. It should be noted that Stasi\\'{n}ska also questioned the existence of such metal-rich H\\,{\\sc ii} regions \\citep[Section 2.2 in][]{stasinska05}.  Another factor that can affect T$_{\\rm e}$-based abundance determinations is the temperature fluctuations inside H\\,{\\sc ii} regions \\citep{peimbert67}. If they are important, the (O/H)$_{\\rm T_e}$ abundance would be a lower limit. In most cases the effect appears to be small, probably under 0.1 dex \\citep{pagel03}.  Fortunately, there is a way to verify the validity of the T$_{\\rm e}$ -- method at high metallicities  \\citep{vybor}. High-resolution observations of the weak interstellar O{\\sc i}$\\lambda$1356 absorption lines towards stars allow one to determine the interstellar oxygen abundance in the solar vicinity with a very high precision. It should be noted that this method is model-independent.  These observations yield a mean interstellar oxygen abundance of 284 -- 390 O atoms per 10$^6$ H atoms (or 12+log(O/H) = 8.45 -- 8.59) \\citep{meyeretal98,sofiameyer01,cartledgeetal04}. \\citet{oliveiraetal05} have determined a mean O/H ratio of 345 $\\pm$19 O atoms per 10$^6$ H atoms (or 12+log(O/H) = 8.54) for the Local Bubble. Thus, an oxygen abundance 12+log(O/H) = 8.50 in the interstellar gas at the solar galactocentric distance seems to be a reasonable value. The value of the oxygen abundance at the solar galactocentric distance traced by (O/H)$_{\\rm T_e}$ abundances in H\\,{\\sc ii} regions is then in good agreement with the oxygen abundance derived with high precision from the interstellar absorption lines towards stars (open square in the bottom panel of Fig.~\\ref{figure:mwg}). This is strong evidence that the classic T$_{\\rm e}$ -- method provides accurate oxygen abundances in high-metallicity H\\,{\\sc ii} regions.  The measured oxygen abundance in the solar vicinity provides an indirect way to check our derived central oxygen abundances. The oxygen abundance in the solar vicinity is 12+log(O/H)=8.50, and a gas mass fraction $\\mu$ $\\sim$ 0.20 seems appropriate. The simple model of chemical evolution of galaxies predicts that a decrease of $\\mu$ by 0.1 results in a increase of the oxygen abundance by $\\sim$ 0.12 dex in the range of $\\mu$ from $\\sim$ 0.75 to $\\sim$ 0.05. When all the gas in the solar vicinity will be converted into stars, the oxygen abundance will increase by 0.2 -- 0.25 dex and will reach a value around 12+log(O/H) = 8.70 -- 8.75. This value is in excellent agreement with the central oxygen abundances obtained above for the Milky Way and other galaxies. There is no room for central oxygen abundances as large as 12+logO/H = 9.00 or greater.  The present study suggests that there is an upper limit to the oxygen abundances in spiral galaxies. If true, the existing luminosity -- metallicity relation should be revisited. One can expect the metallicity to increase with luminosity up to 12+(O/H) $\\sim$ 8.75, but to remain approximately constant for higher metallicities, resulting in a flattening of the luminosity -- metallicity relation. This suggestion will be investigated in a future paper.  The maximum gas-phase oxygen abundance in spiral galaxies is 12+log(O/H) $\\sim$ 8.75. The central oxygen abundance in the Milky Way is similar to that in other large spirals. Some fraction of the oxygen is locked into dust grains in H\\,{\\sc ii} regions. \\citet{estebanetal98} found that the fraction of the dust-phase oxygen abundance in the Orion nebula is about 0.08 dex. Then, the true gas+dust maximum value of the oxygen abundance in H\\,{\\sc ii} regions of spiral galaxies is 12+log(O/H) $\\sim$ 8.85. This value can increase up to $\\sim$ 8.90 -- 8.95 if temperature fluctuations inside H\\,{\\sc ii} regions are important.  \\subsection*"
    },
    "0601/astro-ph0601314_arXiv.txt": {
        "abstract": "New images of M31 at 24, 70, and 160~\\micron\\ taken with the Multiband Imaging Photometer for Spitzer (MIPS) reveal the morphology of the dust in this galaxy.  This morphology is well represented by a composite of two logarithmic spiral arms and a circular ring (radius $\\sim$~10~kpc) of star formation offset from the nucleus.  The two spiral arms appear to start at the ends of a bar in the nuclear region and extend beyond the star forming ring.  As has been found in previous work, the spiral arms are not continuous but composed of spiral segments.  The star forming ring is very circular except for a region near M32 where it splits.  The lack of well defined spiral arms and the prominence of the nearly circular ring suggests that M31 has been distorted by interactions with its satellite galaxies.  Using new dynamical simulations of M31 interacting with M32 and NGC 205 we find that, qualitatively, such interactions can produce an offset, split ring like that seen in the MIPS images. ",
        "introduction": "\\label{sec_intro}  M31 is the nearest \\citep[d~$\\sim 780$~kpc,][]{Stanek98,Rich05} external spiral galaxy \\citep[Sb,][]{Hubble29} making it one of the premier laboratories for understanding spiral galaxies like our own. Tracing the spiral structure in M31 is an obvious first step, which has been done with varying degrees of success for many years. However, the high inclination (74-78$\\arcdeg$) of M31 makes it difficult to determine its large-scale structure.  In particular, there is no clear consensus regarding the organization of spiral structure in this important galaxy.  The earliest quantitative studies focused on optical tracers such as HII regions \\citep{Arp63, Baade64}, OB associations \\citep{Hodge79, vandenBergh91}, and dark nebulae \\citep{Hodge80}.  These studies generally found structures consistent with a two-armed spiral \\citep[for the case for a one-armed spiral see][]{Simien78}, but invoked disturbances such as varying inclinations and warps to account for the deviations from simple logarithmic spirals.  More recently, HI \\citep{Guibert74, Braun91} and CO \\citep{Loinard99, Nieten05} have been used.  The velocity information available for these gas tracers can be used to disentangle the pileup of information due to M31's high inclination.  These studies have found strong evidence for a two-armed spiral structure, but still suffer from the difficulty in connecting spiral arm segments into a coherent pattern.  The difficulty in finding coherent spiral structure in M31 is plausibly related to the interactions of M31 with its two nearest satellites, M32 and NGC 205.  Disturbances in M31 are revealed by tidal features far outside M31's disk visible in deep, wide-field optical imaging \\citep{Ibata01,Ibata05,Ferguson02, Lewis04}. M32 is understood to be the perturber of the spiral arms \\citep{Byrd78, Byrd83} while NGC 205 has been modeled as the cause of the warp seen in the optical and HI disk of M31 \\citep{Sato86}.  Quantifying the effects of M31's satellite galaxies is complicated by the lack of measured proper motions and uncertainties in relative distances \\citep[e.g.][]{McConnachie05}.  Another difficulty in quantifying M31's spiral structure has been the lack of appropriate tracers in the innermost regions of the galaxy. Optical tracers are overwhelmed by the bright stellar bulge, while HI and CO, which would be easily seen through the bulge, are largely absent \\citep{Brinks84, Melchior00} as the interstellar medium (ISM) in the central region is composed mainly of ionized gas \\citep{Devereux94}.  The far-infrared (far-IR) emission can circumvent these difficulties as it traces the dust in M31, which penetrates the stellar bulge and should trace all phases of the ISM.  \\begin{figure*} \\epsscale{1.05} \\plotone{f1_small.eps} \\caption{The M31 MIPS images are shown.  The images are scaled using a quadratic root ($x^{0.25}$) between -0.1 -- 3 MJy/sr (24~\\micron), 0.1 -- 35 MJy/sr (70~\\micron), and 0 -- 90 MJy/sr (160~\\micron).  The bottom image gives the 3 color composite of the MIPS images (with scaling adjusted to emphasize the differences between bands); the color gives an indication of the dust temperature.  The images are oriented with the nominal position angle of M31 ($38\\arcdeg$) horizontal and are cropped to a common size.  The nearly-vertical streaks present in the 70 and 160~\\micron\\ images are residual Ga:Ge detector instrumental signatures.  M32 is clearly detected at 24~\\micron, but not in the other two MIPS bands.  NGC~205 is detected in all three MIPS bands, but is located NW of M31 beyond the boundary of the cropped images. The two spots near the nucleus marked on the 160~\\micron\\ image are discussed in the text. \\label{fig_m31_fullres} } \\end{figure*}  The first far-IR images of M31 were obtained with the Infrared Astronomical Satellite (IRAS) at 12, 25, 60, and 100~\\micron.  The most prominent feature in these images is the bright, $\\sim$10~kpc radius, ring of star formation \\citep{Habing84, Rice93} seen previously in H$\\alpha$ images \\citep{Arp63, Devereux94}.  The Infrared Space Observatory (ISO) imaged M31 at 175~\\micron\\ finding both the bright ring and a fainter ring at larger radii \\citep{Haas98}.  The Midcourse Space Experiment (MSX) provided an image of M31 at 8.3~\\micron\\ revealing possible spiral structure inside the $\\sim$10~kpc ring.  The resolutions of the IRAS ($\\sim 1 \\arcmin$) and ISO ($1\\farcm 3$) images reveal the overall morphology, but are not high enough to reveal the structure of the spiral arms.  We present new infrared images of M31 taken with the Multiband Imaging Photometer for Spitzer \\citep[MIPS,][]{Rieke04MIPS} instrument on the Spitzer Space Telescope \\citep[Spitzer,][]{Werner04Spitzer}.  The MIPS spatial resolution at the distance of M31 is $6\\arcsec$/23~pc, $18\\arcsec$/68~pc, and $40\\arcsec$/151~pc at 24, 70, and 160~\\micron, respectively.  This paper concentrates on the newly-revealed details of M31's infrared morphology as shown in the MIPS images ",
        "conclusions": "We have presented new MIPS images of M31 at 24, 70, and 160~\\micron\\ which have higher spatial resolution and sensitivity than previous infrared images.  These new images reveal the presence of two spiral arms in addition to a ring of star formation.  The spiral arms appear to emanate from a bar in the nuclear region and can be roughly traced out beyond the $\\sim$10~kpc ring of star formation. The ring of star formation forms an almost complete circle which is offset from the nucleus and split near the location of M32.  This ring of star formation and the disturbed spiral arms are indications that interactions may be affecting the morphology of M31.  Dynamical models of M31 interacting with M32 predict morphology qualitatively similar to M31's (while models with NGC~205 instead of M32 give only small perturbations).  Specifically, both the splitting and the offset of the star forming ring are seen as well as copious spiral arms.  It seems that M31 is not an undisturbed normal spiral galaxy, but one which has been significantly affected by interactions.  This paper presents a preview of the possibilities becoming available to understand M31.  Data covering the majority of this galaxy are available in the X-ray (XMM, Trudolyubov et al., in prep.), ultraviolet \\citep[GALEX,][]{Thilker05}, optical (Engelbracht et al., in prep.), mid-infrared (Spitzer/IRAC, Barmby et al., in prep.), infrared (this paper), HI (Braun et al., in prep.), CO \\citep{Nieten05}, and radio continuum \\citep{Beck89}; these will enable us to make detailed models on the impact that satellite galaxies have on the morphology of M31."
    },
    "0601/astro-ph0601587_arXiv.txt": {
        "abstract": "The Spectroscopy of Plasma Evolution from Astrophysical Radiation (or the Far-ultraviolet Imaging Spectrograph) instruments, flown aboard the STSAT-1 satellite mission, have provided the first large-area spectral mapping of the cosmic far ultraviolet  (FUV, $\\lambda$ 900-1750)  background. We observe diffuse radiation from hot (10$^{4}$ -- 10$^{6}$ K) and ionized plasmas, molecular hydrogen, and dust scattered starlight. These data provide for the unprecedented detection and discovery of spectral emission from a variety of interstellar environments, including the general medium, molecular clouds, supernova remnants, and super-bubbles. We describe the mission and its data, present an overview of the diffuse FUV sky's appearance and spectrum, and introduce the scientific findings detailed later in this volume. ",
        "introduction": "The Spectroscopy of Plasma Evolution from Astrophysical Radiation instruments (hereafter {\\em SPEAR}, also known as the Far-ultraviolet Imaging Spectrograph or {\\em FIMS}) have provided the first large-area spectral sky survey of cosmic far ultraviolet (FUV) radiation. The FUV band ($\\lambda$ 900-1750) includes important astrophysical diagnostics such as  strong atomic cooling lines from hot (T=10$^{4}$ K -- 10$^{6}$ K) and photoionized plasmas as well as radiation from molecular hydrogen (H$_{2}$) fluorescent emission and dust scattered starlight. We describe the {\\em SPEAR} mission and its science motivation, introduce the data by presenting the spatial distribution of the FUV ($\\lambda$ 1360 -1710) sky brightness and the spectra of the entire observed sky, and provide a brief introduction to other $SPEAR$ science results described in the accompanying papers in this Volume.  {\\em SPEAR} is the primary payload on the STSAT-1 (S-1) satellite launched 2003 September 27. The mission has thus far observed $\\sim$80\\% of the sky and conducted deep pointed observations toward numerous selected targets. {\\em SPEAR} contains dual imaging  spectrographs optimized for the measurement of diffuse FUV emission. The spectrographs, referred to as the Short wavelength band 'S'  ($\\lambda\\lambda $ 900 -- 1150, 4.0$^{\\circ}$ x 4.6$\\arcmin$ view) and Long wavelength band 'L':  $\\lambda\\lambda$ 1350 -- 1750, 7.4$^{\\circ}$ x 4.3$\\arcmin$ view), each have a spectral resolution of \\ensuremath{\\lambda /\\Delta \\lambda \\sim 550} half-energy width and an imaging resolution of 5$\\arcmin$. Each instrument uses a collecting mirror, a diffraction grating and an open faced, photon counting micro-channel plate detector. The {\\em SPEAR} instruments, their on-orbit performance, and the basic processing of instrument data  are described in detail in the following paper (Edelstein et al. \\textit{ibid}.). ",
        "conclusions": "The first FUV spectral imaging survey of the large fractions of the sky has been taken with the {\\it SPEAR} mission. The resulting map of total FUV radiation shows that diffuse FUV cosmic radiation is concentrated where both hot stars and scattering dust coexist, such as in the Galactic plane, young stellar associations, and the Magellanic clouds. The spectra of the total-sky contains prominent lines from ionized gas, including the previously observed \\CIV\\, and \\SiIV\\, emission, and from newly discovered \\CI, \\SiIIs, \\AlII, \\HeII \\, emission, as well as from H$_{2}$ fluorescence. Similar interstellar FUV emission lines are found in the general ISM, as sampled at the NEP, and toward the Eri-Ori superbubble and its interface to the ISM. Intense FUV emission lines from shocks are visible in nearby SNR. Diffuse H$_{2}$ fluorescence is ubiquitous across observed Galactic environments. Much work remains to elaborate upon the character of the FUV sky and to analyze the physical conditions of detailed objects."
    },
    "0601/astro-ph0601534_arXiv.txt": {
        "abstract": "We have obtained a near infrared spectrum of the brightest satellite of the large Kuiper Belt Object, 2003 EL61.  The spectrum has absorption features at 1.5 and 2.0 microns, indicating that water ice is present on the surface. We find that the satellite's absorption lines are much deeper than water ice features typically found on Kuiper Belt Objects.  We argue that the unusual spectrum indicates that the  satellite was likely formed by impact and not by capture. ",
        "introduction": "Satellites in the Kuiper Belt provide unique insights into the internal properties of icy bodies as well as the early dynamical history of the trans-Neptunian region. To date, over a dozen satellites have been found around Kuiper Belt Objects (KBOs) and the overall satellite fraction is estimated to be 11$\\pm ^5_2 $\\% (Stephens et al., 2005). The first KBO satellites discovered (excluding the Pluto-Charon system) were found to have high eccentricities, large orbital separations, and have similar brightness to the primary, indicating that the satellites are nearly as large as their parent.  Various capture mechanisms were proposed to account for the high specific angular momentum of these systems (Goldreich et al., 2002). Recently,  Brown et al. (2006) have found several satellites around the largest known KBOs, and these satellites appear to be different than those discussed above. The satellite fraction appears to be higher around large KBOs and, like the Pluto-Charon system, the satellites have lower fractional brightness compared to the primary. The orbit of the brightest satellite around 2003 EL61 is nearly circular and  Brown et al. (2006) suggest that these systems likely formed by a different mechanism and may be the result of giant impacts (McKinnon 1989, Canup 2005). The density, albedo, and surface composition of these satellite systems may help to clarify the origins of satellites in the Kuiper belt.   Near infrared spectroscopy is a particularly well suited method for remotely studying the surface composition of KBOs. Many of the abundant species in the outer solar system, including CH$_4$, CO$_2$, and H$_2$O ices, have absorption features in the near IR. A small amount of water ice has been observed on a number of KBOs  and some show evidence for the presence of other ices (Fronasier et al. 2004 and references therein).  The spectra of satellites may shed light their origin. Captured objects are likely to have spectra similar to other KBOs, but satellites formed from other mechanisms could show unusual or altered surface compositions.  2003 EL61 is currently the third brightest known KBO and is elliptical in shape with axes of 1960  and 2500 km (Rabinowitz et al., 2006). Its satellites offer an excellent opportunity to study the formation and dynamics of KBO binaries. The brighter satellite, S/2005~(2003~EL61)~1, has a period 49.12 $\\pm$ 0.03 days and a maximum separation of 1.4 arcseconds (Brown et al. 2005). The orbit is nearly circular, with an eccentricity of 0.05.  The inner satellite, S/2005~(2003~EL61)~2, appears to have an orbit of 34.7 $\\pm$ 0.1 days, though there is currently insufficient data to reliably fit a non-circular orbit.  The inner satellite has a fractional brightness of 1.5\\% compared to 2003 EL61, making it difficult to study with spectroscopic techniques. The outer, brighter satellite has a fractional brightness of 5\\% and is a viable target for low resolution infrared spectroscopy. We report here on the near-IR (J through K band) spectrum of  2003 EL61's brighter satellite. We compare the spectrum of this satellite to 2003~EL61's as well as to other KBOs.  We discuss possible formation scenarios for the satellite and find that the surface composition is consistent with formation by giant impact. ",
        "conclusions": "The satellite spectrum shows strong absorption features at 1.5 and 2 microns, which are consistent with absorption due to water ice (Fig \\ref{fig2}). However, the resolution is too poor to detect the 1.65 $\\mu$m crystalline water ice absorption feature seen on 2003 EL61 (Trujillo et al., submitted). The depth of the water lines suggests that much of the satellite's surface is covered with smooth water ice. We calculated a model water ice spectrum using laboratory optical constants from Grundy \\& Schmitt (1998), assuming a temperature of 50 K. Our model uses Hapke theory (Hapke 1981) to transform the optical constants to a reflectance spectrum. We vary the grain size of ice in our models and find that the grain size of 250 microns provides a good fit to the absorption lines at 1.5 and 2 microns.  A better fit may be achieved with the addition of a second blue component or a multiple grain size water ice model.   The absorption features in the spectrum of  the satellite are significantly deeper than water ice features typically observed on KBOs. The ratio of flux at 1.7 $\\mu$m to the flux at 2.0 $\\mu$m is $\\sim$ 80\\%. Most KBOs with water ice features have depths of less than $\\sim$20\\% and the deepest absorptions ($\\sim$60 \\%) are found on 2003~EL61, Charon and 1996~TO66 (Brown et al. 2000, Fronasier et al 2004, Dotto et al 2003 and references therein).  As such, the brighter satellite of 2003 EL61 appears to be an unusual object in the Kuiper belt. Captured objects are likely to resemble the parent population where they are formed, thus the satellite's spectrum suggests that formation through capture is unlikely.   An alternative to capture is that the satellite could have formed via impact.  Models of the Charon-forming impact by Canup (2005) are consistent with many of the features of the 2003 EL61 system.  In these models, Charon is formed from a grazing impact where the impactor remains largely intact and creates a large satellite.  However, the models also show that a more direct impact creates a disk around the target body and smaller satellites can eventually coalesce from the disk material.  The small satellites around 2003 EL61 are more likely to be formed by the latter kind of impact than by the Charon-type grazing impacts. In models where a disk was created, the satellite preferentially forms from the less dense disk material. In the outer solar system, this trend will lead to water-enriched satellites around KBOs,  which may account for the strong water ice features seen in the satellite's spectrum.   Direct impacts can also spin up the target body, which is consistent with the high rotation rate observed for 2003 EL61 (Rabinowitz et al., 2005).   2003 EL61 has a mean density of 2.6-3.4 g/cm$^3$, indicating  that it has lost most of its volatiles (Rabinowitz et al. 2005). Further modeling is needed to show if a high energy impact can eject a substantial amount of ice off of 2003 EL61 to account for its high rock-to-ice ratio.  In sum, 2003 EL61 and its brighter satellite exhibit many of the features predicted by Kuiper belt impact models. Further analysis of other satellites may provide a greater understanding of the dynamical history of the early Kuiper belt.      {\\it Acknowledgments:}  This research is funded by the California Institute of Technology and is also supported by the NASA Planetary Astronomy program. 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 authors thank Chad Trujillo, David Rabinowitz, and Robin Canup for their helpful discussions regarding this paper. We also thank Antonin Bouchez for his help during observations."
    },
    "0601/astro-ph0601702_arXiv.txt": {
        "abstract": "{There are several possibilities to constrain the value of the magnetic field in the ICM, the most direct ones being the combination of inverse Compton and synchrotron observations, and the Faraday rotation measures. Here we discuss on the possibility to provide constraints on the magnetic field in the ICM from the analysis of the statistical properties of the giant radio halos, Mpc--scale diffuse radio emission in galaxy clusters. Present observations of a few well studied radio halos can be interpreted under the hypothesis that the emitting relativistic electrons are re-accelerated on their way out. By using statistical calculations carried out in the framework of the re-acceleration model we show that the observed radio--power vs cluster mass correlation in radio halos can be reproduced only by assuming $\\mu$G fields in the ICM and a scaling of the magnetic field with cluster mass $B \\propto M_v^b$, with $b \\geq 0.6$. We also show that the expected occurrence of radio halos with mass and redshift, and their number counts are sensitive to the magnetic field intensity in massive galaxy clusters and to the scaling of B with cluster mass. Thus future deep surveys of radio halos would provide constraints on B in galaxy clusters. ",
        "introduction": "Giant radio halos (size $\\sim 1$ $h_{50}^{-1}$ Mpc, GRHs elsewere) are the most spectacular evidence for non-thermal components (relativistic particle and magnetic field) in the intra cluster medium (ICM). They are diffuse synchrotron radio emission extending on Mpc scale, have no obvious connection with the cluster galaxies, but are rather associated with the ICM (e.g., Giovannini \\& Feretti 2000). Such a synchrotron radio emission requires a population of GeV relativistic electrons (and possibly positrons) and cluster magnetic fields on $\\mu$G levels. An independent evidence of the existence of cluster magnetic fields in the ICM is given by the rotation measurement studies of radio galaxies located within or behind the clusters (e.g., Govoni \\& Feretti 2004). One possibility to explain the presence of GeV electrons diffused on Mpc scales is given by the so-called {\\it re-acceleration} model. In this model the radio emitting electrons injected in the ICM (by AGN, starbursts, supernovae and/or galactic winds, and hadronic collisions) are re-accelerated {\\it in situ} by some kind of turbulence generated in the cluster volume during cluster mergers (e.g., Brunetti et al. 2001, 2004; Petrosian 2001; Fujita et al. 2003; Brunetti \\& Blasi 2005). Observations indicate that the detection rate of GRHs shows an abrupt increase with increasing the X-ray luminosity of the host clusters. In particular about 30-35\\% of the galaxy clusters with X-ray luminosity larger than $10^{45}$ erg/s show diffuse non-thermal radio emission (Giovannini \\& Feretti 2002); these clusters have also high temperature (kT $>$ 7 keV) and large mass ($\\gtsim\\, 2\\times$ $10^{15} M_{\\odot}$). Furthermore, GRHs are always found in merging clusters (e.g., Buote 2001). Although the {\\it re-acceleration} model seems to reproduce the observational features of the diffuse radio emission, a theoretical investigation of the statistical properties of the GRHs in galaxy clusters in the framework of this model has not yet carried out extensively. Only recently, Cassano \\& Brunetti (2005; CB05 elsewhere), have calculated the expected occurrence of GRHs as a function of the cluster mass and dynamical status in the framework of the re-acceleration model. CB05 follow cluster formation using the Press \\& Schechter (1974; PS hereafter) formalism and assume that a fraction, $\\eta_t$, of the $PdV$ work done by the merging subclusters in going through the main one is injected in the form of fast magnetosonic (MS) waves which accelerate relativistic electrons in the ICM. They show that GRHs are {\\it naturally} expected only in the more massive clusters with an expected occurrence (at $z\\ltsim\\, 0.2$) which can be reconciled with the observed one under viable assumptions ($\\eta_t\\,\\simeq\\,0.24-0.34\\,$). More recently, Cassano, Brunetti, Setti (2006, CBS06 elsewhere) extended this investigation and calculated the expected evolution with cosmic time of the fraction of galaxy cluster with GRHs, the expected luminosity function of GRHs and their number counts. In CBS06 it was assumed a scaling of the rms magnetic field with cluster mass $B \\propto M^b$ ($b<2$) and it was shown that the available radio-X correlations can be accounted for by assuming $b\\ge0.5-0.6$. A scaling of the magnetic field with cluster mass is indeed expected by numerical MHD simulations (e.g., Dolag et al. 2004) where the seed field is amplified by the effect of shear flows driven by the cluster formation process. These simulations show that the resulting magnetic field strengths depend on the final cluster mass and/or temperature as $B\\propto T^a$ with $a\\sim2$ (Dolag et al. 2002). Here we focus our attention on the possibility to make constraints on the value of the magnetic field strength in galaxy clusters by comparing the observed and the predicted statistical properties of GRHs. In addition, we show how expectations for future observations will be powerful tools to test this re-acceleration model and also to provide additional constraints on B in the ICM. We focus our attention on GRHs only (size $\\sim 1$ $h_{50}^{-1}$ Mpc, GRH elsewhere). The adopted cosmology is: $\\Lambda$CDM ($H_{o}=70$ Km $s^{-1}$ $Mpc^{-1}$, $\\Omega_{o,m}=0.3$, $\\Omega_{\\Lambda}=0.7$, $\\sigma_8=0.9$). ",
        "conclusions": "In this contribution we have shown that the observed correlations between radio and X-ray properties of galaxy clusters, together with other statistical properties of GRHs (occurrence as a function of mass and redshift and number counts), may be useful tools to test the particle acceleration model for the formations of GRHs. For the first time we have discussed in some details the magnetic field intensity and scaling with mass required in the framework of this model.  We made use of the statistical calculations developed in CBS06, which are based on the re-acceleration of a seed population of relativistic electrons by MS waves injected in the ICM during merger events (CB05) and which assume a dependence of the rms magnetic field intensity on cluster mass, $B =B_{<M>} (M/<M>)^b$. The main results of this paper are:  {\\it i)} A comparison of the expected correlation between the radio power and the cluster mass with the observed one allows the definition of a permitted region of the parameter space ($B_{<M>}$,b) in galaxy clusters, with typically small values of B for sublinear scaling ($b<1$) and higher values of B for superlinear scaling ($b>1$). A lower bound is found at $b\\sim 0.5-0.6$, while a lower bound $B_{<M>}=0.2\\, \\mu$G is also obtained from IC arguments (Sect.~2). A superlinear scaling of $B$ with mass, as expected by MHD simulations (Dolag et al.~2004) falls within the allowed region. The values of B in the superlinear case are close to (slightly smaller than) those obtained from rotation measurements (e.g., Govoni \\& Feretti 2004), which, however, generally sample regions which are even more internally placed than those spanned by GRHs.  {\\it ii)} We have shown that the occurrence of GRHs as a function of mass and redshift is sensitive to B in galaxy clusters. This probability depends on the merging history of clusters and on the relative importance of the synchrotron and IC losses, which indeed depends on the value of the magnetic field. We show that typically in the case of sublinear scaling of the magnetic field with cluster mass ($b \\sim 0.6-0.9$) the probability to have GRHs increases with cluster mass and decreases with redshift, whereas in the case of superlinear scalings ($b \\sim 1.2-1.7$) more complex behaviors of the probability with mass and redshift are present.   {\\it iii)} We have derived the integral number counts of GRHs at 1.4~GHz in both the superlinear and sublinear cases. We find that at higher fluxes ($>30-40$~mJy) the predicted number counts are dominated by the $z\\le0.2$ clusters discovered in the NVSS. We estimate that the number of GRHs which would be discovered below 30~mJy by future deeper radio surveys (by LOFAR, LWA, and SKA) will be up to $\\sim 100$ if superlinear scalings of the mass with B hold, while they will be up to $\\sim 50$ in the case of sublinear scaling."
    },
    "0601/astro-ph0601228_arXiv.txt": {
        "abstract": "s{ The statistical properties of interstellar turbulence are studied by means of three-dimensional high-resolution HD and MHD simulations of a SN-driven ISM. It is found that the longitudinal and transverse turbulent length scales have time averaged (over a period of 50 Myr) ratios of 0.5-0.6, almost similar to the one expected for isotropic homogeneous turbulence. The mean characteristic size of the larger eddies is found to be $\\sim 75$ pc. Furthermore, the scalings of the structure functions measured in the simulated disk show unambiguous departure from the Kolmogorv (1941) model being consistent with the latest intermittency studies of supersonic turbulence (Politano \\& Pouquet 1995; Boldyrev 2002). Our results are independent of the resolution, indicating that convergence has been reached, and that the unresolved smaller dissipative scales do not feed back on the larger ones.} ",
        "introduction": "Interstellar turbulence is mainly driven by the energy injected into the ISM by supernovae, with the driving scale still being uncertain. It is also unclear what the statistical properties of the turbulent interstellar gas are, if the full available range of energies is taken into account. In this paper we discuss the statistical properties of the interstellar turbulence in the Galactic disk based on three-dimensional adaptive mesh refinement (AMR) simulations of the ISM, which include the disk-halo-disk circulation [\\cite{avi00}]. In particular we explore the injection scales (section 2) and the scalings of the velocity structure functions (section 3) of the interstellar turbulent gas and discuss (section 4) their implications. ",
        "conclusions": ""
    },
    "0601/astro-ph0601472_arXiv.txt": {
        "abstract": "We present the results of our investigation of three samples kinematically representative of the thin and thick disks and the Hercules stream using the catalogue of Soubiran \\& Girard (2005). We have observed abundance trends and age distribution of each component. Our results show that the two disks are chemically well separated, they overlap greatly in metallicity and both show parallel decreasing trends of alpha elements with increasing metallicity, in the interval -0.80 $<$ [Fe/H] $<$ -0.30. The thick disk is clearly older than the thin disk with a tentative evidence of an AMR over 2-3 Gyr and a hiatus in star formation before the formation of the thin disk. In order to improve the statistics on the disk's abundance trends, we have developed an automatic code, TGMET{\\Large$\\alpha$}, to determine (Teff, logg, [Fe/H], [$\\alpha$/Fe]) for thousands of stellar spectra available in spectroscopic archives. We have assessed the performances of the algorithm for 350 spectra of stars bieng part of the abundance catalogue. ",
        "introduction": "We have compiled a large catalogue of stars from several studies from the literature presenting determinations of O, Na, Mg, Al, Si, Ca, Ti, Fe, Ni abundances (Soubiran \\& Girard 2005). Because authors of the different studies do not use the same scales and methods in their spectral analyses, systematics between their results have been investigated (for more details on the construction of the catalogue see Soubiran \\& Girard 2005). The final catalogue of abundances includes 743 stars.  In order to study kinematical groups of the Milky Way's disk we need velocities and orbits of stars from the abundance catalogue. We first cross-correlated the catalogue with Hipparcos (ESA 1997), selecting stars with $\\pi >$10 mas and $\\sigma_{\\pi} \\over \\pi$ $<$0.10. Then we have searched for radial velocities in several sources. Distances, proper motions and radial velocities have been combined to compute the 3 components (U, V, W) of the spatial velocities with respect to the Sun and the orbital parameters for 639 stars. In addition the derivation of ages was kindly done by Fr\\'ed\\'eric Pont making use of the Bayesian method of Pont \\& Eyer (2004). Good estimations of ages were obtained for 322 stars from the abundance catalogue. ",
        "conclusions": "The abundance trends for each kinematical group are shown in Fig. \\ref{f:figure_abu}. In addition we have represented the distribution of ages vs. [Fe/H] for stars having well-defined ages for the thin disk, the thick disk and the Hercules stream (see Fig. 12 in Soubiran \\& Girard 2005).\\\\ Our results confirm previous well established findings: \\begin{itemize} \\item The thin disk and the thick disk overlap in metallicity and exhibit parallel slopes of [$\\alpha$/Fe] vs [Fe/H] in the range -0.80 $<$ [Fe/H] $<$ -0.30, the thick disk being enhanced. \\item The thick disk is older than the thick disk. \\end{itemize} \\begin{figure*}[] \\centering \\includegraphics[width=6cm]{girard_fig1.eps} \\includegraphics[width=6cm]{girard_fig2.eps} \\includegraphics[width=6cm]{girard_fig3.eps} \\includegraphics[width=6cm]{girard_fig4.eps} \\caption[]{Averaged [X/Fe] vs [Fe/H] per bin of metallicity in the thin disk (green squares) in the thick disk (red filled circles) and in the Hercules stream (blue open circles). Errors bars correspond to the standard deviations around the mean value in each bin. Note that Si, Ca, Ti and Na elements are not represented here but can be found in Soubiran \\& Girard (2005).} \\label{f:figure_abu} \\end{figure*} We bring new constraints on more controversial issues: \\begin{itemize} \\item The thin disk extends down to [Fe/H]=-0.80 and exhibits low dispersions in its abundance trends. \\item The thick disk also shows smooth abundance trends with low dispersions. The change of slope which reflects the contribution of the different supernovae to the ISM enrichment is visible in [Si/Fe] vs [Fe/H] and [Ca/Fe] vs [Fe/H] at [Fe/H] $\\simeq$ -0.70, less clearly in [Mg/Fe] vs [Fe/H]. \\item Al behaves as an $\\alpha$ element. \\item $\\lbrack$O/Fe$\\rbrack$ decreases in the whole metallicity range with a change of slope at [Fe/H] = -0.50 for the 3 populations. \\item An AMR is visible in the thin disk, the most metal-poor stars having 6.2 Gyr on average, those with solar metallicity 3.9 Gyr. \\item  Ages in the thick disk range from 7 to 13 Gyr with an average of 9.6 $\\pm$ 0.3 Gyr. There is a tentative evidence of an AMR extending over 2-3 Gyr. \\item  The most metal rich stars assigned to the thin disk do not follow its global trends. They are significantly enhanced in all elements (particularly in Na and Ni) except in O which is clearly depleted. They have also a larger dispersion in age. Half of these stars are probable members of the Hyades-Pleiades supercluster, two others are surprisingly old. \\end{itemize}  \\begin{figure}[h] \\centering \\includegraphics[width=8.0cm , angle=90]{girard_fig5.eps} \\caption{The [$\\alpha$/Fe] ratio from TGMET{\\Large$\\alpha$} vs. the same ratio from the reference catalogue of abundances.} \\label{figure_alf} \\end{figure}"
    },
    "0601/astro-ph0601191_arXiv.txt": {
        "abstract": "The recently proposed Cooperstock-Tieu galaxy model claims to explain the flat rotation curves without dark matter.  The purpose of this note is to show that this model is internally inconsistent and thus cannot be considered a valid solution.  Moreover, by making the solution consistent the ability to explain the flat rotation curves is lost. ",
        "introduction": "Cooperstock and Tieu model a galaxy in general relativity as an axially symmetric, pressure free dust cloud with metric \\begin{equation} \\label{metric} ds^2=-e^{w}(cdt-Nd\\varphi)^2+r^2e^{-w}d\\varphi^2+e^{v-w}(dr^2+dz^2) \\end{equation} where $w$, $v$, and $N$ are functions only of $r$ and $z$, thus the metric is stationary \\citep{CT}.  They further assume their coordinates to be co-moving with the galactic dust, thus \\begin{equation} u^\\mu=e^{-w/2}\\delta_t^\\mu \\end{equation} is the four-velocity.  By performing a local diagonalization of the metric they obtain a relation between the metric and local angular velocity as \\begin{equation} \\omega=\\frac{cNe^w}{r^2e^{-w}-N^2e^w}\\approx\\frac{cN}{r^2}\\label{rot} \\end{equation} where the approximate value is appropriate to the order they consider.  The $tt$ and $t\\varphi$ Einstein equations are respectively \\begin{mathletters} \\begin{eqnarray} r^{-2}\\left({N_{,r}^2+N_{,z}^2}\\right) &=& \\frac{8\\pi G \\rho}{c^2}\\label{dens}\\\\ N_{,rr}-\\frac{1}{r}N_{,r}+N_{,zz}&=&0. \\label{neqn} \\end{eqnarray} \\end{mathletters} The second equation is equivalent to the formula \\begin{equation} \\label{lap} \\nabla^2\\Phi=0 \\end{equation} where they have defined\\footnote{It should be noted that this $\\Phi$ is not the Newtonian gravitational potential.} \\begin{equation} \\Phi=\\int\\frac{N}{r}dr. \\end{equation}  Thus the observed rotation curve becomes a boundary condition for the solution to Laplace's equation (\\ref{lap}) which they take in the form \\begin{equation} \\Phi=\\sum_nC_ne^{-k_n|z|}J_0(k_nr) \\end{equation} where the $k_n$ are chosen for orthogonality over the radius of the galaxy.  Once $N$ is found by fitting this function to the obsersved rotation curve, they derive the density by (\\ref{dens}) and in this way they obtain an excellent fit to the data while obtaining a density profile that accords with observation.  However, it has been pointed out \\citep{c1,c2} that, since the solution depends on $|z|$, equation (\\ref{neqn}) is not satisfied, but rather yields a singular contribution to the $z=0$ plane, which has the properties of an exotic form of matter.  It may be wondered whether this singular disk can be removed by choosing a different solution form or by increasing the complexity of the model.  However, in the following analysis we will show that this is not possible and that model fails to accord with general relativity. ",
        "conclusions": "It has been shown that the Cooperstock-Tieu galaxy model is inconsistent as a general relativistic model and that a proper model fails to account for the flatness of the rotation curves without the dark-matter hypothesis.  This failure is due to the weakness of the metric coupling to the angular momentum of the galaxy.  However, the flat rotation curves seem to imply a large inertial induction effect, where the rotating inner matter boosts the rotation of the outer matter, leveling off the rotation curve, which is what the Cooperstock-Tieu model attempts to describe within general relativity.  Since their solution predicts a matter density well within visible limits it is quite possible that their solution represents an alternative, more Machian, gravitational theory where inertial induction effects are much larger than in General Relativity."
    },
    "0601/astro-ph0601158_arXiv.txt": {
        "abstract": "Cosmic rays represent one of the most fascinating research themes in modern astronomy and physics. After almost a century since their discovery, a huge amount of scientific literature has been written on this topic and it is not always easy to extract from it the necessary information for somebody who approaches the subject for the first time. This has been the main motivation  for preparing this article, which is a concise and self-contained review for whoever is interested in studying cosmic rays. The priority has been given here to well established facts, which are not at risk to get obsolete in a few years due to the fast progress of the research in this field. Also many data are presented, which are useful to characterize the doses of ionizing radiation delivered to organisms living on the Earth due to cosmic rays. The technical terms which are often encountered in the scientific literature are explained in a separate appendix. ",
        "introduction": "Cosmic rays (CR) represent a fascinating subject of research, which is recently witnessing a growing interest within the scientific community. With the next generation of cosmic ray detectors, like for example the Pierre Auger Observatory which is currently under construction, there is the hope that the questions posed by these cosmic particles will find soon some satisfactory answer. One should also note that the attention to CR is not only restricted to the traditional fields of high energy physics and astroparticles. This article is for instance the result of the authors' efforts to investigate the mutagenic effects of cosmic rays on cells. To this purpose, one needs to characterize very precisely the fluxes and intensities of particles arriving to the ground as an effect of the interaction of CR with the atmosphere of the Earth.  The increasing popularity of CR is accompanied by an increasing demand of information on this topic, but it is not always easy to extract this information from the huge amount of literature which has been written on CR during a period of almost a century after their discovery. An additional problem is that, looking at the scientific literature on CR, often concepts like the integral vertical intensity or the differential integrated flux are encountered. These terms sound somewhat puzzling and exotic for a reader that encounters them for the first time. These motivations have compelled us to prepare this short review, which would like to be a compact but self-contained reference on CR. The preference has been given here to well established facts, which are not in danger to become obsolete in a few years due to the fast progress in this topic. Moreover, a considerable effort has been done in order to explain in details and with the help of figures the technical terms used in the current scientific literature. The puzzles raised by the existence of  ultra-high energy cosmic rays with energy of 10$^{19}$eV or higher, like the mystery of their origins or the apparent violation of the theory of relativity which is connected with them, have been just briefly mentioned. Hopefully, these puzzles will be solved with the next generation of cosmic rays detectors. Another aim  of this work is to present data concerning the physical parameters, e.~g. types of radiation, delivered effective doses and dose rates, fluxes and intensities of incoming  particles, which characterize the radiation to which the organisms living on the Earth are exposed because of CR. These data are certainly of interest for scientists working in life sciences. Cosmic rays are in fact the source of an almost uniform background of ionizing radiation which is present everywhere on the Earth. Most of their energy arrives to the ground in the form of kinetic energy of muons. The latter particles are very penetrating and are able to travel for kilometers in water and for hundreds of meters in rock. Since ionizing radiation is mutagenic, it is very likely that the radiation of cosmic origin has shaped in some way the evolution of life on our planet, generating some kind of adaptive response in cells.  Finally, as already mentioned, this article is self-contained, but of course it is far from being complete, because it is impossible to cover all the immense literature which exists on cosmic rays. To integrate the material presented here, suggested further readings are for instance \\cite{mewaldt,pdg,biermann,cronin,battistoni,Anchor} and references therein.   \\section[Cosmic rays and...]{Cosmic rays and the natural background radiation} \\subsection{Introduction to cosmic rays} Cosmic rays, first discovered by Victor Hess in 1912, are charged particles accelerated at very high energies by astrophysical sources located anywhere beyond the atmosphere of the Earth. $89$\\% of cosmic rays consists of protons, followed by $\\alpha-$particles ($\\sim 10$\\%) and heavier nuclei ($\\sim 1$\\%) \\footnote{ These data are taken from Ref.~\\cite{mewaldt}. Of course, as it always happens in the case of experimental data, there is some uncertainty due to measurement errors. For this reason, different authors report slightly different values as in the case of \\cite{obrien}, which gives $95$\\% for protons, $3.5$\\% for $\\alpha$ particles and $1.5$\\% for all the rest. It should also be noted that the flux of particles below the energy of $10$GeV strongly depends on the 11--year solar cycle.}. All elements of the periodic table, including  transuranic elements, have been detected. More details on the composition of cosmic rays can be found in Ref.~\\cite{pdg} and references therein.  Within the flux of incoming particles one can distinguish different components. The most relevant ones are galactic CR (GCR), solar energetic particles (SCR) and the anomalous CR (ACR) \\cite{nasaacr}. GCR originate from sources located in our galaxy outside the solar system. It is believed that GCR are a consequence of astrophysical events like stellar flares, stellar coronal mass ejections, supernova explosions, particle acceleration by pulsars \\cite{obrien}. Very often, the word cosmic rays refers only to GCR \\cite{mewaldt}. The smallest detectable energy of GCR is about 1~GeV. Below this energy, the screening effect of the solar wind is too strong to allow them to penetrate the heliosphere \\cite{gapnote}. A detailed information on the interaction of GCR with the magnetic fields of the Sun and of the Earth is contained in Ref.~\\cite{obrien} and references therein. Here we would just like  to mention the action on CR of the galactic magnetic fields which are generated by the spinning of the Milky Way. These fields are relatively weak, because the average magnetic field in our galaxy is of the order of $10^{-10}$T. However, since CR consist of charged particles traveling along huge distances, at the end these magnetic fields are able to bend their trajectories in a relevant way. At this point one should  note that there are two different components in the magnetic fields of our galaxy, a regular one and a turbulent one, see Ref.~\\cite{RandKul}. The strengths and directions of the magnetic fields belonging to the turbulent component are random \\cite{gapnote,RandKul}. Due to this randomness, the trajectories of CR are randomly bent. As a consequence, the flux of CR  which arrives on the Earth is also random or, more precisely, isotropic. For this reason, it is not easy to ascertain where GCR are coming from. Besides being isotropic, the particle flux is constant in time too, so that GCR form an almost uniform background of ionizing radiation on the surface of the Earth. CR of energies up to  $10^{21}$eV have been observed. These are considerable energies for a microscopic particle. For example, the upper energy limit of $10^{21}$eV corresponds in SI units approximately to 160 Joules. This is comparable to the kinetic energy of a ball of $0.8$kg thrown at the speed of $50$km/hour. The origin of such ultra high energy CR (UHECR) is so far unknown, but there are strong hints that they could be produced outside of our galaxy. Candidate sources of UHECR could be relativistic plasma jets from supermassive black holes \\cite{lofar}, explosions of galactic nuclei \\cite{obrien}, but other possibilities have been proposed, like magnetic monopoles, see for example \\cite{exotic}. The main evidence which suggests that UHECR are of extragalactic origin is provided by the fact that the magnetic fields which are present in the Milky Way are not able to trap CR of that energy. Indeed, already protons of energy higher than 10$^{15}$eV are able to escape the galactic confinement\\footnote{ In the case of heavier nuclei, the threshold energy for escaping the galactic confinement is higher than that of protons. It is for this reason that one observes a greater proportion of heavier nuclei with respect to protons in CR with energy above 10$^{15}$eV.}. As a consequence, if protons of energies of 10$^{19}$eV or higher would be produced by sources located in our galaxy, they would escape it in all possible directions following trajectories which are almost straight lines. For this reason, the ultra high energy protons which reach the Earth should arrive along directions which are approximately parallel to the galactic plane. However, this conclusion is not confirmed by observations, see Refs.~\\cite{cseven,unscearb} and references therein. Observations  show in fact that the spatial distribution of ultra high energy protons is isotropic, so that their directions are not aligned with the galactic plane. This is of course a strong hint that UHECR are of extragalactic origin.  On the other side, there is at least one argument which seems to point out that the sources of UHECR are not very far from our galaxy. In fact, it has been noted that protons of energies above 5$\\cdot$10$^{19}$eV would lose their energy by interacting with the photons of the microwave (big bang) background. This effect was predicted in 1966, one year after the detection of the microwave background radiation, by Kenneth Greisen, Vadem Kuzmin and Georgi Zatsepin. The energy threshold of 5$\\cdot$10$^{19}$eV is called the GZK-limit, from the names of its discoverers. Protons with energies above that threshold will be slowed down during their travel to the Earth by the mechanism of energy-loss pointed out by Greisen, Kuzmin and Zatsepin until their energy falls below the GZK-limit. This mechanism is so efficient that, in practice, protons with energies higher than 5$\\cdot$10$^{19}$eV should not be observed on the Earth if their source is located at distances which are greater than 50Mpc~\\footnote{Mpc stands for megaparsec. 50Mpc are approximately 150 millions of light years. This is about 1500 times the diameter of a galaxy and it is not a big distance in comparison with the cosmic scale of distances.}. Since ultra high energy protons have instead been detected, this implies that they originate from sources which are within the range of 50Mpc. Yet the known cosmic objects which could be able to accelerate protons to such high energies are at least at a distance of about 100 Mpc or more. So cosmic ray protons above this energy should not arrive on the Earth or we should explain why their formidable sources, which should be relatively near to us, remain invisible. These contradictions are known under the name of GZK paradox. On these points see Ref.~\\cite{noelsta}.  Let's now discuss the remaining components of CR. Energetic solar events, like for instance solar flares, are able to accelerate particles up to some GeV very efficiently within the time of 10 seconds. The SCR are mainly protons, heavier nuclei and electrons. Finally, it is worth  mentioning also the case of ACR. This kind of CR is mainly characterized by ions of elements which are difficult to ionize, including He, N, O, Ne and Ar. Moreover, ACR have a relatively low energy, up to a few hundreds of MeV \\cite{klecker}. It is thus improbable that CR of such a low energy could originate  from the violent phenomena which produce GCR. Besides, we recall that CR of energy below 1GeV coming from outside the heliosphere cannot penetrate very deeply inside the solar system, because they are deflected by the solar magnetic fields.  There are indeed evidences that  ACR may be neutrally charged dust particles present in the interstellar gas near the border of the solar system \\cite{mewaldt,nasaacr,jokipii}. When these particles enter in contact with the far edges of the heliosphere, they are ionized by UV solar photons or by interactions with the particles within the solar wind. Once these dust particles have been ionized, they are accelerated by the shock waves formed when the solar wind encounters the interstellar plasma. Particles escaping the shocks may diffuse toward the inner heliosphere and arrive on the Earth as ACR. Actually, there are many open questions on the mechanism of ACR acceleration. Hopefully, these questions will be answered in 2007. By that date, in fact, the Voyager~1 should reach the region in which the shocks occur, which is thought to be somewhere between 75 and 100 AU from the Sun \\footnote{At the time in which this article was finished, a distance up to about 90 AU has been explored by Voyager 1. Some of the results of the observations can be found in Refs.~\\cite{voy1,voy2,voy3}.}. This will be the first time that an example of CR acceleration will be observed directly.  \\subsection{Interaction of CR with the Earth's atmosphere}\\label{subsec:intatm} When CR arrive near the Earth, they hit the nuclei of the atoms of the atmosphere, in particular nitrogen and oxygen, producing in this way secondary particles.  The first interaction of the CR primary particle takes place in the top 10\\% of the atmosphere \\cite{clay}. The most relevant reactions\\footnote{One should remember that the collisions of CR with the atmosphere gives raise also to less relevant reactions, which produce  particles like  kaons, $\\eta$ particles and even resonances.},  remembering that 90\\% or more of CR consists of protons, are: \\begin{eqnarray} pp\\longrightarrow pn\\pi^+\\qquad &\\mbox{or}&\\qquad pp\\longrightarrow pp\\pi^0\\label{reaserone}\\\\ pn\\longrightarrow pp\\pi^-\\qquad& \\mbox{or}\\qquad pn\\longrightarrow pn\\pi^0\\qquad \\mbox{or}& \\qquad pn\\longrightarrow nn\\pi^+\\label{reasertwo}\\\\ \\end{eqnarray} In the above reactions all the secondary particles are hadrons, namely protons $p$, neutrons $n$ and pions in all their charged states $\\pi^\\pm,\\pi^0$. Pions may in turn decay according to the following processes: \\begin{eqnarray} \\pi^+\\longrightarrow \\mu^+\\nu_\\mu&\\mbox{and}& \\pi^-\\longrightarrow\\mu^-\\bar \\nu_\\mu\\label{decaserone}\\\\ \\pi^0&\\longrightarrow&\\gamma\\gamma\\label{decasertwo} \\end{eqnarray} where the $\\mu^\\pm$'s are muons, the $\\gamma$'s are photons and $\\nu_\\mu,\\bar \\nu_\\mu $ are respectively muonic neutrinos and their anti-particles. The mean life of pions is 26ns for $\\pi^\\pm$ and 10$^{-16}$ seconds in the case of $\\pi^0$. For this reason, charged pions may still collide with air atoms before decaying, but it is very unlikely that this happens in the case of neutral pions, which have a very short average life. Other secondary particles, like protons, neutrons  and photons interact very frequently with the atoms of the atmosphere giving raise in this way to a cascade of less and less energetic secondary particles. At the end, these particles are  stopped by the atmosphere or, if the energy of the primary particle was sufficiently high, they can reach the ground.      The main mechanism of energy loss \\footnote{Here and in the following energy means the kinetic energy of the particles.} of high energetic hadrons consists in the disintegration of the molecules of the atmosphere \\cite{milagro}, see Fig~\\ref{shower}.  This leads to the creation of new particles through nuclear interactions like those shown in Eqs.~\\ceq{reaserone} and \\ceq{reasertwo}. At lower energies, dissipative processes become predominant, in which the molecules of the atmosphere get either ionized or excited. The most relevant process of this kind in the case of heavy charged particles is the ionization of the molecules of the atmosphere. Lighter charged particles like electrons and positrons lose their energies not only by ionization, but also by bremsstrahlung. This consists in the radiative loss of energy of charged particles moving inside matter, when they are deflected by the electrostatic forces of the positive charged nuclei of the surrounding molecules. Photons and neutrons, the remaining  relevant particles in the cascade, are examples of indirectly ionizing radiation. Their interaction mechanisms are more complicated than those of charged particles and will not be described here. The interested reader may find a more detailed account on the way in which radiation of different kinds interacts with matter in Ref.~\\cite{doereport}.  The total number of secondary particles $N_{sec}$ within the cascade grows rapidly, mainly sustained by the processes of bremsstrahlung and pair production due to electrons, positrons and photons. Hadrons like  protons and, to a less extent, neutrons, are easily stopped by the atmosphere, so that they increase relevantly the number of particles by disintegrating the molecules of air only during the first stages of the formation of the cascade. The decays of pions given in Eqs.~\\ceq{decaserone} and \\ceq{decasertwo} produce  muons and photons of considerable energies. The muons are very penetrating particles and do not interact very much with the air. They lose a small fraction of their energy before reaching the ground by ionizing the molecules of the atmosphere. The photons give raise instead to electron-positron pairs $e^+e^-$. In turn, electrons and positrons create other electrons by ionization or other photons due to bremsstrahlung. In this way, while the cascade propagates inside the atmosphere, the number of its electrons, positrons and photons grows almost exponentially. The maximum number of  particles inside the cascade is attained when the average energy per electron reaches the threshold $E_T\\sim$80MeV. When the energy of electrons in air falls below that threshold, ionization starts to prevail over bremsstrahlung as the main mechanism of energy loss of electrons in air and the process for increasing the number of particles described above ceases to be effective.    If the energy of the primary particle is below $10^{14}$eV, essentially only the penetrating muons and neutrinos are able to arrive to the sea level, while the other particles in the cascade are absorbed at higher altitudes. Actually, neutrinos interact so rarely with matter that they could pass through a light year of water without undergoing any interaction. Thus, if one is concerned with the dose of ionizing radiation delivered by CR to the population, the contribution of neutrinos  can simply be neglected. Muons are more dangerous for the health. They have a short mean life at rest (2.2ns), but since they travel at very high speeds, they manage to reach the surface of the Earth due to the relativistic dilatation of time. Part of these muons can still decay, giving raise to electrons $e^-$ or positrons $e^+$, mainly according to the process: $\\mu^\\pm=e^\\pm+\\nu_e+\\nu_\\mu$.  When the energy of the primary particles is instead above $10^{14}$eV, the cascade of secondary particles arrives to the ground before being stopped by the atmosphere. In that case, the cascade is called with the name of air shower, see Fig.\\ref{shower}. To be precise, the effects of an air shower started by a primary particle of 10$^{14}$eV are relevant up to altitudes comparable to that of Mount Everest \\cite{allan,lofara}. \\begin{figure}[bpht] \\centering \\includegraphics[width=\\textwidth]{shower.eps} \\caption{This figure illustrates schematically how air showers generate from cosmic rays. The high energetic  primary particle, usually a proton, whose trajectory has been denoted in black, starts to interact with the molecules of the upper atmosphere. In this way, secondary particles are produced, which give raise to other particles (tertiary, quaternary etc.) via other interactions with the atmosphere or via decay processes. The total flux of particles can be divided in a {\\it electromagnetic component} (photons, electrons and anti-electrons or positrons), see the blue trajectories in the figure, in a {\\it muon component}, see the cyan and red trajectories, and finally in a {\\it nucleonic component} (mainly protons, neutrons, rarely pions) denoted in brown. At the sea level, the air shower has the form of a pancake, whose height is around two meters, while its radius usually around a few hundreds of meters, but may reach some tents of kilometers in the case of very energetic cosmic rays. The nucleonic component is usually confined in a narrow cone centered along the direction of the incoming primary particle. }\\label{shower} \\end{figure} Only air showers generated by primaries of energy of about $10^{15}$eV or higher are able to reach also the typical altitudes of inhabited areas and arrive to the sea level. The frequency of these air showers is relatively high, because the  total flux of primary particles with energy $E\\ge 10^{15}$eV is of about 100 particles per m$^2$ per year. Giant air showers produced by primaries of energies beyond $10^{20}$eV are much more rare: their total  flux is of 1 particle per km$^2$ per century \\cite{mewaldt}. More data concerning fluxes of incoming primary particles and an explanation of how these data are measured using ground detectors, can be found in Refs.~\\cite{Anchor,eheas}.  In the air shower we distinguish a nucleonic component, a muon component and an electromagnetic component. The nucleonic component is generated by high energetic protons and neutrons, which  disintegrate the atoms of the air giving raise to other protons and neutrons. The fluxes of electrons, positrons and photons started  by the decay of the $\\pi^0$'s,  together with the electrons and positrons coming from the decay of muons or from the ionization due to the hadrons, form the electromagnetic component. In the early times of CR research, electrons and positrons have been called the soft component, while muons coming from the decay of the charged pions have been called the hard component. These names originate from the fact that muons are very penetrating and thus they may be regarded as ``hard'' particles. For instance, at the energy of 1 GeV, the range\\footnote{The range is defined as the average depth of penetration of a charged particle into a material before it loses all its kinetic energy and stops. The concept of range has a meaning only in the case of charged particles whose energy is kinetic energy which is lost continuously along their paths due to ionization and bremsstrahlung processes \\cite{doereport}.} of muons is  2.45$\\cdot$10$^{5}$g~cm$^{-2}$. This implies that in water, which has a density of 1~g~cm$^{-3}$ at 4~C, muons run along an average distance of 2.45 km before being stopped completely. In standard rock, which has a density of 2.65~g~cm$^{-3}$ \\cite{pdg}, this average distance reduces to about 900 meters.   At the sea level, the air shower has approximately the form of a pancake with an height of 1-2  meters. Its extension in the other two directions, defined as the distance in which 90\\% of the total energy of the shower is contained, is given by the so-called Moli\\`ere radius. For example, in the case of an air shower started by a primary particle of an energy of $10^{19}$eV (10 EeV) the Moli\\`ere radius  has a length of about 70 meters. The real extension of the shower is  much bigger and some of the muons may be detected up to a distance of a few kilometers from the core \\cite{eheas}. Usually the nucleonic component, which is composed by heavier particles than those of the muon and electromagnetic components, is less deflected from the direction of  the incident primary particle by the interactions with the atmosphere  and it is thus concentrated in a narrow cone inside the air shower. The center of the cone is roughly aligned with the direction of the original primary particle. The number of secondary particles which arrives on the ground with an air shower is huge. Considering particles whose energies are greater that 200keV, an air shower generated by a 10EeV primary particle contains up to 10$^{10}$ particles, mostly photons, electrons and positrons. Electrons outnumber positrons in the ratio 6 to 1. The maximum number of particles, i. e. the so-called point of shower maximum or simply shower maximum, is attained at an altitude of 2--3 km above the sea level. Many other data and diagrams describing the propagation of air showers in the atmosphere may be found in Ref.~\\cite{TalkPierog}.  When air showers approach the ground, about 85\\% of their energy is concentrated in the electromagnetic component. The contribution of the muon and nucleonic components is thus much less relevant. The situation changes completely if we consider all CR, and not only those which have sufficient energy to give raise to an air shower. As we see in Fig.~\\ref{percent}, muons are in fact responsible for about 85\\% of the total equivalent dose delivered by CR to the population at the sea level. As a consequence, globally it is the muon component the most significant from the energetic point of view and not the electromagnetic component. The reason of this fact is that primary particles with energy $E\\ge 10^{15}$, namely those which can produce air showers, form just a minimal fraction of the total amount of CR arriving on the Earth. For example, the total flux of particles with energy $E\\ge 10^{12}$eV (1 TeV), is of 1~particle~per~m$^2$~per~second, i. e. a factor of 3$\\cdot10^5$ higher than the  total flux of CR with energy $E\\ge 10^{15}$ reported above. In other words, there is an overwhelming number of CR with energy lower than $10^{15}$eV which are not able to start an air shower, but may still generate energetic muons. Being very penetrating particles, these muons are not stopped easily by the atmosphere and reach the sea level altitude, where they represent the biggest source of ionizing radiation of cosmic origin. Other particles which deliver relevant doses of ionizing radiation to the population on the surface of the Earth are photons, electrons and neutrons.  The percent contributions to the total equivalent dose of the various  components of CR as a function of the altitude is given in Fig.~\\ref{percent}. In that figure, which has been published in 1996, the curve concerning neutrons should be taken with some care, because the data on neutron fluxes in the atmosphere were still sparse at that time \\cite{unscearb}. Other data about energies and fluxes of particles due to CR will be given in the next Section. \\begin{figure}[bpht] \\centering \\includegraphics[width=10cm]{percent.eps} \\caption{Percent contribution of the various CR components to the total equivalent dose at different altitudes. This figure is based on data of Ref.~\\cite{obrien}.}\\label{percent} \\end{figure} We note that protons and neutrons prevail at higher altitudes, but they are rapidly absorbed by the atmosphere and muons become dominant at lower altitudes.   \\subsection{Intensities and fluxes of CR} CR are the source of an avalanche of secondary particles which hit constantly the surface of the Earth. To determine the energy and number of these particles, their directions of arrival and their distribution in time, quantities like the integral vertical intensity or the differential directional intensity are measured. The meaning of these quantities is explained in details in a separate Appendix at the end of this article. In this section, we present some experimental data which are useful in order to characterize the contribution of the muon, electromagnetic and nucleonic components to the background of ionizing radiations on the ground due to CR.      The  integral vertical intensity (IVI in the Appendix, see Eq.~\\ceq{ividef}) of the muon component with energy above 1 GeV at sea level is approximately $I_{ivi}^{hard}(\\theta=0)\\sim$0.70$\\cdot 10^{-2}$ cm$^{-2}$s$^{-1}$sr$^{-1}$\\cite{pdg}. The integral directional intensity of muons in the other directions, which are at an angle $\\theta$ with respect to the vertical direction, has the following behavior: $I_{idi}^{hard}(\\theta)\\propto I_{ivi}^{hard}(\\theta=0)\\cos^2\\theta $. More complete phenomenological formulas for the angular distribution of CR intensities may be found in Refs.~\\cite{pdg,allkofer,dar}. As Fig.~\\ref{ziegler} shows,  muons arriving to the ground are very energetic. The most frequent muon energy is 500MeV, while the average muon energy is 4GeV. \\begin{figure}[bpht] \\centering \\includegraphics[width=10cm]{ziglermio.eps} \\caption{Differential integrated flux (see Eq.~\\ceq{fluangthe} of the Appendix) of the different components of CR related radiation at the sea level. This figure has been created on the basis of an analogous figure appeared in Ref.~\\cite{schumacher}. The original data are taken from Ref.~\\cite{ziegler}.}\\label{ziegler} \\end{figure} There is almost no protection from this source of radiation, since, as we have seen before, high energetic muons are able to penetrate thick layers of concrete and rocks.  The integral vertical intensity for electrons plus positrons with energies greater than 10, 100 and 1000 MeV is very approximately given by\\cite{pdg}: \\begin{eqnarray} I_{ivi}^{el+pos}(\\theta=0, E_{min}>10MeV)&\\sim& 0.30\\cdot 10^{-2} \\mbox{cm}^{-2}\\mbox{s}^{-1}\\mbox{sr}^{-1}\\\\ I_{ivi}^{el+pos}(\\theta=0, E_{min}>100MeV)&\\sim& 0.06\\cdot 10^{-2} \\mbox{cm}^{-2}\\mbox{s}^{-1}\\mbox{sr}^{-1}\\\\ I_{ivi}^{el+pos}(\\theta=0, E_{min}>1000MeV)&\\sim& 0.02\\cdot 10^{-3} \\mbox{cm}^{-2}\\mbox{s}^{-1}\\mbox{sr}^{-1} \\end{eqnarray} Moreover, the total flux of electrons plus positrons amounts approximately to 30{\\tt \\char`\\%} of the total particle flux which is reaching the ground due to cosmic rays. Always according to \\cite{pdg}, the ratios of photons to electrons + positrons is approximately 1.3 above 1 GeV and 1.7 below the critical energy $E_T\\sim 80$MeV mentioned above. The differential flux on the ground of the various components of radiation related to CR can be seen in Fig.~\\ref{ziegler}. One may observe that different particles become predominant at different energies. In the lowest portion of the energy range, electrons are predominant. At energies of around 1~MeV electrons are taken over by fast  neutrons which arise mainly due to the de-excitation of atmospheric nuclei following compound-nucleus reactions \\cite{desalin} \\footnote{Roughly speaking, in a compound nucleus reaction a neutron or a proton, but also an $\\alpha-$particle interacting with the nucleus of an atom creates a nucleus of higher atomic number which is metastable and decays after a short period of time. For example, a possible compound-nucleus reaction is: $p+{}^{63}Cu\\longrightarrow{}^{64}Zn^*$. The de-excitation of the metastable nucleus  produces others neutrons and protons, e. g.: ${}^{64}Zn^*\\longrightarrow {}^{63}Zn+n$ or ${}^{64}Zn^*\\longrightarrow {}^{62}Cu+n+p $. Compound-nucleus reactions become possible only if the energies of the incoming nucleons or alpha particles are such that the de Broglie wave lengths of these particles are comparable with the size of the hit nucleus.}. Electrons become once again predominant in the energy range going from a few MeV up to some tens of MeV. Starting from energies approximately above 200MeV, the number of muons becomes overwhelmingly high in comparison to that of the other particles. In considering the above data, one should of course take into account the fact that, at low energies, let's say below 10MeV, particle fluxes are strongly dependent on many factors, including weather conditions, so that there are big uncertainties in their measurement up to an order of magnitude \\cite{ziegler}. Moreover, all the  data presented so far in this Section refer to the sea level altitude. With increasing altitudes, the contribution to the particle flux given by protons and neutrons, the nucleonic component of CR, becomes more and more relevant (see Fig.~\\ref{percent}) and is predominant above atmospheric depths   of approximately 500~g~cm$^{-2}$~\\footnote{The atmospheric depth is a quantity which is often used to measure the altitude. For convenience, we give in Fig.~{\\ref{altdep}} a diagram which is useful to make the conversion from atmospheric depth in g/cm$^2$ to altitude in kilometers. }. \\begin{figure}[bpht] \\centering \\includegraphics[width=10cm]{altitvsdepth.eps} \\caption{Conversion diagram from atmospheric depth to altitude. The data corresponding to the dots are taken from Ref.~\\cite{usat}.}\\label{altdep} \\end{figure}  To conclude this Section, we provide also some data regarding the particle flux outside the heliosphere.  The total flux of CR in the galaxy is large, about 100000 particles m$^{-2}$s$^{-1}$ \\cite{zieglerb}. Much lower is for example the integral directional flux of CR primaries with $E>2\\cdot 10^{15}$eV is about $\\Phi_{idf}=75000$ particles km$^{-2}$sr$^{-1}$day$^{-1}$. This datum has been derived using a formula for the integral directional flux given in Ref.~\\cite{lofara}, which is based on the results of measurements and simulations of incoming cosmic ray fluxes reported in Ref.~\\cite{fowler}. The dependence of this flux on the direction of the incoming particles is minimal since, as mentioned before, the distribution of these particles is isotropic due to the presence of random  magnetic fields in the Milky Way. Finally, the integral directional flux of particles with energies above $10^{20}$eV is $\\Phi_{idf}=1 \\mbox{ particle} \\mbox{ km}^{-2}\\mbox{ sr}^{-1}\\mbox{ century}^{-1}$ \\cite{kitpmini}. With such a small flux, the investigation of CR with energy beyond the GZK limit requires detectors which cover a very large area. Perhaps with the next generation of detectors it will be possible to solve the puzzles connected with UHECR.   \\subsection{Dose of ionizing radiation from CR in present times} In present times\\footnote{The data concerning the present levels of radiation coming from CR are taken from the United Nations Report of the year 2000 \\cite{unscearb}. } the effective dose rate (EDR) delivered by CR to the human population varies from a minimum of 300$\\mu$Sv~year$^{-1}$ to a maximum of 2000$\\mu$Sv~year$^{-1}$. This wide range depends on many factors, the most important one being the altitude. At sea level, the value which usually is given for the EDR is 270~$\\mu$Sv~year$^{-1}$, or equivalently 31~nSv~h$^{-1}$. One should keep in mind however that this estimation is the result of population-weighing average. As a matter of fact, even if the altitude is fixed at the sea level, the EDR still changes with the latitude within a range of variation of approximately 10\\%. Strictly speaking, 270~$\\mu$Sv~year$^{-1}$ is the dose rate received by the population living at a latitude which is near the 30$^\\circ$ parallel. It turns out in fact from the distribution of the population on the Earth that this is the average latitude at which people are living. Besides, the above value of EDR takes into account only the contributions of  muons and of the electromagnetic component. The nucleonic component, which at the sea level is essentially consisting of neutrons, gives to the average EDR at the sea level an additional contribution of 48~$\\mu$Sv~year$^{-1}$ or, equivalently, 5.5~nSv~h$^{-1}$. The figures concerning neutrons should be taken with some care, since up to the time in which Ref.~\\cite{unscearb} was released, the available data on neutron fluxes were sparse. If one considers also the different altitudes in which the population lives, the population-weighted EDR is of 380~$\\mu$Sv~year$^{-1}$, corresponding to an average habitant of the planet Earth living approximately near the 30$^\\circ$ parallel and at an altitude of 900 meters above the sea level. ",
        "conclusions": "In this review a short but thorough account has been provided on what it is known about cosmic rays, starting from their origin in space and arriving to the doses of ionizing radiation delivered by them to the human population. Only the sources of CR have not been discussed here, because this argument is outside the aims of this article. Much attention has been dedicated to fluxes and intensities of CR and of the particles arriving on the ground as an effect of the cascades  initiated by CR in the atmosphere. These quantities are  of interest for scientists working in different subjects. Apart from research in high energy physics and astronomy, the fluxes and intensities of particles of cosmic origin are also studied  for radioprotection purposes \\cite{obrien,unscearb} and for their capability of causing potentially harmful failures in computers and electronic storage devices \\cite{zieglerb}. Fluxes of CR  are also carefully measured due to their relevance to space explorations, see for example \\cite{ssareport}. Finally, in  Appendix A the various kinds of fluxes and intensities of CR and related particles which one encounters in research articles about CR  have been defined and their meaning has been illustrated. Concrete expressions for these quantities have been given in terms of mass densities, velocity distributions  and energy densities. Both relativistic and non-relativistic cases have been treated. Up to now, a systematic classification and explanation of these quantities such as that provided in this work was missing in the scientific literature on CR. The necessity of filling this gap justifies the length of this Appendix. \\begin{appendix}"
    },
    "0601/astro-ph0601644_arXiv.txt": {
        "abstract": "{We report on the results of a deep 1.6 Ms INTEGRAL observation of the Cassiopeia region performed from December 2003 to February 2004. Eleven sources were detected with the imager IBIS-ISGRI at energies above 20 keV, including three new hard X-ray sources. Most remarkable is the discovery of hard X-ray emission from the anomalous X-ray pulsar 4U~0142+61, which shows emission up to $\\sim$150 keV with a very hard power-law spectrum with photon index $\\Gamma = 0.73 \\pm 0.17$. We derived flux upper limits for energies between 0.75~MeV and 30~MeV using archival data from the Compton telescope COMPTEL. In order to reconcile the very hard spectrum of 4U~0142+61 measured by INTEGRAL with the COMPTEL upper limits, the spectrum has to bend or break between $\\sim$75 keV and $\\sim$750 keV.  1E~2259+586, another anomalous X-ray pulsar in this region, was not detected. INTEGRAL and COMPTEL upper limits are provided. The new INTEGRAL sources are IGR~J00370+6122 and IGR~J00234+6144. IGR~J00370+6122 is a new supergiant X-ray binary with an orbital period of $15.665 \\pm 0.006$ days, derived from RXTE All-Sky Monitor data. Archival BeppoSAX Wide-Field Camera data yielded four more detections. IGR~J00234+6144 still requires a proper identification. Other sources for which INTEGRAL results are presented are high-mass X-ray binaries 2S~0114+650, $\\gamma$~Cas, RX~J0146.9+6121 and 4U~2206+54, intermediate polar V709~Cas and 1ES~0033+595, an AGN of the BL-Lac type. For each of these sources the hard X-ray spectra are fitted with different models and compared with earlier published results. ",
        "introduction": "\\label{sec:intro}  The {\\em INTErnational Gamma-Ray Astrophysics Laboratory} INTEGRAL \\citep{Winkler03_IGR} is ESA's currently operational space-based hard X-ray/soft gamma-ray telescope. Since its launch in October 2002 it has been observing the sky at photon energies between 3~keV and 2~MeV. Aboard INTEGRAL there are two main hard X-ray/soft gamma-ray instruments, the imager IBIS \\citep{Ubertini03_IBIS} and the spectrometer SPI \\citep{Vedrenne03_SPI}, supplemented by two X-ray monitors, JEM-X \\citep{Lund03_JMX} and an optical monitor, OMC \\citep{Mas-Hesse03_OMC}. These instruments are all wide-field instruments with observational coverage over a broad spectral band. In this work only INTEGRAL results from the low-energy detector of IBIS, called ISGRI \\citep{Lebrun03_ISGRI}, will be presented.  IBIS-ISGRI is a coded-mask instrument with arcminute-resolution \\citep[12\\arcmin \\, FWHM;][]{Gros03_IBIS} imaging capabilities in the energy range from $\\sim$20~keV to $\\sim$300~keV. Its field of view (FOV) is $29^\\circ \\times 29^\\circ$ (full width zero response). The most sensitive part of the FOV of a coded-mask instrument is the fully-coded FOV, which is $9^\\circ \\times 9^\\circ$ for IBIS-ISGRI \\citep[see][ for detailed information on coded-mask instruments]{IntZand92_codedmask}.   The main targets of our INTEGRAL observations are the young Galactic supernova remnants Cassiopeia~A and Tycho. In this paper these sources will not be addressed in detail; we refer to \\citet{Vink05_casa} and \\citet{Renaud04_tycho} for more detailed studies. This paper presents results on compact hard X-ray sources detected in the deep 1.6-Ms observation of the Cassiopeia region. In this direction of the Galactic plane the line of sight crosses the Perseus arm at a distance of $\\sim$3 kpc. One may expect that the majority of the sources are likely to be associated with this spiral arm. The detected sources are: (Anomalous X-ray Pulsar; AXP) 4U~0142+61, (high-mass X-ray binaries) 2S~0114+650, $\\gamma$~Cas, RX~J0146.9+6121, 4U~2206+54, the newly discovered IGR~J00370+6122, (intermediate polar) V709~Cas, (BL-Lac) 1ES~0033+595, and the currently unidentified newly discovered INTEGRAL source IGR~J00234+6141. INTEGRAL-flux upper limits are provided for some transient sources and a second AXP (1E~2259+586) in this field.  For the two AXP sources, 0.75--30 MeV archival data from COMPTEL \\citep{Schoenfelder93_comptel} are reanalyzed, delivering constraining upper limits for 4U~0142+61. In an attempt to reveal the nature of IGR~J00370+6122 archival data from the RXTE All-Sky Monitor and the BeppoSAX Wide-Field Cameras have been studied, as well as a Target of Opportunity (ToO) observation with RXTE PCA. ",
        "conclusions": "\\label{sec:disc}  We report results from a first deep observation of the Cassiopeia region in the hard X-ray energy window of 20--200 keV, exploiting the good imaging capabilities and the wide field of view of the IBIS-ISGRI telescope aboard INTEGRAL.  We report results on nine detected sources of various kinds: Anomalous X-ray Pulsar 4U~0142+61, several high-mass X-ray binaries, intermediate polar V709~Cas, \\mbox{BL-Lac} 1E~0033+595, and a still unidentified new INTEGRAL source IGR J00234+6141. In addition the upper limits are presented of the non-detected AXP 1E~2259+586 and three more high-mass X-ray binaries.  The discovery in this work of the very hard spectral tail between 20 and 150 keV of AXP 4U~0142+61, increases the number of IBIS-ISGRI AXP detections to three, the others being 1E~1841-045 \\citep{Molkov04_sagarm, Kuiper04_1841} and 1RXS~J170849-400910 \\citep{Revnivtsev04_gc}. This suggests that such a hard, highly luminous spectral component is a common property of AXPs. Of these three AXPs, 4U~0142+61 appears to exhibit the most extreme spectral shape (see Fig.~\\ref{fig:0142sp}): below 10 keV it has the softest spectrum with a black-body temperature of $kT = 0.470 \\pm 0.008$ keV and a very soft power-law component with photon index $3.40 \\pm 0.06$. Above 10 keV this drastically changes into a very hard power-law component with photon index 0.73$\\pm$0.17 (possibly even as hard as $-0.11 \\pm 0.43$ till $\\sim$70 keV if the indication for a break around that energy is real, see the fits in Fig. \\ref{fig:0142sp}). The X-ray luminosity between 20 and 100 keV for the single power-law fit is $5.9 \\times 10^{34}$ erg s$^{-1}$, a factor of 440 higher than the rotational energy loss, amounting $1.3 \\times 10^{32}$ erg s$^{-1}$ (assuming $M = 1.4\\, \\rm{M}_{\\odot}$, $R = 10$ km for a homogeneous sphere and a distance of \\mbox{3 kpc}). Such a high hard X-ray luminosity has to find its origin in the magnetosphere of the magnetar, fed from the magnetic energy reservoir in the strong magnetic field. The underlying scenario of production processes is not yet well understood. To constrain different model scenarios, more detailed observational parameters are needed. One of these is the detailed spectral shape of the hard X-ray component. The reported COMPTEL upper limits for energies between 0.75 MeV and 30 MeV already indicate that the power-law spectrum of 4U 0142+61 has to bend/break below this energy window. So far AXP 1E~2259+586 has not been detected with INTEGRAL, however, the flux upper limits extracted from INTEGRAL and COMPTEL data do not exclude the existence of a high-energy tail.  INTEGRAL has increased the number of known supergiant X-ray binaries significantly \\citep[see e.g.][]{Kuulkers05_HMXB, Negueruela06_hmxb, Lutovinov06_hmxb}.  We introduce a new SXB, i.e. IGR~J00370+6122. We detected this source in two INTEGRAL revolutions, at a position consistent with that of the ROSAT source 1RXS J003709.6+612131, which has a BN0.5II-III bright-giant counterpart. Analysis of archival RXTE-ASM data revealed the $15.665 \\pm 0.006$ day orbital period. Follow-up studies using a ToO RXTE observation and BeppoSAX-WFC archival data rendered a consistent picture of a new, by an order of magnitude variable, wind-accreting SXB. The source exhibits below 20 keV a power-law spectrum with index $\\sim$2.7, and a hard power-law spectrum with index $1.9 \\pm 0.3$ in the INTEGRAL window 20--80 keV.  We report the detection of a second SXB, 2S~0114+650. For the first time the spectrum has been measured up to $\\sim$150 keV. The good statistics in the time-averaged spectrum above 70 keV, showed that a single bremsstrahlung model cannot fit the data satisfactorily. In addition, a power-law component with index $2.0 \\pm 0.7$ is required, or the spectrum can be explained with a single power law with index $3.01 \\pm 0.06$.  Our deep survey of the Cassiopeia region rendered also detections of three Be X-ray binaries: $\\gamma$~Cas, RX~J0146.9+6121 and 4U~2206+54. For the first two, the INTEGRAL spectra were measured to higher energies than published earlier, and are discussed in comparison to these earlier reports.  The intermediate polar V709~Cas was detected up to $\\sim$70 keV.  The best model fit was obtained for a power-law model yielding a photon index of $3.00 \\pm 0.12$, different from the bremsstrahlung-model fits reported in literature.  Finally, two more hard-X-ray sources are found: 1) IGR~J00234+6141, discovered as a weak ($\\sim$1 mCrab) source up to 100 keV. No counterpart could be found yet. Since it is located close to the Galactic plane and there is no radio counterpart, it is likely an X-ray binary, rather than an AGN. 2) 1ES~0033+595, an AGN of the BL~Lac type, also detected in this work at the 1-mCrab level."
    },
    "0601/hep-ph0601147_arXiv.txt": {
        "abstract": "We discuss the cosmological implications of an antisymmetric tensor field when it experiences a kind of ``Higgs mechanism,'' including nonlocal interaction terms, on D-branes. Even when a huge magnetic-condensation-inducing anisotropy is assumed in the early universe, the expansion of the universe leads to an isotropic $B$-matter-dominated universe. ",
        "introduction": "If we are interested in the very early universe near or slightly below the Planck scale, the cosmological model driven by string theory should be taken into account. The first signal of string cosmology is given by dilaton gravity instead of Einstein gravity~\\cite{Pol}. In string theory, another indispensable bosonic degree is the antisymmetric tensor field of rank two~\\cite{Kalb:1974yc,Cremmer:1973mg}. When the ten-dimensional bulk is compactified to a (1+3)-dimensional spacetime, the cosmological effect of the massless antisymmetric tensor field has been taken into account. The homogeneous, but anisotropic, Bianchi-type universe is supported due to the antisymmetric tensor field~\\cite{Goldwirth:1993ha}. Once the anisotropy is generated, its effect can usually survive in the present universe, despite the sufficient expansion of the universe. The observed data of cosmic microwave background radiation (CMBR) threatens the viability of this string-driven cosmological model with the antisymmetric tensor field.  The development of D-branes and related topics for the last ten years has opened another possibility in cosmology that our universe may be identified with a D3-brane, or parts of higher-dimensional D-branes in a nine-dimensional spatial bulk. A clear distinction between string cosmology~\\cite{copeland:1995} in the bulk without D-brane and that on the D-brane appears through the mass generation of the antisymmetric tensor field on the D-brane~\\cite{Witten:1995im}. In the context of Einstein gravity, the cosmological effect of this massive antisymmetric field was studied in detail~\\cite{Chun:2005ee}. Specifically, the anisotropy induced by condensation of the antisymmetric tensor field is diluted during the expansion of the universe~\\cite{Kim:2004zq}. If the expansion is sufficient, the resultant present universe becomes isotropic and is free from the constraint of the CMBR.  In this paper, we will investigate the effect of the massive interacting antisymmetric tensor field. In Section II, we introduce the model of our interest, including graviton, dilaton, and antisymmetric tensor field in both the string frame and the Einstein frame. In Section III, we consider only a single magnetic component of the homogeneous antisymmetric tensor field and study its oscillation without damping in a flat spacetime. In Section IV, we consider another flat system of the antisymmetric tensor field with a dilaton. For the case of negative cosmological constant, the stabilization of the dilaton is achieved for some initial values. In Section V, the cosmological evolution of the antisymmetric tensor field is studied without and with a D3-brane. In the case of a D-brane universe, solutions of the isotropic $B$-matter-dominated universe are obtained. We conclude in Section VI with a summary and discussion. ",
        "conclusions": "We considered the antisymmetric tensor field on a D3-brane. In a flat spacetime with a stabilized dilaton, the condensed homogeneous magnetic field oscillated without damping. When the linear dilaton was turned on, we analyzed the difference between the string frame and the Einstein frame. Particularly, the effect of the cosmological constant from (1+9)-dimensions was drastic; i.e., the dilaton was stabilized for some initial values and a negative cosmological constant~\\cite{Kim:2005kr}.  In the early universe with gravitation, the existence of the D3-brane is significant. In the bulk without the D-brane, the anisotropy induced by the massless antisymmetric tensor field cannot be washed out through the cosmological evolution, but it remains as an observable fossil, which goes against the observed fact in the present universe. Once the mass and the self-interaction terms are included due to the presence of D3-brane in our universe, even the huge condensation of the homogeneous magnetic component is diluted as the universe expands. When the cosmological expansion is sufficiently large, the universe becomes isotropic and $B$-matter dominated.  It is intriguing to study cosmological evolution when all the closed string degrees in the bulk are included. Such a string cosmology of the graviton, the dilaton, and the antisymmetric tensor field or that of unstable D-brane~\\cite{Chung:2004ep} need further study."
    },
    "0601/astro-ph0601650_arXiv.txt": {
        "abstract": "{Characterizing the spatial and velocity structure of molecular clouds is a first step towards a better understanding of interstellar turbulence and its link to star formation. } {We present observations and structure analysis results for a large-scale ($\\sim$ 7.10 deg$^{2}$) $^{13}$CO J = 2--1 and $^{12}$CO J = 3--2 survey towards the nearby Perseus molecular cloud observed with the KOSMA 3m telescope.} {We study the spatial structure of line-integrated and velocity channel maps, measuring the $\\Delta$-variance as a function of size scale. We determine the spectral index $\\beta$ of the corresponding power spectrum and study its variation across the cloud and across the lines.} {We find that the spectra of all CO line-integrated maps of the whole complex show the same index, $\\beta \\approx 3.1$, for scales between about 0.2 and 3\\,pc, independent of isotopomer and rotational transition. A complementary 2MASS map of optical extinction shows a noticeably smaller index of $2.6$. In contrast to the overall region, the CO maps of individual subregions show a significant variation of $\\beta$. The $^{12}$CO 3--2 data provide e.g. a spread of indices between 2.9 in L\\,1455 and $3.5$ in NGC\\,1333. In general, active star forming regions show a larger power-law exponent. We find that the $\\Delta$-variance spectra of individual velocity channel maps are very sensitive to optical depth effects clearly indicating self-absorption in the densest regions. When studying the dependence of the channel-map spectra as a function of the velocity channel width, the expected systematic increase of the spectral index with channel width is only detected in the blue line wings. This could be explained by a filamentary, pillar-like structure which is left at low velocities while the overall molecular gas is swept up by a supernova shock wave.} {} ",
        "introduction": "The structure of the interstellar medium (ISM) is random to a large degree, with a complex spatial density distribution and velocity field. This is evident in large-scale surveys of spectral line tracers, such as CO, made for parts of the Galactic plane \\citep[e.g.][]{heyer1998,simon2001} and individual molecular clouds \\citep[e.g.][]{bally1987,ut87}. The spatial structure of the emission has been quantified in terms of its power spectrum \\citep{s87,sbhoz98,e99}. When fitting the azimuthally averaged power spectrum with a power law, the slope of the power law $\\beta$ provides information on the relative amount of structure at the linear scales resolved in the image. A pure power law is expected for structure on the linear scales of a self-similar turbulent energy cascade. Deviations from a power law are expected at scales where physical processes insert energy in the turbulence cascade (outflows, supernovae, super-bubbles, galactic rotation) and at scales of turbulence dissipation \\citep{es04}. Thus, a study of the scaling behavior of the cloud structure and the velocity field as traced by the power spectrum of observed spectral line maps can help to constrain the turbulent energy cascade in the ISM. A number of power-spectrum studies have been carried out for the atomic medium using \\ion{H}{i} observations \\citep{cd83,g93,ssdss99,dmsgg01,eks01}. They find spectral indices between about 2.7 and 3.7 for the line-integrated maps \\citep{fhl2004}.   The $\\Delta$-variance is an alternative method to determine the index of the power spectrum of isotropic images \\citep{sbhoz98}. In contrast to the power spectrum, the $\\Delta$-variance can be computed {\\it in the spatial domain}. It allows for a better separation of the intrinsic cloud structure from contributions resulting from the finite signal-to-noise in the data and the telescope beam. In addition, problems related to the discrete sampling of the data can be avoided. \\citet{bso01} presented a $\\Delta$-variance analysis for CO maps of the Polaris Flare and six other molecular clouds. The index $\\beta$ is close to 3 for most clouds, but it steepens in the Polaris Flare from 2.5 to 3.3 for maps with a linear resolution increasing from $\\ga 1$\\,pc to $\\la 0.1$\\,pc. \\citet{okh01} used the $\\Delta$-variance to study numerical models of self-gravitating, supersonically turbulent clouds and compared the results to observations. They showed that the $\\Delta$-variance traces deviations from an inertial scaling behavior at the scales of driving and dissipation \\citep[see also][]{om02}.   Constraints on the velocity structure of the ISM can be obtained from an analysis of individual channel maps of a spectral line cube, and by comparing the results with those obtained for the line-integrated maps.  \\citet{sl01} have analyzed \\ion{H}{i} observations of the Small Magellanic Cloud in this way, noting a systematic variation of the measured index $\\beta$ with the width of the velocity-channels, thus confirming theoretical predictions by \\citet{lp00}. Using the Boston University/Five College Radio Astronomy Observatory (BU/FCRAO) Galactic ring survey, \\citet{bs01} found that the index of the channel maps is smaller than that of the line-integrated maps. It showed a significant variation as a function of the channel velocity. But $\\beta$ did not increase with increasing velocity channel width, in contrast to the predictions by \\citet{lp00}. Apart from this first attempt, no systematic structure analysis of velocity-width dependent channel maps of molecular lines has been performed yet.   We present a new CO survey of the Perseus molecular cloud complex and analyze the cloud structure in the line integrated and the channel maps. The observations of the $^{12}$CO 3--2 and $^{13}$CO 2--1 transitions were made with the 3m telescope of the K\\\"olner Observatorium f\\\"ur Sub-Millimeter Astronomie (KOSMA) and cover the entire Perseus molecular cloud complex (7.1 square degrees). The Perseus cloud is one of the best examples of a nearby star-forming region. The cloud is related to the Perseus OB\\,2 association \\citep{bc86,ut87}, which is located in the area mostly free of CO emissions northeast of the cloud \\citep{ut87}. It includes a region where intermediate-mass stars form (NGC\\,1333), a young open cluster (IC\\,348), and a dozen dense cloud cores with low levels of star-formation activity (L\\,1448, L\\,1455, B\\,1, B\\,1 EAST, B\\,3 and B\\,5). 91 protostars and pre-stellar cores have been identified in a 3 square degree survey of the dust continuum at 850 and 450 $\\mu$m made with the James Clerk Maxwell Telescope, JCMT \\citep{hrfqlc05}.  Imaging observations of the Perseus complex in molecular cloud tracers exhibit a wealth of substructure, such as cores, shells, filaments, outflows, jets, and a large-scale velocity gradient \\citep{pbbjn99}. \\citet{pbbjn99} compared the structure traced by $^{13}$CO 1--0 observations to synthetic spectra and find that the motions in the cloud  must be super-Alfv\\'enic, with the exception of the B1 core, where \\cite{gchmt89} and \\cite{ctghkm93} detected a strong magnetic field. \\citet{pad03a} find that the structure function of the line-integrated $^{13}$CO 1--0 map follows a power law for linear scales between 0.3--3\\,pc, and \\citet{pad03b} compared the velocity structure of Perseus to MHD simulations.  A complete census of the stellar content of nearby ($\\la 350$\\,pc) molecular clouds (Perseus, Serpens, Ophiuchus, Camaeleon, and Lupus) is currently obtained by the Spitzer legacy project ``Cores to Disks'' \\citep[c2d,][]{eab03}. Large-scale $^{12}$CO, $^{13}$CO 1--0 and A$_{v}$ maps of the northern clouds were recently obtained by the COMPLETE team \\citep{goodman04}. The KOSMA survey of Perseus in higher CO transitions traces the warmer and denser gas due to the elevated critical densities and excitation energies ($\\sim$ 10$^{5}$ cm$^{-3}$ and 33.2\\,K for CO 3--2) relative to the J = 1--0 transition. Moreover, $^{12}$CO is largely optically thick, while $^{13}$CO, being a factor $\\sim$ 65 less abundant \\citep{lp90}, is often optically thin, thus tracing column densities. We plan to extend this work to the other nearby clouds using the KOSMA and NANTEN\\,2\\footnote{http://www.ph1.uni-koeln.de/nanten2} observatories.   The KOSMA observations are described in Sect.\\ 2. The general properties of the CO data sets are discussed in Sect.\\ 3. Sect.\\ 4 presents the results of the $\\Delta$-variance analysis. The discussion of the results and a summary are given in Section 5 and 6, respectively. ",
        "conclusions": "\\subsection{Integrated intensity maps}  Besides the $\\Delta$-variance, other tools have been used to characterize interstellar cloud structure. The second-order structure function for an observable $s(\\vec{r})$ is $S_\\mathrm{2} = \\langle|s(\\vec{r}) - s(\\vec{r}+\\delta \\vec{r})|^{2}\\rangle_{\\vec{r}}$ which is treated as a function of the absolute value of the increment $|\\delta \\vec{r}|$ \\citep{es04}. \\citet{pad03a} computed the structure function of the integrated intensity map of $^{13}$CO 1--0 in Perseus.  A power-law fit to S$_{2}$ $\\propto$ $\\delta$r$^\\mathrm{\\zeta}$ over a range of 0.3 to 3\\,pc provided an index $\\zeta$ of 0.83.  The index of the structure function is related to the power spectral index by $\\zeta=\\beta-2$ for $2<\\beta<4$ \\citep{sbhoz98} resulting in $\\beta=2.83$. This result of \\citet{pad03a} agrees within the error bars with the indices found by the $\\Delta$-variance analysis of the Perseus maps of integrated CO intensities over almost the same linear range (see Table~\\ref{deltatable}).   However, the $\\Delta$-variance spectra of individual regions show significant variations of the spectral index as discussed above. For $^{13}$CO 2--1, these span the range between $\\beta=2.86$ in L1455 and 3.76 in NGC\\,1333 (Table\\,\\ref{deltatable_indi}). The analysis of different sub-sets in molecular cloud complexes thus provides additional and complementary information on the structure of the cloud complex.  \\subsection{Velocity channel maps}   The velocity channel analysis (VCA), was used previously by \\citet{dmsgg01} to study \\ion{H}{i} maps of two regions in the 4th Galactic quadrant. One of the regions is rich in warm \\ion{H}{i} gas, the other is rich in cool \\ion{H}{i} gas. For the warm gas, \\citet{dmsgg01} find a systematic increase of the mean index with velocity channel width. The cold gas at lower latitudes behaves differently and shows rather constant indices of 2.7--3.1.   The latter results resemble the outcome of the VCA of the $^{12}$CO and $^{13}$CO channel maps in Perseus presented above. We find similar indices for the velocity integrated maps of the full region (cf. Table\\,\\ref{deltatable}) and the CO data show no significant variation of the index with velocity channel width when averaged over all velocity bins (Figure~\\ref{slopevel}a). The indices stay relatively constant. Since the bulk of the molecular gas traced by CO is even colder than cold \\ion{H}{i} gas, these results suggest a sequence of a reduced dependence of the spectral index of the channel maps on the bin widths from the warm to the cold ISM. The constancy of the average spectral index could also be explained by optical depth effects. \\citet{lp04} have shown that absorption can lead to an effective slice broadening, which leads in extreme cases to slice indices that become independent from the actual channel width.   Studying the spectral index of individual velocity bins of the CO data across the line profile (Fig.\\,\\ref{deltachan}), we find that the power-law indices increase with the velocity channel width in the blue wing (Fig.\\,\\ref{slopevel}b) while staying rather constant in the red wing.  No corresponding analysis was conducted for the \\ion{H}{i} data by \\citet{dmsgg01}. One possibility to explain the asymmetry between the blue and red wings might be a shock expansion of the CO gas in Perseus. \\citet{san74} found an expanding shell of neutral hydrogen which was created by a supernova in the Per\\,OB2 association a few $10^6$ years ago. At the location of the molecular cloud complex, the expansion is directed away from the Sun. Most of the associated molecular gas has been swept up by the shock, but pillar-like filaments have been left at the backside of the shock. They are visible in the channel maps (Fig. \\ref{velchan}) and produce the velocity dependence seen in the VCA of the blue line wings, i.e. the increase of indices with size of the velocity bin because of a gradual increase of large scale contributions across the line wing. We illustrate this scenario in Fig.~\\ref{explode}.   \\begin{figure} \\centering \\includegraphics[width=\\linewidth]{figures/h2_explode.eps}  \\caption{Sketch adopted from Fig.~3 of \\citet{san74} illustrating the spatial arrangement and motion of the Perseus cloud complex. The gas is swept up by a shock expansion with 12~km\\,s$^{-1}$. Due to the overall curvature, the line-of-sight velocity is 9~km\\,s$^{-1}$ for IC\\,348, but only 7~km\\,s$^{-1}$ for NGC1333. The diameter of the cloud is $\\sim$ 20\\,pc. Pillar-like structures are left at lower velocities as remainders of high-density regions which were not accelerated to the same velocity.} \\label{explode} \\end{figure}  The quantitative results from the velocity channel analysis can be interpreted in terms of the power spectrum of the velocity structure. \\citet{lp00} showed that the spectrum of velocity slices as a function of the velocity channel width is determined by the power spectral indices of the density structure $\\beta$ and the velocity structure $m$. They obtain different regimes for shallow ($\\beta < 3$) and steep ($\\beta > 3$) density power spectra: \\begin{center} $ P(k) \\propto \\left\\{% \\begin{array}{ll} k^\\mathrm{-\\beta+\\frac{m-3}{2}}, & \\hbox{thin slices,} \\\\ k^\\mathrm{-\\beta}, & \\hbox{thick slices,} \\\\ k^\\mathrm{-\\beta}, & \\hbox{very thick slices;} \\\\ \\end{array}% \\right. (\\beta < 3)$ \\end{center} \\begin{center} $ P(k) \\propto \\left\\{% \\begin{array}{ll} k^\\mathrm{-\\frac{9-m}{2}}, & \\hbox{thin slices,} \\\\ k^\\mathrm{-\\frac{3+m}{2}}, & \\hbox{thick slices,} \\\\ k^\\mathrm{-\\beta}, & \\hbox{very thick slices.} \\\\ \\end{array}% \\right. (\\beta > 3)$ \\end{center} Thin slices have a velocity width less than the local velocity dispersion at the studied scale; thick slices have a width larger than the velocity dispersion and very thick slices essentially correspond to the integrated maps \\citep{lp00}. We can assume that a single channel of our data corresponds to thin slices as they are much narrower than any observed line width. We use $\\beta$ from the integrated maps and the index measured in the single channels to derive $m$. From the indices obtained by averaging over the full line profile (Fig.~\\ref{slopevel}a), we obtain $m$ values of 3.9$\\pm$1.6, 3.9$\\pm$1.9 and 3.8$\\pm$1.7 for $^{12}$CO 3--2 in IC\\,348, NGC\\,1333 and L\\,1455, respectively; while $m$ is 3.9$\\pm$2.0, 3.5$\\pm$2.5 and 4.1$\\pm$1.8 for $^{13}$CO 2--1 in the same three regions. In the blue wings (Fig.~\\ref{slopevel}b), we obtain $m$ values of 3.2$\\pm$1.6, 3.4$\\pm$1.4 and 3.6$\\pm$1.5 for $^{12}$CO 3--2; while $m$ is 4.3$\\pm$1.9, 4.3$\\pm$1.7 and 4.1$\\pm$1.5 for $^{13}$CO 2--1 in the three regions. These values have large error bars, so that they are not directly suited to discriminate between different turbulence models. At least, we find that all values are consistent with Kolmogorov turbulence that gives $m \\sim 3.7$ \\citep{kol41}."
    },
    "0601/astro-ph0601466_arXiv.txt": {
        "abstract": "{We present evidence that the mid infrared ( MIR , rest frame 5-30 $\\mu$m ) is a good tracer of the total infrared luminosity, L(IR)($=~L[8-1000\\,\\mu{\\rm m}]$), and star formation rate (SFR), of galaxies up to $z\\sim$ 1.3. We use deep MIR images from the Infrared Space Observatory (ISO) and the Spitzer Space Telescope in the Northern field of the Great Observatories Origins Deep Survey (GOODS-N) together with VLA radio data to compute three independant estimates of L(IR). The L(IR,MIR) derived from the observed 15 and/or 24\\,$\\mu$m flux densities using a library of template SEDs, and L(IR,radio), derived from the radio (1.4 and/or 8.5 GHz) using the radio-far infrared correlation, agree with a 1-$\\sigma$ dispersion of 40\\,\\%. We use the k-correction as a tool to probe different parts of the MIR spectral energy distribution (SED) of galaxies as a function of their redshift and find that on average distant galaxies present MIR SEDs very similar to local ones. However, in the redshift range $z=$ 0.4-1.2, L(IR,24\\,$\\mu$m) is in better agreement with L(IR,radio) than L(IR,15\\,$\\mu$m) by 20\\,\\%, suggesting that the warm dust continuum is a better tracer of the SFR than the broad emission features due to polycyclic aromatic hydrocarbons (PAHs). We find marginal evidence for an evolution with redshift of the MIR SEDs: two thirds of the distant galaxies exhibit rest-frame MIR colors (L(12\\,$\\mu$m)/L(7\\,$\\mu$m) and L(10\\,$\\mu$m)/L(15\\,$\\mu$m) luminosity ratios) below the median value measured for local galaxies. Possible explanations are examined but these results are not sufficient to constrain the physics of the emitting regions. If confirmed through direct spectroscopy and if it gets amplified at higher redshifts, such an effect should be considered when deriving cosmic star formation histories of dust-obscured galaxies. We compare three commonly used SED libraries which reproduce the color-luminosity correlations of local galaxies with our data and discuss possible refinements to the relative intensities of PAHs, warm dust continuum and silicate absorption. ",
        "introduction": "Luminous infrared phases, during which galaxies radiate the bulk of their light in the infrared (IR) and are producing stars at rates larger than 20 M$_{\\odot}$ yr$^{-1}$, play a dominant role in the cosmic star formation history (CSFH) and in the production of the cosmic infrared background (CIRB) \\citep{CE,E02,Elbaz99,xu03,fran01,fran03,L04,chary04,papovich04}. Those studies suggest that at least half of present-day stars were formed in ``infrared luminous galaxies'', which we define here as the combination of luminous (LIRGs, $10^{11}$L$_{\\sun}\\leq L(IR)<10^{12}$L$_{\\sun}$) and ultra-luminous (ULIRGs, $L(IR)\\geq 10^{12}~$L$_{\\sun}$) infrared galaxies. Indeed, the comoving density of mid (15 \\& 24\\,$\\mu$m) and total IR luminosities ($L(IR)=~L[8-1000\\,\\mu{\\rm m}]$) due to infrared luminous galaxies was about 70 times larger at $z\\sim$ 1 than it is today \\citep{Lefloch,CE}. However, existing instruments do not allow direct measurement of $L(IR)$ and one must rely on correlations between other wavelength ranges and $L(IR)$ to derive its value as well as the obscured star formation rate (SFR) that it quantifies \\citep{Kennicutt98}.  $L(IR)$ can be derived from a single measurement in the mid IR (MIR, rest frame 5-30 $\\mu$m) for local galaxies with a 68\\,\\% uncertainty of 30-40\\,\\% depending on the MIR wavelength (see Figs.1\\&2 from \\cite{CE}). Similarly, the radio continuum exhibits an impressive correlation with L(IR) over four orders of magnitude with an uncertainty of only 5\\,\\% over nearly 2000 galaxies \\citep{Yun01}. However, it must be noted that when limited to the 162 infrared luminous galaxies of the sample of \\cite{Yun01}, this dispersion rises to 46\\,\\% which is comparable to the one observed between the mid-infrared and total IR luminosities. The origin of both correlations, linking the radio and MIR to L(IR), is not yet completely understood since several physical processes are generally advocated to explain the emission in both domains.  It is generally acknowledged that the radio emission is produced by synchrotron radiation from relativistic electrons and free-free emission from HII regions. Most of the energy required for the acceleration of electrons is produced by the supernova remnants of stars more massive than 8 M$_{\\odot}$, which also dominate the ionization in H II regions. However, whether this is enough to explain the extent of the correlation remains a matter of debate.  The correlation between mid and far IR luminosities of galaxies is more direct since both are produced by the re-emission of UV photons by dust, but the physics of the emitting sources is quite complex.  Above 60\\,$\\mu$m, the radiation of a galaxy is dominated by the grey body emission of ``big dust grains'' reaching an equilibrium temperature of 20-50 K. At shorter wavelengths, the spectrum of a galaxy results from the combination of :  - the continuum emission due to stochastically heated, warm dust grains, fluctuating around temperatures of a few hundred degrees. The carriers associated with this emission are commonly considered to be Very Small Grains \\citep[VSG]{DBP90}.  - the PAH (polycyclic aromatic hydrocarbon) bands due to stochastically heated molecules producing broad features in emission at 3.3, 6.2, 7.7, 8.6, 11.3 and 12.7\\,$\\mu$m.  - the absorption features due to silicates which are centered at 9.7 and 18\\,$\\mu$m.  Extragalactic surveys in the MIR, FIR, sub-millimeter and radio have all been used to trace star formation without being affected by dust extinction. However, with existing datasets and instrument sensitivities, the MIR is the most efficient wavelength to detect galaxies in the LIRG regime up to $z\\sim$ 2. Sub-millimeter surveys are limited to the mJy level where LIRGs are only detected up to $z\\sim$ 0.5 and with a bias towards objects with cooler dust temperatures for a fixed IR luminosity \\citep{chapman05}. Deep radio surveys provided a picture consistent with the one derived from the MIR but to a shallower depth \\citep{Haarsma00,E02,Elbaz03}. As a result, it is crucial to determine whether MIR properties of distant infrared luminous galaxies are similar to local ones and assess if the MIR can be used as a tracer of star formation in the distant universe. This is the goal of the present paper.  We will first compare the $L(IR)$ derived independantly from three tracers, namely the radio continuum at 1.4 and/or 8.5 GHz, the 15\\,$\\mu$m flux density from ISOCAM \\citep{Cesarsky96} onboard the Infrared Space Observatory \\citep[ ISO]{Kessler96} and the 24\\,$\\mu$m flux density from MIPS \\citep[Multiband Imaging Photometer for Spitzer][]{Rieke04} onboard the Spitzer Space Telescope \\citep{Werner04}. The flux density and variation of the k-correction at the two MIR wavelengths are used to constrain the relative contributions of the three components of MIR light : VSGs, PAHs and Silicate absorption. Using a similar strategy, \\cite{Elbaz05} and \\cite{Teplitz05} showed evidence that the MIR SED distant LIRGs did exhibit PAH features in emission. We will extend this work statistically by producing a direct comparison of the positions of local and distant galaxies in luminosity-luminosity diagrams at the same rest-frame wavelengths. This will allow us to test a possible variation of the relative roles played by the previous three ingredients of MIR SEDs which would have strong implications on the derivation of the CSFH.  While the derivation of L(IR) from a radio measurement depends only on the radio--far infrared (FIR) correlation and radio spectral index, which is well constrained for star forming galaxies, it is strongly dependant on the library of template SEDs that is chosen when using the MIR. Hence we will compare three commonly used SED libraries -- CE \\citep{CE}, DH \\citep{DH02} and LDP \\citep{L04} -- which present some noticeable differences in terms of the relative contributions of the three major components of MIR light.  We adopted a $H_{0}$=75~km~s$^{-1}$~Mpc$^{-1}$, $\\Omega_{M}=0.3$, $\\Omega_{\\Lambda}=0.7$ cosmology throughout this paper. ",
        "conclusions": "\\label{SEC:discussion}   We have presented a comparison of the mid infrared (MIR) and radio properties of local and distant galaxies (up to $z\\sim$ 1.3) with particular emphasis on IR luminous galaxies, i.e. LIRGs and ULIRGs. Our goal was to assess the possibility of deriving a total IR luminosity, L(IR), corresponding to a star formation rate (SFR), on the sole basis of a measured flux density in the MIR for those galaxies. We derived three L(IR) estimates using independantly the measured ISOCAM-15\\,$\\mu$m, MIPS-24\\,$\\mu$m and VLA-1.4 GHz flux densities and found that all three luminosities were consistent at the 40\\,\\% confidence level, i.e. an uncertainty similar to the one measured for the local MIR-FIR and radio-FIR correlations for IR luminous galaxies. This first result, obtained for a sample of 49 galaxies with redshifts up to $z\\sim$ 1.3, and a sub-sample of 18 galaxies detected in the radio, suggests that the MIR can be used as a tracer of the SFR of distant actively star forming galaxies. However, the L(IR) computed from 15\\,$\\mu$m shows a 10-20\\,\\% larger scatter than the one computed from 24\\,$\\mu$m in comparison to the radio derived L(IR). This result suggests that the warm dust continuum and neutral PAH is a better tracer of the SFR than is the ionized PAH emission at rest-frame 7.7\\,$\\mu$m.  The redshift evolution of the 24 over 15\\,$\\mu$m colors show a bump around $z\\sim$ 0.7 as one would expect from the entrance of PAH emission centered around 7.7\\,$\\mu$m into the 15\\,$\\mu$m passband. This behavior reinforces previous studies by \\cite{Elbaz05} and \\cite{Teplitz05} which were also interpreted as evidence for the presence of PAHs in distant IR luminous galaxies.  We attempted to probe the same rest-frame wavelength range in local galaxies and high redshift galaxies. We selected galaxies within two redshift bins centered at $z\\sim$ 0.6 and 1, where the observed 15\\,$\\mu$m (24\\,$\\mu$m) passband corresponds to the rest-frame 10 and 7\\,$\\mu$m (15 and 12\\,$\\mu$m) respectively.  In both redshift bins, the distant galaxies fall in the same region as local galaxies of equivalent luminosities and provides an extension of local correlations to higher luminosities. We find marginal evidence of evolution with redshift of the MIR SEDs: two thirds of the distant galaxies exhibit rest-frame MIR colors (L(12\\,$\\mu$m)/L(7\\,$\\mu$m) and L(10\\,$\\mu$m)/L(15\\,$\\mu$m) luminosity ratios) below the median value measured for local galaxies. Possible explanations are examined below. It is indeed to be expected that an evolution should take place if not at the $z\\sim$ 1 level, then at higher redshifts. The first cause that one may consider is the role played by metallicity. It was recently showed by \\cite{engelbracht} and \\cite{madden05} that the PAH over VSG continuum ratio was increasing as a function of metallicity, with a trend for galaxies with metallicities below $log(Z)$= 12$+log$(O/H)=8.2 to exhibit nearly no PAH contribution. \\cite{Liang04} found the gas metallicity of LIRGs ranging between half-solar to solar ($log(Z)$=8.63--8.93) at similar redshifts. Hence, the influence of metallicity is not expected to be strong in our sample. Furthermore, it would be responsible for a stronger $L_{12}/L_{7}$ for the distant galaxies instead of the lower values that we tend to find. We measure a $L_{12}/L_{7}$ ratio for the distant LIRGs which is similar to the one found for \"normal\" nearby galaxies. This may suggest that distant LIRGs, which are probably more gas rich than local ones, could also be triggered by less violent causes, i.e. minor mergers, gas infall or even spiral waves instead of the major mergers found in local galaxies, as suggested by their morphologies \\cite[see][]{bell05}. We note that the dust temperature measured in the far infrared also depends on the geometry of the regions of star formation \\cite[][]{chanial05,Wang92}. The lower $L_{10}/L_{15}$ ratio may result from a deeper 9.7\\,$\\mu$m silicate absorption. Another possibility would be that the level of ionization of PAHs in distant galaxies is higher:  more ionised PAHs tend to present a lower $L_{12}/L_{7}$ \\citep[e.g.][]{Dartois97}. Other causes might also play a role such as a different grain size distribution but it is not possible to disentangle between all those possibilities with existing data. Extended samples of objects at higher redshifts with Spitzer 16\\,$\\mu$m and 24\\,$\\mu$m imaging as well as mid-infrared spectroscopy will soon be available which might help discriminate between the various effects.  We tested three commonly used libraries of template SEDs-- CE \\citep{CE}, DH \\citep{DH02} and LDP \\citep{L04} -- in which the relative strengths of broad emission lines from PAHs with respect to warm dust continuum due to VSGs increases from CE to DH and to LDP.  The LDP library predicts values for L(IR) which are 40\\,\\%  too low from the measured 15\\,$\\mu$m flux density and 30\\,\\% too large from the 24\\,$\\mu$m one, with respect to the radio ones. They also exhibit the largest dispersion and inconsistency between L(IR,15) and L(IR,24). This is mostly due to enhanced 8\\,$\\mu$m and decreased the continuum between 11-32 \\,$\\mu$m and  and to a lesser extent, due to a lack of evolution of MIR SEDs as a function of L(IR).  The CE and DH libraries both provide robust fits to both local and distant galaxies but we find evidence for more PAH emission in IR luminous galaxies than predicted in the CE SEDs. As a result, the CE library would overestimate L(IR) for distant galaxies by 20-30\\,\\%. If, as suggested above, distant galaxies do exhibit stronger PAH emission than local ones, then both libraries need to be revised and the CE library would overestimate L(IR) by up to 50\\,\\% for ULIRGs at $z\\sim$ 1. Since LIRGs dominate the star formation rate density per unit comoving volume at least up to $z\\sim$ 1.5, this evolution of MIR SEDs does not significantly affect the derivation of the cosmic star formation history (CSFH). We have checked this by replacing the CE SEDs by the DH ones in the \\cite{CE} model and found two CSFH which are consistent within 30\\,\\%.  The true accuracy with which FIR luminosities can be determined at high redshift are being tested for the most luminous objects with Spitzer, and will be pushed to fainter luminosities with the Herschel satellite scheduled for launch 2007, and the Atacama Large Millimeter Array, ALMA. Direct MIR spectroscopy with the InfraRed Spectrograph onboard Spitzer will also provide very useful constraints if pushed to detect distant LIRGs."
    },
    "0601/astro-ph0601716_arXiv.txt": {
        "abstract": "Radial velocities from the 2.1 m telescope at McDonald Observatory supplemented with radial velocities from the coud\\'e feed telescope at KPNO provide new precise orbits for the double-lined spectroscopic binaries RR~Lyn (A3/A8/A6), 12~Boo (F8IV), and HR~6169 (A2V).  We derive orbital dimensions ($a_1\\sin i$ and $a_2\\sin i$) and minimum masses ($m_1\\sin ^3i$ and $m_2\\sin ^3i$) with accuracies of 0.06 to 0.9\\,\\%. The three systems, which have $V$ magnitudes of 5.54, 4.83, and 6.42, respectively, are all sufficiently bright that they are easily within the grasp of modern optical interferometers and so afford the prospect, when our spectroscopic observations are complemented by interferometric observations, of fully-determined orbits, precise masses, and distances.  In the case of RR~Lyn, which is also a detached eclipsing binary with a well-determined orbital inclination ($i = 87\\fdg45 \\pm 0\\fdg11$; Khaliullin et al. 2001), we are able to determine the semimajor axis of the relative orbit, $a = 29.32 \\pm 0.04$\\,R$_\\odot$, primary and secondary radii of $2.57 \\pm 0.02$\\,R$_\\odot$ and $1.59 \\pm 0.03$\\,R$_\\odot$, respectively; and primary and secondary masses of $1.927 \\pm 0.008$\\,M$_\\odot$ and $1.507 \\pm 0.004$\\,M$_\\odot$, respectively. Comparison of our new systemic velocity determination, $\\gamma = -12.03 \\pm 0.04$ km~s$^{-1}$, with the earlier one of \\citet{k76}, $\\gamma = -11.61 \\pm 0.30$ km~s$^{-1}$, shows no evidence of any change in the systemic velocity in the 40 years separating the two measurements, a null result that neither confirms nor contradicts the presence of the low-mass third component proposed by \\citet{kk02}. Our spectroscopic orbit of 12 Boo is more precise that that of \\citet{betal05} but confirms their results about this system. Our analysis of HR~6169 has produced a major improvement in its orbital elements.  The minimum masses of the primary and secondary are 2.20 $\\pm$ 0.01 and 1.64 $\\pm$ 0.02 M$_{\\sun}$, respectively. Although all three systems have eccentric orbits, the six components of the systems are either pseudosynchronously rotating or very nearly so. ",
        "introduction": "The current generation of ground-based interferometers, such as the Palomar Testbed Interferometer (PTI) \\citep{cetal99}, the Naval Prototype Optical Interferometer (NPOI)\\citep{hetal03}, the Infrared Optical Telescope Array (IOTA3)\\citep{tetal03}, and the Center for High Angular Resolution in Astronomy (CHARA) array \\citep{tbetal03}, is advancing stellar astronomy in a number of ways.  \\citet{q01}, for example, reviewed the state of optical and infrared interferometry.  One direction of progress is the increasing number of spectroscopic binaries that are being resolved as visual binaries.  This allows the determination of their three-dimensional orbits and the derivation of accurate masses for the component stars and distances to the systems, distances that in many cases are more accurate than those from the {\\it Hipparcos} satellite.  In recognition of this development we have started a program to determine substantially improved spectroscopic orbits for bright, field spectroscopic binaries.  The program has two benefits: the provision of new radial velocities and spectroscopic orbits of a quality that matches or exceeds the prospective interferometric observations and, for some binaries, the detection of the secondary spectrum and measurement of secondary radial velocities for the first time.  We now briefly consider these two points in turn.  While some interferometric studies, such as that of 12~Boo \\citep{betal05}, include complementary new radial velocities, the usual practise is to take the radial velocities for the binary concerned from the literature.  The precision of such velocities often falls short of that needed to match the interferometric observations.  For example, in their recent determination of the three-dimensional orbit of the bright spectroscopic binary $\\sigma$~Psc, \\citet{kl04} had to complement their interferometric measurements with radial velocities observed in 1944 and 1945 \\citep{b47}. Their resulting best-fit solution for the three-dimensional orbit has rms velocity residuals of 4.8 and 3.6 km~s$^{-1}$ for the primary and secondary, respectively.  Orbits with large velocity residuals are not exceptional because of the generally lower resolution and low signal-to-noise ratio of spectra obtained in the first roughly three-quarters of the twentieth century,  For example, of the first 100 systems in the {\\em Eighth Catalogue of the Orbital Elements of Spectroscopic Binaries} \\citep{betal89}, 63 have orbits that were published in 1980 or earlier and 24 have orbits that were published in 1950 or earlier, long before the advent of radial velocity spectrometers and charge-coupled device detectors, which can produce spectra with very high signal-to-noise ratios.  Similar proportions must apply for all 1469 systems in the Catalogue\\footnote{ One would like to cite the {\\em Ninth Catalogue of Spectroscopic Binary Orbits} \\citep{petal04} here, but at present it is incomplete with respect to orbits published since 1989 and so is unsuitable for the point we are making.}.  While these proportions will have improved as a result of the substantial number of new spectroscopic binary orbits that have been published since 1989 \\citep{petal04}, most such orbits are for newly discovered binaries.  With respect to the detection of the secondary spectrum, we note that without secondary radial velocities and a determination of the secondary's spectroscopic orbit, the linear separation between the binary components is unknown and the determination of the three-dimensional orbit is incomplete.  Increasing the pool of double-lined spectroscopic binaries (SB2s) thus increases the number of spectroscopic binaries available for fruitful interferometric observation.  In addition, binary systems with components of significantly different masses provide the greatest constraints on evolutionary models.  Considering that the majority of spectroscopic binaries are single-lined spectroscopic binaries (SB1s), there is ample opportunity here.  \\citet{petal04}, for example, found that two-thirds of the spectroscopic binaries in their {\\em Ninth Catalogue} are SB1s.  (There is no reason to think the Catalogue's incompleteness affects this statistic much.)  Our program uses new, high-resolution, red-wavelength spectra obtained with the 2.1 m telescope at McDonald Observatory of the University of Texas and the coud\\'e feed telescope at Kitt Peak National Observatory (KPNO). \\citet{ft04} provided a preliminary description of our program and an initial list of observed stars, which has now been expanded to over 40 systems. These come primarily from a sample of 130 candidate systems obtained by searching the {\\em Eighth Catalogue} for SB2 systems that could profit from renewed spectroscopic observation and SB1 systems with large enough mass functions to suggest that high signal-to-noise ratio spectra might transform them into SB2 systems \\citep[e.g.,][]{sf92}.  The stars are north of $-$40\\arcdeg\\ in declination and generally brighter than $V = 7.5$ mag.  Others have also seen the need for improved radial velocities for spectroscopic binaries.  For example, \\citet{k05} has successfully applied the iodine absorption-cell method for determining very precise radial velocities to the measurement of radial velocities of {\\em both} components in SB2s.  Hitherto, this technique, which uses an iodine absorption cell to impose a reference spectrum on the stellar spectrum and is notable for its use in the discovery of extrasolar planets, has been restricted to the radial velocities of single stars or stars with companions of insignificant relative brightness.  His pioneering investigation, which was carried out on the Keck~I telescope with the HIRES spectrograph, was limited to five objects including a radial-velocity standard and two SB2s.  Among the latter was 64~Psc (HD~4676), a well-known, bright spectroscopic binary (F8V, $P = 13.8$ days) with a three-dimensional orbit determined by \\citet{betal99}, using their own interferometric observations made with PTI and radial velocities from \\citet{dm91}.  Konacki's combined fit of his new radial velocities and the \\citet{betal99} interferometric data leads to better-determined orbital and physical parameters for 64~Psc. In particular, the rms velocity residual of 24 m~s$^{-1}$, determined from his new fit, is a striking improvement compared to the rms residual of 810 m~s$^{-1}$ (average $|O - C| = 650$ m~s$^{-1}$) given by \\citet{dm91} for their radial-velocity solution.  Our new velocities are not as precise as the iodine-cell velocities from the Keck~I telescope --- we will see that for 12~Boo, which has a similar spectral type to 64~Psc, the rms residual for our velocities is 110 m~s$^{-1}$ --- but their quality is still a good match with that of current interferometric observations.  There is thus ample opportunity for radial velocities from small and medium-sized telescopes, measured by traditional methods, to make a worthwhile contribution.  Here we report new orbit determinations for three bright spectroscopic binaries, RR~Lyn, 12~Boo, and HR~6169.  Analysis of our new velocities provides significant improvements in the orbital elements of the systems.  Indeed, in the case of HR~6169 some of the revisions are substantial.  Table~1 gives basic data for the systems; all three are known SB2s, have orbital periods near 10 days, and eccentric orbits. We now briefly look at each system individually.  \\subsection{RR Lyncis} RR~Lyn has long been known as both an eclipsing and spectroscopic binary. Its orbit has a period of 9.95~days and a modest eccentricity, $e = 0.08$. The star's peculiar abundances are indicated in its metallic-lined spectral classification, which reflects primarily the spectral type of the primary. \\citet{am95} classified it as A3/A8/A6, based on its Ca~II K line, its hydrogen Balmer lines, and its metallic lines, respectively, while \\citet{r49} called the hydrogen lines A7 and the metallic lines F0.  \\citet{ketal01} compared the two components with evolutionary tracks and also used their photoelectric photometry on the $WBVR$ system \\citep{ketal85} to place the components of RR~Lyn in a $(B - V, V - R)$ diagram.  From those results they estimated individual spectral types of A6~IV for the primary and F0~V for the secondary.  Thus, the primary has already begun to leave the main sequence, while the secondary is still ensconced within it. The metallic-line status of the primary is not in doubt \\citep{p71,k76,lr95}, while that of the secondary remains uncertain, although the abundances of \\citet{lr95} indicate possible metallicism.  \\citet{h15} published the first spectroscopic orbit.  More recent, but now quite old orbits were published by \\citet{p71} and \\citet{k76}.  Kondo's orbit has the advantages that it sampled the primary and secondary radial velocity curves more fully than Popper's and that it has a single systemic velocity, while Popper determined different systemic velocities for the primary and secondary.  Photoelectric light curves of RR Lyn have been reported by \\citet{h31}, \\citet{mk59}, \\citet{b60}, \\citet{l66}, \\citet{lavetal88}, and \\citet{ketal01}.  The last of these studies used precise photometry in the $W, B, V,$ and $R$ bandpasses \\citep{ketal85}, which are based on but somewhat different from the Johnson $U, B, V,$ and $R$ photometric bandpasses.  \\citet{ketal01} showed that the primary and secondary eclipse depths in $V$ are $\\sim 0.37$ and $\\sim 0.25$, respectively, the fractional luminosities, also in $V$, are $L_1 = 0.7835 \\pm 0.0039$, $L_2 = 1 - L_1$, and the orbital inclination is $87\\fdg45 \\pm 0\\fdg11$.  Both \\citet{l66} and \\citet{b74} needed the addition of third light to solve the light curves, while \\citet{k76}, who also did a photometric analysis, found no requirement for third light.  This last result is supported by \\citet{ketal01}, who demonstrated that their observations in all filters are fitted by the same geometry and that their model matches the observations without any need to invoke third light.  In a later paper, however, \\citet{kk02} presented new evidence for a third body, based {\\it not} on an analysis of the light curve, but on the times of the eclipse minima.  The presence of the third star, which \\citet{kk02} suggested is a very low-mass, low-luminosity object that contributes an insignificant amount of light to the system, is inferred from quasi-periodic oscillations in the times of primary and secondary minima.  They proposed an extemely eccentric orbit ($e = 0.96 \\pm 0.02$) with a rather long orbital period of $39.7 \\pm 4.2$ years for it.  The 64-year timespan of the published photoelectric minima amounts to only one-and-a-half orbital periods, however, so the case for the third body cannot yet be regarded as conclusive.  \\subsection{12 Bootis} The SB2 system 12~Boo is the only one in our trio that already has a determination of its three-dimensional orbit \\citep{betal05}.  The two components have an orbital period of 9.60~days and are of very similar mass, having a secondary to primary mass ratio of 0.97.  The system's combined spectral type is F8~IV \\citep{r50} or F8~V \\citep{getal01}. \\citet{betal05} found that the primary is leaving the main sequence, and with $T_{eff} = 6130 \\pm 100$~K it is now marginally {\\em cooler} than the secondary ($T_{eff} = 6230 \\pm 150$~K).  \\citet{h14} was the first to determine its spectroscopic orbit.  Newer orbital elements have been published by \\citet{al76} and \\citet{du99}. \\citet{betal05} presented the newest set of radial velocities as well as interferometric observations.  Details for these more recent spectroscopic orbits, including the present investigation, are summarised in Table~2. Prior to its publication, we were unaware of the new investigation of 12~Boo by \\citet{betal05}. Thus, in spite of the element of (unintentional) redundancy here, our own study of 12~Boo, based on radial velocities acquired from 2002 to 2005, is still of interest because our velocities turn out to be significantly more precise, allowing us to determine improved spectroscopic orbital elements.  \\subsection{HR 6169} Compared to RR~Lyn and 12~Boo, the previous work on HR 6169 has been quite modest.   This binary has a combined spectral type of A2~V \\citep{cetal69} or A1~IV \\citep{am95}.  The first and, to date, only spectroscopic orbit is that of \\citet{y20}, which was based on radial velocities from blue photographic spectrograms obtained at the Cassegrain focus of the 1.8 m telescope at the Dominion Astrophysical Observatory (DAO).  This SB2 system has an orbit with a period of 10.56 days and a moderate eccentricity, $e = 0.41$. ",
        "conclusions": "We have determined new spectroscopic orbits for the bright SB2 systems RR~Lyn, 12~Boo, and HR~6169.  The accuracies of the minimum masses ($m_1\\sin ^3i$ and $m_2\\sin ^3i$) range from 0.13\\% for the secondary of 12~Boo to 0.9\\% for the secondary of HR~6169, while the accuracies of the linear sizes of the relative orbit ($a\\sin i$) range from 0.05\\% for 12~Boo to 0.25\\% for HR~6169.  These results show our radial velocities provide well-determined spectroscopic orbits. which do justice to interferometric measurements of 12~Boo \\citep{betal05} and, we expect, will also do justice to prospective interferometric measurements of RR~Lyn, HR~6169, and other systems in our program.  For RR~Lyn, an eclipsing binary with a well-determined orbital inclination, we have determined accurate new masses and radii for its components.  The semi-major axis of the relative orbit, $a = 29.32 \\pm 0.04$\\,R$_\\sun$, plus the distance to the system, $73.5 \\pm 2.8$~pc \\citep{ketal01}, lead to a predicted angular size $a = 1.86$~mas for the true orbit and, because $\\omega$ ($= 179\\fdg4 \\pm 0\\fdg6$) is $180^\\circ$ to within its measurement error so the semi-major axis lies across the line-of-sight, this will also be the $a$ of the apparent orbit.  We find that the expected effect of the unseen third star proposed by \\citet{kk02} on our and Kondo's (1976) radial velocities is so slight as to be well below the threshold of our measurement errors.  We have determined a new spectroscopic orbit for 12~Boo, which confirms, and is more precise than, that determined by \\citet{betal05} as part of their measurement of the three-dimensional orbit of this system.  Our investigation of HR~6169 corrects systematic errors in Young's (1920) orbit and leads to significant revisions of the minimum masses and $a_{1,2}\\sin i$.  At the distance of HR~6169, $162 \\pm 23$\\,pc (Perryman et al. 1997), the $a\\sin i$ of 22.05 $\\pm$ 0.05~Gm corresponds to an angular separation of 0.91\\,mas.  In conclusion we find that radial velocities observed with small and medium-sized telescopes and measured by traditional methods provide a suitable spectroscopic contribution to the combination of measurements required for determination of three-dimensional orbits."
    },
    "0601/astro-ph0601520_arXiv.txt": {
        "abstract": "Stars and compact objects that plunge toward a black hole are either 1) captured, emitting gravitational waves as the orbit decays, 2) tidally disrupted, leaving a disc of baryonic material, 3) scattered to a large radius, where they may thereafter avoid encounters with the black hole or 4) swallowed whole, contributing to black hole growth. These processes occur on a dynamical time, which implies that for a static spherically symmetric stellar system, the loss cone is quickly emptied. However, most elliptical galaxies and spiral bulges are thought to be triaxial in shape. The centrophilic orbits comprising the backbone of a triaxial galaxy have been suggested as one way to keep the loss cone around a supermassive black hole filled with stars, stellar remnants, and intermediate mass black holes. We investigate the evolution of the loss cone population in a triaxial galaxy model with high resolution N-body simulations. We find that enough regular orbits flow through angular momentum space to maintain a full loss cone for a Hubble time. This increases the astrophysical capture rate by several orders of magnitude; we derive new capture rates for the Milky Way bulge and a triaxial M32-analogue. In the Milky Way, for example, we find that the white dwarf capture rate can be as high as $10^{-5}$ per year, 100 times larger that previous estimates based on spherical models for the bulge. This is the first step in a multiphase project that will explore the dynamics of supermassive black holes in realistic, fully self-consistent systems. The phase space content and time evolution of the loss cone affects supermassive black hole merger rates, extreme mass ratio inspiral rates, as well as the growth rate of a single supermassive black hole due to stellar accretion, all key observables for the future Laser Interferometer Space Antenna (LISA). ",
        "introduction": "Observations suggest that supermassive black holes are a part of nearly every galaxy center (e.g. Richstone et al. 1998), with masses ranging from $O(10^6)$ - $O(10^{10}) M_\\odot$. Gas accretion is thought to fuel the early stages of black hole growth, which may explain the tight correlation between black hole mass and the host galaxy's global velocity dispersion (Gebhardt et al. 2000a; Ferrarese \\& Merritt 2000). Indeed, a supermassive black hole may be essential to regulate its host galaxy's global structure and kinematics.  However, a single supermassive black hole (SMBH, hereafter) exacts a toll on its host galaxy, steepening the central cusp (Young 1980, Quinlan et al. 1995), and preferentially depleting the system of its low angular momentum material. Stars on plunging orbits interact with the SMBH and are removed from the galaxy by four basic processes:  \\begin{enumerate} \\item Tidal Disruption  A star is disrupted when it passes so close to the SMBH that the tidal force exerted by the SMBH exceeds the star's own surface gravity. These stars are commonly thought to be the main source of the large amplitude X-ray outbursts seen in both active and inactive galactic nucleii (e.g. Donley et al. 2002) as the gaseous debris settles into an accretion disc and falls into the  SMBH (e.g. Rees 1988, Bogdanovic et al. 2004). The tidal disruption radius, $r_{\\rm tidal}$ depends on the stellar mass, radius, and internal structure:  \\begin{equation} {r_{\\rm tidal}} =  {{\\Bigg( {{\\eta^2 {M_{\\bullet} }} \\over {{m_\\star}}}\\Bigg)^{1/3} } r_{\\star}}, \\end{equation}  where $\\eta \\sim 2.21$ for an incompressible, homogeneous system, and $\\eta = 0.844$ for a main sequence star modeled with a $n=3$ polytrope (Magorrian \\& Tremaine 1999; Sridhar \\& Tremaine 1992, Diener et al. 1995). A G2V star would have to pass within $0.64$ AU of the Milky Way SMBH to be disrupted, assuming the current best Milky Way SMBH mass estimate of $3.7 \\times 10^6 M_\\odot $ (Ghez et al. 2005).  \\item Swallow Directly  For high mass SMBHs, such that $M_\\bullet >\\sim 10^8 M_\\odot$, the tidal disruption radius of a main sequence star can be smaller than the Schwarzschild radius, $r_\\bullet = 2 G M_\\bullet / c^2$. In this case, stars plunge directly into the SMBH, possibly without any accompanying electromagnetic event. Since the tidal disruption radii of compact objects are much smaller, they can be swallowed whole by lower mass SMBHs as well: 10-30 $\\%$ of the total flux of stellar objects into a $10^6 M_\\odot $ SMBH can be attributed to compact objects and dark matter particles that have been swallowed whole (Zhao, Haehnelt \\& Rees 2002; though see Hopman \\& Alexander for a larger estimate).  \\item Capture and Inspiral  Those objects that survive a close pericenter pass without disruption can be captured by the SMBH on a bound quasi-Keplarian orbit. Typically, this occurs when a white dwarf, neutron star or black hole (of any mass) passes within $few-100 M_\\bullet$ of the SMBH (in $c=G=1$ units, i.e. $1.5-7.6$ AU for the Milky Way SMBH)\\footnote{For this paper, we ignore captures from neutron stars}. This newly formed close binary has a large dynamic quadrupole potential and begins to emit significant gravitational radiation (Peters 1964) with a characteristic amplitude $h$:  \\begin{equation} {h} \\sim {{4 \\times 10^{-24}} {\\Bigg( {{m_\\star} \\over {d}}\\Bigg)^{1/3} } {\\Bigg( {{M_\\bullet} \\over {10^6}}\\Bigg)^{2/3}} {\\Bigg( {{10^4} \\over {P}}\\Bigg)^{2/3}    }}, \\end{equation}  where the orbital period P is measured in seconds, the stellar remnant mass $m_\\star$ is measured in solar masses, and the distance to the source d is measured in Gpc. As a consequence of losing energy via gravitational radiation, the compact object spirals in and eventually coalesces with the SMBH. These extreme mass ratio inspirals (EMRIs) are a prime observational candidate for LISA, a planned gravitational wave telescope set to launch in 2017. Current estimates suggest that LISA will observe about one stellar mass EMRI per year within 1 Gpc of Earth (Hils and Bender 1995, Sigurdsson \\& Rees 1997, Sigurdsson 1998, Freitag 2003, Ivanov 2002, though see Hopman and Alexander 2005 for a significantly  lower estimated rate of detection.)  \\item Ejection  Interactions between stars can induce changes in each stellar velocity of the order:   \\begin{equation} {\\delta v_1} \\sim {{4.4 \\times 10^{2} {\\rm {km}\\over {s}}} {\\Bigg( {{2 m_2} \\over {m_1 + m_2}}\\Bigg)^{{1}\\over{2}} } {\\Bigg( {{m_2} \\over {1 M_\\odot}}\\Bigg)^{{1}\\over{2}}} {\\Bigg( {{1 R_\\odot} \\over {b}}\\Bigg)^{{1}\\over{2}}    }}, \\end{equation}  where b is the impact parameter between the two stars, $m_1$ and $m_2$ are the masses of stars 1 and 2, respectively (Yu \\& Tremaine 2003). It is quite rare to eject main sequence stars with a large enough velocity to escape the Milky Way entirely by this process, occurring less than once per Hubble time. However, if the two stars were bound as a binary, the SMBH can break the binary apart and eject one of the stars with $O(10^3)$ km/sec velocities (Hills 1988, Gould \\& Quillen 2003, Yu \\& Tremaine 2003). This process occurs more commonly for the Milky Way, at a rate of $10^{-5} (\\eta/0.1)$ per year, where $\\eta$ is the binary fraction (Yu \\& Tremaine 2003).  Though we have described this 3-body removal process in terms of two stars and one SMBH, ejection is a much more efficient process when there are two black holes and one star. If there are two black holes of significant mass bound at the center of the Milky Way, either another lower mass SMBH or an intermediate mass black hole, then the number of stars that can be ejected is much larger, because stars need pass only within the separation between the two black holes to receive energy and angular momentum. Given a second SMBH with a mass much smaller the the first, Yu \\& Tremaine (2003) estimate the $O(10^{-4})$ stars are ejected per year, with as many as 1000 stars having ejection velocities $ > 10^3$ km/sec inside the solar radius (see also Mapelli et al. 2005). \\end{enumerate}   Since stars are lost from the system by these 4 processes, they comprise the 'loss cone' of the galaxy.\\footnote{ Traditionally, the term 'loss cone' has referred only to the tidal disruption process. We have settled on a broader definition. Tidal disruption is the dominant process for main sequence stars interacting with a Milky Way-scale SMBH, while captures are by far the dominant processes for compact objects such as white dwarfs, neutron stars and up to intermediate mass black holes.} The loss cone is rapidly depleted of its stellar reservoir in at most a few dynamical times. With a depletion timescale this rapid, it has become a puzzle to understand how the loss cone could be anything other than empty, overall. In a static, spherical galaxy model, for example, the loss cone can only be refilled by 2-body relaxation, as long as 3-body scattering ejects stars permanently (e.g. Milosavljevi\\'{c} \\& Merritt 2003). Because the timescale for refilling the loss cone via 2-body relaxation can be much longer than a Hubble time,    \\begin{equation}{ {T_{\\rm relax}} \\sim {2 \\times 10^9 {\\rm yr} \\Big({{\\sigma}\\over {200 / {\\rm km/sec}}}\\Big)^3 \\Big({{\\rho}\\over {10^6 / {\\rm M_\\odot/pc^3}}}\\Big)^{-1}},} \\end{equation}  the loss cone remains empty over a galaxy lifetime; this will render any type of stellar interaction with a SMBH very inefficient, at least in static, spherical models.  Filling a loss cone, relies on maintaining a reservoir of low angular momentum stars within the galaxy, so considering more realistic galaxy shapes may be the key. While most loss cone studies have concentrated on static, spherical galaxy models, theory and observations both indicate that dark matter halos and elliptical galaxies are at least mildly triaxial (Bak \\& Statler 2000, Franx, Illingworth, \\& de Zeeuw 1991). Triaxiality is present not only elliptical galaxies, but in disc galaxies as well: a barred galaxy is a prime example of a rotating triaxial ellipsoid, which has a more complex orbital structure and response to a SMBH (Hasan \\& Norman 1990, Sellwood \\& Shen 2004). Over 70 percent of the local disc galaxy population is barred, and early indications suggest that this fraction may stay constant out to redshift 1 (Jogee et al. 2004; Sheth et al. 2004). Our own galaxy hosts a bar with a semi-major axis length of about 2 kpc seen nearly end on from our perspective ($\\phi_{\\odot -{\\rm bar}} \\sim 20^\\circ$; see Gerhard 2001 for a review).   Non-axisymmetry can introduce more stars to the loss cone in several ways. First, stars in even a mildly triaxial potential move in entirely different orbit families than are present in a spheroid (see Figure 1). In particular, there are a rich variety of regular box and boxlet orbits that are centrophilic and comprise the backbone of the galaxy (Miralda-Escude \\& Schwarzschild 1989). These centrophilic orbits can pass formally through, or very near to the SMBH, which make them the primary regular orbital component of the loss cone in a non axisymmetric galaxy.  Besides the regular box and boxlet orbits, chaotic orbits often occur a triaxial galaxy. Though chaotic orbits are generically in a triaxial galaxy, they are particularly prevalent in one that is embedded with a SMBH. In fact, it is a commonly thought that SMBHs cannot exist in a stable triaxial galaxy, because black holes induce chaos in the centrophilic orbit families (Norman, May \\& van Albada 1985, Gerhard \\& Binney 1985, Miralda-Escude \\& Schwarzschild 1989, Merritt \\& Quinlan 1998, Valluri \\& Merritt 1998, Holley-Bockelmann et al. 2002, Poon \\& Merritt 2002). and drive the galaxy toward axisymmetry in a few crossing times. Fully self-consistent, high resolution n-body simulations show, however, that triaxiality is not entirely destroyed by a SMBH. Though SMBHs do indeed incite chaos in box and boxlet orbits at small radii and do rounden the region inside the radius of influence $r_{\\rm inf} \\equiv G M_\\bullet / \\sigma^2 $, most of the centrophilic orbits are simply scattered onto resonant (or sticky) orbits that allow triaxiality to persist over most of the galaxy for many Hubble times (Holley-Bockelmann et al. 2002; see also Poon \\& Merritt 2002).  Due to their stochastic motion through phase space, chaotic orbits can refill the loss cone as well. Stars on chaotic orbits encounter the SMBH once per crossing time with a pericenter separation that varies each time, eventually passing close enough to the SMBH to be captured or tidally disrupted. Merritt \\& Poon (2004) showed that, for triaxial models where the fraction of chaotic orbits is substantial ( $\\sim 50 \\%$), the capture rate deep inside the sphere of influence of the SMBH meets and even exceeds that of a full loss cone for a spherical isotropic model. Lower energy orbits, however, experience a capture rate several orders of magnitude lower than a full spherical loss cone. Despite its efficiency close to the black hole, it requires a chaotic orbit fraction that may not exist in real galaxies or in N-body generated models (Holley-Bockelmann et al. 2002).   \\begin{figure} \\epsfig{file=fig1.ps, height=3in, width=7in} \\caption{Planar Centrophilic orbits in a triaxial potential with no SMBH. The left panel is a resonant 7:6 boxlet, the middle panel is a 4:3 resonant boxlet (a.k.a. pretzel),  and the right panel is a non-resonant box. These orbits will pass through the center of the potential. The box orbit, in fact, has no net angular momentum and can pass formally through zero. In general, angular momentum is not conserved for any orbit in a triaxial potential.} \\label{fig:orbitsnobh} \\end{figure}  In this paper we investigate a new loss cone refilling mechanism, uniquely present in any non-axisymmetric galaxy. Since orbits in a triaxial potential do not conserve angular momentum, stars can stream rapidly though angular momentum space. This organized flow can occur for any orbit, regular or chaotic, as a consequence of the lack of symmetry in the system which prevents angular momentum from being a constant of the motion. Since any star can flow through angular momentum space, it can be a highly efficient means of transporting stars through zero angular momentum. Even if only a very small fraction of these stars are present in the loss cone, it may be enough to keep it filled.     Since most galaxies are so far from spherical, it is important to determine how more realistic galaxy models effect the structure and time evolution of the loss cone. This is the first paper in a series that will investigate the behavior of the loss cone in general for non-spherical galaxies. Here, we quantify how angular momentum flow refills the 'capture and inspiral' loss cone in a triaxial potential and allows compact objects to be more efficiently captured onto orbits that emit significant gravitational radiation. We demonstrate this mechanism in section 3 and determine the fraction of orbits that can stochastically participate in refilling the capture and inspiral loss cone in an N-body generated triaxial galaxy model. We also explore the astrophysical consequences of this angular momentum flow on the EMRI rate, the SMBH growth rate, and the final parsec problem in sections 4-6. We find that this process alone can replenish the capture and inspiral loss cone as fast as it empties, resulting in an EMRI rate in the Milky Way that is 100 times larger than the canonical estimate for a spherical, isotropic Milky Way bulge model.     Though we first consider angular momentum flow for all orbits in a triaxial potential, more traditionally diffusive processes can refill a loss cone as well. We include dynamical friction (the first order Fokker-Planck term) and 'kicks' (2nd-order Fokker-Plank diffusion) using a hybrid SCF/Fokker-Planck code in our next paper. We are also generating a live $10^7$ particle triaxial model to study the time-dependent, non-equilibrium loss cone refilling rate, and we are using these models to explore the decay of binary supermassive black holes explicitly. ",
        "conclusions": "At relatively large radii, a significant fraction of orbits in a triaxial galaxy stream so rapidly in angular momentum space that they replenish a loss cone as fast as it is depleted. Hence, we have shown that, in principle, loss cones in a triaxial galaxy may always be full at all energies. Nearly every orbit in a triaxial potential participates in this rapid angular momentum flow, reshuffling about $5\\%$ of the phase space content near $r_{\\rm inf}$ over 1000 dynamical times, or on a timescale of $10^7$ years for the Milky Way. Well inside $r_{\\rm inf}$ the potential is rounder, and most orbits that make significant changes in angular momentum are chaotic. We found that the loss cone can be replenished through this process even deep in the potential, where the figure is nearly spherical.   When we consider angular momentum flow for orbits triaxial potential, we find that the capture rate is several orders of magnitude higher than spherical astrophysical capture estimates at nearly every radius. If our astrophysical white dwarf capture rate for the Milky Way of $10^{-5}$ per year is correct, it implies that up to $20 \\%$ of our SMBH could have been grown from white dwarf accretion alone. We emphasize that the fractional number of white dwarfs accreted is still quite small: though about $25\\%$ of the white dwarf population passes through pericenter in a orbital time (by definition), only about 1 in 100000 are captured by the SMBH in our model.   Since dark matter halos are thought to be triaxial in shape, then if there is dark matter in galactic central cusps, this angular momentum flow can also move dark matter onto an analogous, though much narrower, capture loss cone. Some estimates indicate that the density of dark matter dominates the stellar potential at $0.001$ parsec (e.g. Bertone \\& Merritt 2005). If dark matter is less effective at scattering, particularly if it is not self-interacting, triaxiality may be an excellent mechanism to drive significant amounts of dark matter into the SMBH.    The capture and inspiral rate for an eternally full loss cone in a spherical galaxy has often been used as an upper limit, with the argument that a full loss cone cannot be more full. Note, however, that since triaxial galaxies have many more centrophilic orbits, the distribution function of a full triaxial loss cone can be 'supersaturated' compared to the full loss cone for a spherical isotropic galaxy. This could occur even the orbits were only able to change angular momentum via 2-body processes.   The surprisingly large replenishment rates imply large tidal disruption rates as well, since the tidal loss cone width is about 1000 times larger than the capture and inspiral loss cone. This suggests tidal disruption rates as high as $1 \\times 10^{-3}$ stars per year in a triaxial system with the same $\\rho(r)$ as M32.  In fact, this tidal disruption rate is consistent with what is inferred by the large amplitude X-ray outbursts observed at the centers of galaxies (Donley et al. 2002). Our inferred tidal disruption rate is rather large when our model is scaled to the Milky Way bar; however, it may be true that the angular momentum reshuffling process is less efficient in rotating bar potentials. Our numerical models also suggest that captures and tidal disruption events can originate from radii much larger than what is typically assumed; most centrophilic orbits in our model had apocenters much larger than 100 pc and though each centrophilic orbit contributes very little to the total rate, there are enough centrophilic orbits in our triaxial model to influence the capture or tidal disruption rate at large radii.  We have shown that triaxiality can increase the capture rate of single stars, but it may be even more important for binary systems. Binaries are funneled from large distances to the SMBH along centrophilic orbits in a triaxial galaxy. A close pericenter pass tidally separates the binary, with one member of the binary captured by the SMBH, and one ejected at high velocity. Since tidal separation can occur without significant energy loss, the typical captured star from a binary tidal separation can have an apocenter hundreds of times larger than a single stellar capture; this may make binary interactions a much more efficient capture mechanism (Miller et al. 2005), leading to a yet higher EMRI rate. In addition, the binary EMRI signal itself will be unique. When capture occurs at such large separations from the SMBH, there is ample time for the orbit to circularize by the time it becomes detectable by LISA. Hence, there may be two classes of EMRI: the high eccentricity inspiral of a single star, and the zero eccentricity inspiral of a tidally separated binary. Comparing the ratio of the signals can probe the structure, stellar content, and recent kinematic history of the central regions of galaxies."
    },
    "0601/astro-ph0601246_arXiv.txt": {
        "abstract": "Observations of galaxy clusters show that the intracluster medium (ICM) is likely to be turbulent and is certainly magnetized. The properties of this magnetized turbulence are determined both by fundamental nonlinear magnetohydrodynamic interactions and by the plasma physics of the ICM, which has very low collisionality. Cluster plasma threaded by weak magnetic fields is subject to firehose and mirror instabilities. These saturate and produce fluctuations at the ion gyroscale, which can scatter particles, increasing the effective collision rate and, therefore, the effective Reynolds number of the ICM. A simple way to model this effect is proposed. The model yields a self-accelerating fluctuation dynamo whereby the field grows explosively fast, reaching the observed, dynamically important, field strength in a fraction of the cluster lifetime independent of the exact strength of the seed field. It is suggested that the saturated state of the cluster turbulence is a combination of the conventional isotropic magnetohydrodynamic turbulence, characterized by folded, direction-reversing magnetic fields and an Alfv\\'en-wave cascade at collisionless scales. An argument is proposed to constrain the reversal scale of the folded field. The picture that emerges appears to be in qualitative agreement with observations of magnetic fields in clusters. ",
        "introduction": "Clusters of galaxies are vast and varied objects that have long attracted the attention of both observers (in recent decades spectacularly aided by X-ray and radio telescopes) and theoreticians. The observed properties of clusters have proved far from easy to explain as new data has confounded many old theories.  The overall budget of the cluster constituents is roughly as follows: $\\sim75\\%$ of cluster mass is dark matter, whose sole function is assumed to be to provide the gravitational well, $\\sim20\\%$ of cluster mass is the diffuse X-ray-emitting plasma (the intracluster medium, or ICM), while the galaxies have an all but negligible mass. The plasma that makes up the ICM is made of hot and tenuous ionized hydrogen: temperatures are in the range of $1-10$~keV, number densities $\\sim 10^{-1}-10^{-3}~\\text{cm}^{-3}$, with colder, denser material found in the cool cores and hotter, more diffuse one in the outlying regions.\\footnote{The cool cores occupy roughly the central 100~kpc, compared with the overall cluster size of order $1$~Mpc. Such a structure is, in fact, characteristic of the older type of clusters known historically as cooling-flow clusters (the more observationally popular members of this group are Hydra~A, Centaurus and Perseus), while younger, less dynamically relaxed clusters (e.g., Coma) have flatter density and temperature profies. A good summary of the parameters in a number of cooling-flow clusters and the relevant references can be found in Ref.~\\onlinecite{Ensslin_Vogt_cores}.}  Necessarily, the first observations and the first theoretical models of clusters concerned what one might call large-scale features, such as the overall profiles of mass and temperature, the structure formation and the role of the central objects. As observations increased in accuracy and resolution, the ICM was revealed to be much richer than simply a dull cloud of X-ray glow smoothly petering out with distance from the center. A great panoply of features has been detected: bubbles, filaments, ripples, edges, shocks, sound waves, etc., as well as very chaotic density, temperature and abundance distributions.\\cite{Fabian_etal_Perseus1,Fabian_etal_Perseus3,Fabian_etal_Centaurus,Schuecker_etal,Churazov_etal_Perseus1} It is particularly the presence of chaotic fields and the evidence that this chaos exists in a range of scales\\cite{Schuecker_etal,Vogt_Ensslin1,Vogt_Ensslin2} that makes one expect that ICM, like so many other astrophysical plasmas, is in a turbulent state. It is essential to know the properties of this turbulence in order to predict the current and future statistical measurements of the plasma and magnetic fields (spectra, correlation and distribution functions, etc.) and to model correctly the transport processes in the ICM that determine, for example, the overall temperature profiles.\\cite{Voigt_Fabian,Dennis_Chandran}  We shall assume that the physics of small scales is, at least to some degree, independent of large-scale circumstances and that we can, therefore, gain some useful understanding of the turbulence in clusters by ignoring large-scale features and considering a homogeneous subvolume of the ICM. In what follows, after reviewing briefly what is known about fluid motions (\\secref{sec_turb}) and magnetic fields (\\secref{sec_mag}) in clusters, we describe, mostly in qualitative terms, what we consider to be the essential aspects of the small-scale physics of the ICM (for plasma physics, see \\secref{sec_plasma}). This will lead us to a tentative overall picture of the structure of the ICM turbulence (\\secref{sec_sat}) as well as of the origin of its magnetic component (\\secref{sec_dynamo}). We must emphasize that the current state of the debates on the nature of turbulence in clusters (and, indeed, on its very existence\\cite{Fabian_etal_Perseus2}) is such that even a fundamental, conceptual view of the problem has not yet been agreed, and, therefore, a quantitative theory --- even an incomplete one such as exists for hydrodynamic turbulence --- remains a matter of future work. ",
        "conclusions": ""
    },
    "0601/astro-ph0601593_arXiv.txt": {
        "abstract": "{ {We study the large-scale angular correlation signatures of the Cosmic Microwave Background (CMB) temperature fluctuations from WMAP data in several spherical cap regions of the celestial sphere, outside the Kp0 or Kp2 cut-sky masks.} {We applied a recently proposed method to CMB temperature maps, which permits an accurate analysis of their angular correlations in the celestial sphere through the use of normalized histograms of the number of pairs of such objects with a given angular separation versus their angular separation. The method allows for a better comparison of the results from observational data with the expected CMB angular correlations of a statistically isotropic Universe, computed from Monte Carlo maps according to the WMAP best-fit $\\Lambda$CDM model.} {We found that the, already known, anomalous lack of large-scale power in full-sky CMB maps are mainly due to missing angular correlations of quadrupole-like signature. This result is robust with respect to frequency CMB maps and cut-sky masks. Moreover, we also confirm previous results regarding the unevenly distribution in the sky of the large-scale power of WMAP data. In a bin-to-bin correlations analyses, measured by the full covariance matrix $\\chi^2$ statistic, we found that the angular correlations signatures in opposite Galactic hemispheres are anamalous at the 98\\%--99\\% confidence level.}}   \\titlerunning{On the CMB large-scale angular correlations} \\authorrunning{Bernui et al.} ",
        "introduction": "In recent years, after the COBE mission, experiments of observational cosmology evolved both in the accuracy  and in the angular resolution of the Cosmic Microwave Background (CMB) measurements (see, e.g. Bersanelli et al.~\\cite{Bersanelli} and references therein). However, such experiments were designed to map small patches of the celestial sphere and therefore the study of its properties concerning the large-scale angular correlations was not possible.  This situation changed with the Wilkinson Microwave Anisotropy Probe (WMAP), which produced full-sky CMB maps in five frequency bands (termed K, Ka, Q, V, and W maps), although containing different amounts of contaminations from our galaxy (Bennett et al.~\\cite{WMAPa}). This superb quality data set motivated a number of detailed studies concerning the CMB properties. As a result, several recent works have reported claims regarding some anomalies found in these data. These anomalies include the low-order multipole values (Bennett et al.~\\cite{WMAPa}; Efstathiou~\\cite{GE}; Tegmark et al.~\\cite{TOH}; Gazta\\~naga et al.~\\cite{Gaztanaga}; Slozar et al.~\\cite{Slosar}; O'Dwyer et al.~\\cite{ODwyer}), the alignment of some low-order multipoles (Tegmark et al.~\\cite{TOH}; de Oliveira-Costa et al.~\\cite{OTZH}; Eriksen et al.~\\cite{Eriksen04b}; Bielewicz et al.~\\cite{Bielewicz05}; Land and Magueijo~\\cite{Land05}), and an unexpected asymmetric distribution on the sky of the large-scale power of CMB data (Eriksen et al.~\\cite{Eriksen04a},~\\cite{Eriksen05a}; Hansen et al.~\\cite{Hansen04a},~\\cite{Hansen04b},~\\cite{Hansen04c}; Jaffe et al.~\\cite{Jaffe}; Schwarz et al.~\\cite{Schwarz}; Copi et al.~\\cite{Copi05}).  In this work we focus on the study of the large-scale angular correlations (hereafter termed Angular Correlations Signatures ACS) of the CMB temperature maps from WMAP data. Since the vast majority of the cosmological information is contained in the two-point temperature correlation function, it seems natural to study the ACS by looking at such function (actually, an equivalent of it). Our analyses of the ACS in CMB maps, for data outside the region defined by the Kp0 or Kp2 WMAP masks, allow to perform a close inspection of two interesting issues. First, we investigate --through a full-sky map analysis-- the large-scale power in WMAP data, and its connection with the low quadrupole moment issue. Second, we study --through a partial-sky coverage analysis-- the possible uneven distribution on the celestial sphere of the CMB large-scale power. The significance of our results are evaluated by comparing these results against those obtained from Monte Carlo CMB maps produced with the WMAP best-fitting $\\Lambda$CDM model properties (Hinshaw et al.~\\cite{Hinshaw03b}). Finally, using a $\\chi^2$ statistics we assess the confidence level of the North/South asymmetry in WMAP data compared to 1\\,000 Monte Carlo realizations of CMB skies. ",
        "conclusions": "Here we have shown through the full-sky analyses of the ACS of the WMAP data (outside the Kp0 or Kp2 sky cuts) that the `anomalous' lack of power at large angular scales is mainly explained by the low quadrupole-moment value (see Efstathiou~\\cite{GE}, Gazta\\~naga et al.~\\cite{Gaztanaga}, and Slosar et al.~\\cite{Slosar} for more detailed discussions; see also Bielewicz et al.~\\cite{Bielewicz04} for complete calculations, in particular in similar band maps as here). We confirmed this result by analysing also the WMAP data after removing the quadrupole component, obtaining a better (but not exactly) agreement with what is expected in a $\\Lambda$CDM model with zero quadrupole.  We also performed partial-sky coverage analyses and showed that there is significative evidence, at the $98\\%-99\\%$ CL, of an anomalous distribution of the CMB power in opposite Galactic hemispheres. In fact, in less than 2\\% of the Monte Carlo CMB maps analyzed, data in opposite hemispheres present analogous bin-to-bin correlations as those shown by the $90^{\\circ}$-caps with WMAP data. More interestingly, the excess of correlations in the comparison of opposite antipodal caps is manifestly observed in the South/North $90^{\\circ}$-caps, but it is not seen neither in the examination of the antipodal $45^{\\circ}$- or $60^{\\circ}$-caps.  The fact that the South/North asymmetry persists at a significant level in several frequency bands (the Q, V, and W CMB maps) and with different sky cuts (the Kp0 and Kp2 masks) excludes a simple explanation based on systematic effects and foregrounds (see Hinshaw et al. \\cite{Hinshaw03a}; Eriksen et al. \\cite{Eriksen04a},~\\cite{Eriksen04b}; Hansen et al.~\\cite{Hansen04b})."
    },
    "0601/astro-ph0601070_arXiv.txt": {
        "abstract": "The H$\\alpha$ line emission formation region in the circumstellar disc of $\\alpha$~Eri is: a) extended with a steep outward matter density decline during low H$\\alpha$ emission phases; b) less extended with rather constant density distribution during the strong H$\\alpha$ emission. The long-term variation of the H$\\alpha$ emission has a 14-15 year cyclic \\bbe phase transition. The disc formation time scales agree with the viscous decretion model. The time required for the disc dissipation is longer than expected from the viscous disc model. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601300_arXiv.txt": {
        "abstract": "We analyze {\\it Spitzer} and Magellan observations of a star forming core near IRS-2 in the young cluster NGC 2264.  The submillimeter source IRAS 12 S1, previously believed to be an intermediate mass Class 0 object is shown to be a dense collection of embedded, low mass stars. We argue that this group of stars represents the fragmenting collapse of a dense, turbulent core, based on a number of indicators of extreme youth.  With reasonable estimates for the velocity dispersion in the group, we estimate a dynamical lifetime of only a few x 10$^{4}$ years. Spectral energy distributions of stars in the core are consistent with Class I or Class 0 assignments.  We present observations of an extensive system of molecular hydrogen emission knots.  The luminosity of the objects in the core region are consistent with roughly solar mass protostars. ",
        "introduction": "Over the past 20 years, a relatively consistent picture of low-mass star formation has emerged.  In this scenario, a dense molecular core becomes gravitationally unstable, collapses, and begins the process of star formation.  We know, however, that real star formation is significantly more complex.  In particular, stars do not form in isolation, and interactions between different protostars may be quite important in their development.  Recent numerical simulations by \\citet{Bate03}, for example, show that star formation in a turbulent core is a highly dynamic, interactive process. Because this early formation activity occurs in regions of very high obscuration, detailed investigations have been difficult given the limited sensitivity and angular resolution of available far infrared capabilities.  The observational situation has dramatically improved with the launch of {\\it Spitzer}.   NGC 2264 has long been one of the touchstones in star formation studies.  Beginning with the work of \\citet{Herbig54} and \\citet{Walker56}, studies at all available wavelengths have been conducted in this extremely young cluster.  At a distance of 800 pc, NGC 2264 is close enough for detailed investigations of the stellar population down to very low masses.  While more distant than some of the other well studied star formation regions like $\\rho$ Ophiuchi or the Orion molecular cloud, NGC 2264 displays some of the best examples of the interaction of the interstellar medium with the young stars.  Investigators at optical wavelengths \\citep*{Sung97, Flaccomio99, Park00, Rebull02, Dahm05} have fit their data to theoretical isochrones and find typical ages in the 1-3 Myr range for the young stellar population.  They also find, however, that there is a wide range in derived ages for the individual stars, going from 10$^{5}$ to nearly 10$^{7}$ years. The region contains numerous IRAS sources, many of which have been classified as Class I protostars \\citep{MLY89}. Moreover, the presence of Herbig-Haro objects, molecular outflows \\citep{MLS88, Wolf-Chase95}, and Class-I objects \\citep{Wolf-Chase03} attest to continuing, active star formation.  One of these Class I sources, IRS-2 (listed as IRAS 06382+0939 in the IRAS Point Source Catalog) is located 17\\arcmin \\/ south of the O7 star S Mon and 7\\arcmin \\/ north of the Cone Nebula in a dense clump of molecular gas. It was identified as a source of far-infrared radiation by \\citet{Sargent84} and by \\citet{Schwartz85}. Using the Kuiper Airborne Observatory, \\citet{Cohen85} found a second peak of far infrared radiation 2\\arcmin \\/ to the southwest. \\citet{Castelaz88} mapped the region at near infrared wavelengths and associated IRS-2 with a pair of Red Nebulous Objects (RNO-E and RNO-W).  IRS-2 was designated as IRAS 12 by \\citet{MLY89} in their IRAS survey of the NGC 2264 region, and they classified it as a likely Class I protostar.  This classification was supported by the association of IRAS 12 with a high velocity molecular outflow (NGC 2264 Flow D). More recently, this region has been mapped at higher angular resolution at submillimeter wavelengths by \\citet{Williams02} and \\citet{Wolf-Chase03}.  Those maps show that the IRAS source itself is weak at 450 and 850 $\\mu$m, but that it is part of a complex collection of submillimeter cores.  The latter authors argue that these cores have a range of evolutionary stages.  \\citet{Teixeira06} presented {\\it Spitzer} observations of NGC 2264 IRS-2 region. They found that bright 24 $\\mu$m sources were aligned in a spoke-like pattern appearing to emanate from IRAS 06382+0939 and coinciding with submillimeter dust filaments. Additionally they found that the nearest neighbor spacing between these sources had a peak in the distribution of approximately 20\\arcsec, very close to the Jeans length of the cloud.  They argued that the 24 $\\mu$m sources represented collapsing, fragmenting filaments.  The brightest of these cores in the submillimeter is IRAS 12 S1, located near the position of the \\citet{Cohen85} secondary peak. IRAS 12 S1 also corresponds to the millimeter source D-MM1 in the survey of \\citet*{Peretto06}. \\citet{Wolf-Chase03} suggested that IRAS 12 S1 is an intermediate mass Class 0 protostar.    In this paper we present {\\it Spitzer} and Magellan observations that demonstrate that S1 is, in fact, a complex of extremely young objects. ",
        "conclusions": "The simulations of \\citet{Bate03} and \\citet{Bate05} follow the collapse of a Jeans unstable turbulent core.  They show that star formation commences on time scales comparable to the global free-fall time, and it proceeds by forming dense cores that fragment into stars and brown dwarfs.  Dynamical interactions are important in the development of the ensemble.  More recently, \\citet{Kurosawa04} have simulated the appearance of the \\citet{Bate03} simulations in the {\\it Spitzer} bands.  We will argue that the S1 stars represent just such a fragmenting core. For this hypothesis to be plausible, it will be necessary to demonstrate the extreme youth of the objects.  Figure 4 is a $K_s$ PANIC image of the S1 region with the sources identified.  Overlaying the image are the 850 $\\mu$m contours from \\citet{Wolf-Chase03}.   As shown in the figure, the \"micro-cluster\" is located precisely at the peak of the IRAS-12 S1 submillimeter source mapped by \\citet{Wolf-Chase03} (also denoted Core C by \\citet{Williams02}). The entire core has a composite spectral energy distribution (SED) that peaks in the far-infrared. \\citet{Wolf-Chase03} derived an envelope mass of 17.6 $M_\\sun$ within a 28\\arcsec \\/diameter. Based on the ratio of submillimeter to bolometric luminosity, they classified S1 as an intermediate-mass Class 0 protostar. \\citet{Williams02} observed the region in $HCO^+$ and detected an emission line with a strong, red-shifted absorption feature. They interpreted this spectral signature as indicative of infall.  Hence, both continuum and line observations suggest that IRAS 12 S1 is undergoing active collapse. The submillimeter contours of \\citet{Wolf-Chase03} show that S1 is somewhat extended but unresolved at both 850 and 450 \\micron.  Rather than being a single protostar, however, it is clear from the new shorter wavelength observations that S1 is made up of a cluster of discrete sources.  Figure 5 shows the SEDs of the sources encompassed by the submillimeter peak. Except for \\#34, which appears to be a unreddened foreground star, all the sources are steeply rising to longer wavelengths and can be classified as Class 0 or Class I. While all these stars appear to have significant amounts of circumstellar matter, it is not possible to unambiguously assign the relative contributions of the stars to the 24 $\\mu$m flux since in many cases the sources are badly confused.  For example, the flux given to \\#38 really is a composite that can include contributions from \\#41 and others.  The SEDs in S1 are markedly different from those just outside the dense submillimeter core. Figures 6 and 7 show the SEDs for the regions 1\\arcmin \\/ West and 1\\arcmin \\/ Northeast of S1, respectively. There are a number of flat SEDs in those two adjacent regions, but the majority of the detected objects have SEDs consistent with highly reddened low-mass stars.  A second indication of the youth of the micro-cluster is the presence of extensive molecular hydrogen flows.  In the S1 core itself, several knots of molecular hydrogen are present, for example, between sources \\#47 and \\#49 and between sources \\#40 and \\#37. The most extensive molecular hydrogen line emission is away from the core, however.  A north-south series of knots (indicated by the arrows in Figure 3) is quite prominent in the continuum subtracted $H_2$ image. The morphology is reminiscent of bipolar Herbig-Haro flows. However, we find no evidence for any candidate sources at 24 or 70 $\\mu$m other than possibly \\#26.  It is a relatively weak flat-spectrum source with no significant emission at 70 $\\mu$m.  It is possible that one of the Class I sources in the micro-cluster is the source of the $H_2$ emission, although that interpretation would be difficult to reconcile with the North-South geometry of the knots.  That possibility will likely require proper motion studies of the knots to prove conclusively.  A number of the sources appear to have opaque dust lanes that are likely caused by circumstellar disks.  The best example is \\#44, which has a distinctly bipolar appearance.  Moreover, this source is associated with molecular hydrogen emission that emanates from the \"pole\" and terminates in a bright knot 7\\arcsec to the northeast.  An upper limit to the age of the S1 group can be determined from dynamical arguments.  The characteristic non-thermal velocity dispersion of the gas in the core has been measured by \\citet{Williams02} to be 0.7 - 1.0 km s$^{-1}$. \\citet{Williams02} also identified five other submillimeter cores in the main S1 complex.  The measured core to core velocity dispersion is 0.9 km s$^{-1}$. If we adopt a 0.7 km s$^{-1}$ 1-dimensional dispersion for the stars, the dispersal time in the S1 micro-cluster is only 40000 years for a 10\\arcsec radius. The 0.7 km s$^{-1}$ velocity dispersion may, in fact, be an underestimate if we consider velocity dispersions found in other star formation regions. For example, \\citet{Jones88} measured a 1-dimensional dispersion of 2 km s$^{-1}$ for the Orion nebula stars.  The simulations of turbulent core collapse exhibit many examples of dynamical interactions between the forming stars. In particular, the lower mass members often attain a high enough velocity to be ejected from the core region. \\citet{Bate05} found 1-dimensional velocity dispersions between 1 and 2.5 km s$^{-1}$, depending on initial conditions. Hence, our use of 0.7 km s$^{-1}$ for the velocity dispersion is likely to be conservative.  What can be said about the masses of the members of the micro-cluster?  Unfortunately, the wide range of parameters afforded by extinction (both disk and envelope), inclination, and accretion rate makes fitting of individual sources to evolutionary models problematic.  Limits can be established, however, if we consider the ensemble properties of the micro-cluster.  \\citet{Wolf-Chase03} derive an envelope mass of 17.6 $M_{\\sun}$ and a bolometric luminosity of 107.5 $L_{\\sun}$, leading to their association of the source with an intermediate mass protostar.  The {\\it Spitzer} observations show, however, that the luminosity is actually shared by perhaps a dozen lower-mass objects.  For extremely young objects, the far-infrared and submillimeter luminosity is expected to be a valid measure of the bolometric luminosity since they will still be surrounded by envelopes that intercept the bulk of the short wavelength emission.  Applying the pre-main sequence tracks of \\citet{Siess00}, a star with a luminosity of 10 $L_{\\sun}$ corresponds to a mass of only 0.9 $M_{\\sun}$ at an age of 10$^{5}$ years.  Thus, we may be seeing the fragmentation of a molecular core into solar or lower mass stars."
    },
    "0601/astro-ph0601136_arXiv.txt": {
        "abstract": "The absolute visual magnitudes of three Type~IIb, 11~Type~Ib and 13~Type~Ic supernovae (collectively known as stripped--envelope supernovae) are studied by collecting data on the apparent magnitude, distance, and interstellar extinction of each event.  Weighted and unweighted mean absolute magnitudes of the combined sample as well as various subsets of the sample are reported.  The limited sample size and the considerable uncertainties, especially those associated with extinction in the host galaxies, prevent firm conclusions regarding differences between the absolute magnitudes of supernovae of Type~Ib and Ic, and regarding the existence of separate groups of overluminous and normal-luminosity stripped--envelope supernovae.  The spectroscopic characteristics of the events of the sample are considered. Three of the four overluminous events are known to have had unusual spectra.  Most but not all of the normal luminosity events had typical spectra.  Light curves of stripped--envelope supernovae are collected and compared.  Because SN~1994I in M51 was very well observed it often is regarded as the prototypical Type~Ic supernova, but it has the fastest light curve in the sample.  Light curves are modeled by means of a simple analytical technique that, combined with a constraint on $E/M$ from spectroscopy, yields internally consistent values of ejected mass, kinetic energy, and nickel mass. ",
        "introduction": "In a recent paper \\citep[][hereafter R02]{richardson02} we carried out a comparative study of the absolute magnitudes of all supernovae (SNe) in the Asiago Catalog.  Because of the large number of SNe in the sample, we did not attempt to estimate the extinction of each SN in its host galaxy, and we did not assign uncertainties to individual SN absolute magnitudes.  In this paper we look more closely at the absolute--magnitude distributions of stripped--envelope supernovae (SE~SNe) by assigning uncertainties to each of the quantities that enter into the absolute--magnitude determination, including host galaxy extinction. By SE~SNe we mean SNe of Types IIb, Ib and Ic. (The subset containing only SNe~Ib and Ic is referred to as SNe~Ibc.)  The progenitors of SE~SNe are stars that have lost most or all of their hydrogen envelopes. This can happen by strong winds such as in Wolf Rayet stars or through mass transfer to a companion star such as in Roche--lobe overflow or a common--envelope phase.  The light curves of SE~SNe are powered by the radioactive decay of $^{56}$Ni, so the absolute magnitudes are closely related to the ejected $^{56}$Ni masses, and in turn to the stellar progenitors and explosion mechanisms. Since the discovery of the apparent association of GRB980425 with the peculiar Type~Ic SN~1998bw \\citep{galama99}, and the confirmation by spectra of a SN~1998bw--like event associated with GRB030329 \\citep{garnavich03}, SE~SNe have become of intense interest in connection with GRBs. Contamination of high--redshift samples of Type~Ia supernovae by SE SNe also is an important issue \\citep{homeier05}.  We also study the $V$--band light curves (LCs) of SE~SNe.  Data were collected from the literature for two SNe~IIb, seven SNe~Ib and 11 SNe~Ic including three that had unusually broad and blueshifted absorption features in their spectra (we will refer to these as hypernovae). The LCs show considerable diversity in peak brightness, in the width of the peak, and the slope of the late--time tail. In order to relate all of the LCs to total ejected mass, ejected nickel mass, and kinetic energy in an internally consistent way, we fit the data to a simple LC model.  The peak absolute--magnitude data and analysis are described in \\S2. The light curve data and model fits are presented in \\S3.  A brief summary is given in \\S4. ",
        "conclusions": "We have used the available data to characterize the absolute--magnitude distributions of the SE~SNe in the current observational sample.  Most SE~SNe have a ``normal'' luminosity, which at $M_V=-17.77 \\pm 0.06$ is about a magnitude and a half dimmer than SNe~Ia.  One sixth of the current sample of SE~SNe are overluminous, i.e., more luminous than SNe~Ia, but these are strongly favored by observational selection so the true fraction of SE~SNe that are overluminous is much lower than one sixth.  The small size of the sample and the considerable absolute magnitude uncertainties, especially those due to host galaxy extinction, still prevent an absolute magnitude difference between SNe~Ib and Ic from being firmly established.  Three of the four (or four of the five, counting SN~1999as) overluminous SE~SNe are known to have had unusual spectra; a few of the normal--luminosity SE~SNe also had unusual spectra.  Much more data on SE~SNe are needed in order to better determine the absolute--magnitude distributions, and to correlate absolute magnitudes with spectroscopic characteristics.  Absolute light curves in the $V$ band (some fragmentary) are available for two SNe~IIb, seven SNe~Ib, and 12 SNe~Ic including three hypernovae.  Two of the SNe~Ib, SNe~1984L and 1991D, have LCs that are quite different from those of the others.  The light curves of the SNe~Ic are rather diverse.  The light curve of SN~1994I, often considered to be a typical SN~Ic, actually is the most rapidly declining light curve in the SN~Ic sample.  The simple analytical light--curve model was applied to two SNe~IIb, five SNe~Ib and 10 SNe~Ic.  Instead of assuming a kinetic energy, such as the canonical one foe, an $E_k$/$M_{ej}$ ratio was estimated on the basis of spectroscopy, and the model fits then produced internally consistent values of $E_k$, $M_{ej}$, and $M_{Ni}$.  Reasonably good fits were obtained for the two SNe~IIb and three of the five SNe~Ib.  The slowly decaying tail of the SN~1984L LC and the rapid decline from the peak of the SN~1991D LC cannot be fit by the model.  With the exception of the hypernova SN~1997ef, reasonable fits were obtained for the SNe~Ic, with a considerable range in the parameter values.  As expected, the hypernovae have high $E_k$ and somewhat high $M_{ej}$, but only one of the three has high $M_{Ni}$.  Our values of $M_{ej}$ (and $E_k$) tend to be lower than those obtained by others by means of numerical LC calculations, while our values for $M_{Ni}$ are a little higher.  Some of the other numerical calculations are based on an assumed canonical value of $E_k =1$ foe. Because our spectroscopic constraint on $E_k$/$M_{ej}$ together with our LC model leads to lower $E_k$, and lower $E_k$ makes the LC peak dimmer, we need slightly higher $M_{Ni}$ values to make the LC peaks as bright as observed.  The diversity among SN~Ic LCs is of special interest in connection with the ongoing search for SN signals in GRB afterglows.  In a forthcoming paper we will apply the same modeling technique used in this paper to the putative SN bumps in GRB afterglow light curves. This will enable us to infer internally consistent SN parameter values for comparison with the results of this paper, and to investigate the issue of whether the GRBs and the associated SNe are coincident in time or whether, as in the supranova model, the SN precedes the gamma-ray burst.  \\vspace{7mm}  We would like to thank Rollin C. Thomas for helpful comments and suggestions. Support for this work was provided by NASA through grants GO-09074 and GO-09405 from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS 5-25255.  Additional support was provided by National Science Foundation grants AST-9986965 and AST-0204771.  \\clearpage"
    },
    "0601/gr-qc0601061_arXiv.txt": {
        "abstract": "Neutron stars are discussed as laboratories of physics of strong gravitational fields. The mass of a neutron star is split into  matter energy and gravitational field energy contributions. The energy of the gravitational field of neutron stars is calculated with three different approaches which give the same result. It is found that up to one half of the gravitational mass of maximum mass neutron stars is comprised by the gravitational field energy. Results are shown for a number of realistic equations of state of neutron star matter. ",
        "introduction": "Gravity makes all selfgravitating fluid bodies spherical. Thus the spherical symmetry of non-rotating neutron stars, which are real objects, is their intrinsic property, not an idealization. Neutron stars are relativistic objects, with mass of order $1M_{\\odot}$ enclosed within 10-20 km.  Gravity of such a compact mass distribution is strong enough to require relativistic description. In the following we focus on gravitational field contribution to the neutron star masses in General Relativity.  The problem of the energy distribution of gravitational field remains an unsettled issue in general relativity \\cite{Cooperstock} mainly because of the lack of proper energy-momentum tensor for gravitational field. However, for static, spherically symmetric space-time no controversy exists  and relevant matter and gravity energies can be unambiguousely defined \\cite{MTW, Weinberg}. Our main concern here is to study if neutron stars could shed some new light on the problem of gravitational field energy and its localization.  To calculate both matter and gravitational field  contributions to the neutron star mass we solve the Einstein's field equations for the static, spherically symmetric metric and a given equation of state of dense matter. The metric is \\begin{equation} {ds}^2 = e^{\\nu }c^2{dt}^2 - e^{\\lambda }{dr}^2- r^2\\left({d\\theta}^2 + {\\sin{\\theta}^2}{d\\phi}^2 \\right), \\label{metrics} \\end{equation} where $\\nu$, $\\lambda$ are functions of the radial coordinate $r$.   Dense matter in neutron stars is thought to be in a liquid state, hence it is described by the perfect fluid energy-momentum tensor, \\[T^{\\mu \\nu}=(\\rho c^2 + P)u^{\\mu}u^{\\nu}-Pg^{\\mu \\nu},\\] where $\\rho c^2$ is the energy denisty, $P$ is the pressure and $u^{\\mu}$ is the four velocity. For the neutron star matter both energy density and pressure are uniquely determined by the baryon number density, $\\rho=\\rho(n_B), P=P(n_B)$, hence the equation of state is of barotropic form, $P=P(\\rho)$. In this paper we employ a number of realistic equations of state, described in some details in Sect. \\ref{EOSNM}.  Using the Einstein's equations one finds that functions $\\lambda$ and $\\nu$ satisfy ordinary differential equations  \\begin{eqletters} \\label{myeq} \\begin{eqnarray} \\lambda'(r)= {1 \\over r}\\left(1 - e^{\\lambda } + {8\\,\\pi\\,G\\,e^{\\lambda }\\,r^2\\,\\rho(r) \\over c^2}\\right) \\label{me1} \\\\ \\nu'(r)={1 \\over r}\\left(-1 + e^{\\lambda } + {8\\,\\pi\\,G\\,e^{\\lambda }\\,r^2\\,P(r) \\over c^4}\\right) \\label{me2} \\\\ P'(r)=-{1 \\over 2}(P(r) + \\rho c^2)\\nu'. \\label{me3} \\end{eqnarray} \\end{eqletters}  The metric (1) has to be continuous in the whole space. For this to happen, at the stellar boundary, $r=R$, the internal metric has to go over into the external Schwarzschild metric. There is a vacuum outside the star and the matter energy-momentum tensor vanishes. The metric functions at the boundary should satisfy the condition $e^{\\nu(R)}=e^{-\\lambda(R)}=1-2GM/Rc^2$, where $M$ is the mass of the star.  From Eqs.(2) one derives the Tolman-Oppenheimer-Volkoff (TOV) equation that describes  nonrotating compact objects in the hydrostatic equilibrium: \\begin{eqnarray} \\lefteqn{ {dP(r)\\over dr}= - {GM(r)\\rho(r)\\over r^2}\\times } \\nonumber \\\\ & & \\left(1+ {P(r)\\over c^2\\rho(r)}\\right)\\left(1+{4\\pi r^3P(r)\\over c^2M(r)}\\right)\\left(1-{2GM(r)\\over c^2r}\\right)^{-1} \\end{eqnarray} Here we denote by $M(r)$ the integral $M(r)=\\int_0^r\\rho(r')d^3r'$. Usually $M(r)$ is referred to as \"the mass within the radius $r$\". This name is rather misleading, as it suggests that only matter energy density is responsible for building the neutron star mass $M=M(R)$. It is our prime goal here to calculate the gravitational energy contribution to $M$. In the following we shall be using the differential form of the $M(r)$ definition, \\begin{equation} {dM(r)\\over dr}= 4\\pi r^2\\rho(r). \\label{eqM} \\end{equation} Choosing an appropriate equation of state of dense matter \\begin{equation} P = P(\\rho), \\end{equation} and the boundary conditions we  solve this set of equations for different values of central density $\\rho_c$ obtaining the whole family of configurations of compact objects in hydrostatic equilibrium. \\begin{equation} \\rho(r=0)=\\rho_c \\ \\mathrm{and} \\ P(r=R)=0, \\end{equation}  \\begin{figure}[!hbt] \\begin{center} \\includegraphics[height=5.5cm]{Fig1.eps} \\end{center} \\caption{Neutron star masses \\textit{versus} central density for four equations of state.} \\label{fig:ns_konf} \\end{figure} We report here calculations for four equations of state, which we denote as A18, A18+$\\delta$v+UIX*, FPR and MS (see Sect. \\ref{EOSNM}). These four EOS have been chosen from about twenty EOS we studied solely to make the plots clear. The results of our calculations are shown in  Fig. \\ref{fig:ns_konf}.  Numerical values of the neutron stars masses for these equations of states are given in the Table \\ref{tab:masses}. We show parameters corresponding to the standard neutron star mass $M=1.44M_{\\odot}$, and to the maximum mass neutron star, which is different for every equation of state. \\begin{table}[t] \\begin{center} \\begin{tabular}{l c c }  Equation of state   &  Central density [$10^{15}\\,g/cm^3$] & Mass [$M_\\odot$]   \\\\  \\hline A18 & 2.54 & $1.44$   \\\\ A18 & 4.23&$1.68$  \\\\ \\hline  A18+$\\delta$v+UIX* & 1.02 & $1.44$  \\\\ A18+$\\delta$v+UIX* & 2.74 & $2.22$ \\\\ \\hline FPR & 1.36 & $1.44$ \\\\ FPR & 3.38 & $1.80$ \\\\ \\hline MS & 1.01 & $1.44$ \\\\ MS & 2.84 & $2.04$ \\\\  \\hline \\end{tabular} \\caption{Computed neutron star masses. For each equation of state the first line shows the value of the central density of $1.44\\ M_{\\odot}$ neutron star, and in the second row the central density of the maximum mass neutron star and the value of this mass are given. \\label{tab:masses}} \\end{center} \\end{table} ",
        "conclusions": ""
    },
    "0601/astro-ph0601026_arXiv.txt": {
        "abstract": "ASAS is a long term project to monitor bright variable stars over the whole sky.  It has discovered 50,122 variables brighter than ${\\rm V < 14 }$ mag south of declination ${\\rm + 28^{\\circ}}$, and among them 11,099 eclipsing binaries.  We present a preliminary analysis of 5,384 contact, 2,957 semi-detached, and 2,758 detached systems. The statistics of the distribution provides a qualitative confirmation of decades old idea of Flannery and Lucy that W UMa type binaries evolve through a series of relaxation oscillations: ASAS finds comparable number of contact and semidetached systems.  The most surprising result is a very small number of detached eclipsing binaries with periods ${\\rm P < 1 }$ day, the systems believed to be the progenitors of W UMa stars.  As many (perhaps all) contact binaries have companions, there is a possibility that some were formed in a Kozai cycle, as suggested by Eggleton and his associates. ",
        "introduction": "\\label{sect:intro1}  ASAS - All Sky Automated Survey, is a long term project dedicated to detection and monitoring variability of bright stars.  This paper presents the results of several years of observations done at the Las Campanas Observatory with a single instrument: a telescope with the aperture of 7 cm, the focal length of 20 cm, done through a standard V-band filter and a ${\\rm 2K \\times 2K }$ CCD camera with ${\\rm 15 \\mu }$ pixels from Apogee (Pojma\\'nski 1997, 1998, 2000, 2002, 2003, Pojma\\'nski and Maciejewski 2004, 2005, Pojma\\'nski, Pilecki and Szczygiel 2005).  More information about ASAS is provided on the WWW:  \\centerline{ http://www.astrouw.edu.pl/\\~{}gp/asas/asas.html } \\centerline{ http://archive.princeton.edu/\\~{}asas/ }  \\begin{figure*} \\vspace{-0.9\\linewidth} \\includegraphics[width=\\linewidth]{fig1.ps} \\caption{ The distribution of 50,122 ASAS variables in the sky in Galactic coordinates.  The Milky Way is clearly seen, as well as the patches of interstellar extinction.  The equator is shown with a dashed line. The distribution is limited to declination ${\\rm < +28^{\\circ}}$. } \\end{figure*}  The variable stars were discovered quasi - uniformly for declination ${\\rm < + 28^{\\circ} }$, covering almost 3/4 of the full sky. Fig. 1 shows the distribution of ASAS variables in the sky in the Galactic coordinates.  The Milky Way is clearly visible, together with the dust lanes.  The distribution of ASAS stars as a function of V-mag is shown on Fig. 2 for all stars (17,000,000), all variable stars (50,122), and all eclipsing binaries, the latter divided into Eclipsing Contact binaries (EC, 5,384), Eclipsing Semi Detached binaries (ESD, 2,957), and Eclipsing Detached binaries (ED, 2758). The statistics for stars with ${\\rm 8 < V < 12 }$ mag appears to be approximately complete, but the efficiency falls rapidly within the range of ${\\rm 12 < V < 14 }$ mag as the detection limit is approached.  Also, the statistics deteriorates for ${\\rm 8 < V }$ because of saturation  effects.  \\begin{figure} \\includegraphics[width=\\linewidth]{fig2.ps} \\caption{ The distribution of ASAS stars as a function of V-band magnitude. The total number of stars is about 17,000,000.  The total number of variable stars is 50,122.  The total number of eclipsing binaries is 5,384 for contact systems (EC), 2,957 for semi-detached systems (ESD), and 2,758 for detached systems (ED).  Notice, that the efficiency of discovering variable stars declines for ${\\rm V > 12 }$ because the detection limit is approached, and for ${\\rm V < 8 }$ because of saturation effects. } \\end{figure}  All stars were observed for about 5 years, with a small subset for 8 years. Typical number of photometric measurements was several hundred.  The distribution of this number is shown in Fig. 3. The total number of photometric measurements was 2,916,000. As ASAS continues its operation the number of measurements will increase, approximately 100 V-band photometric measurements per year per variable. We intend to continue the project indefinitely, with some upgrades.  While only V-band results are reported here, the I-band photometry was accumulated, and several years of data are already stored on a RAID-5 disk system.  However, it will take another year to process I-band data, well over a Tera-byte.  There are more stars detectable in the I-band, so we expect that the number of variables will more than double.  We are also planning an expansion of ASAS to the northern hemisphere, to fully cover the whole sky.  \\begin{figure} \\includegraphics[width=\\linewidth]{fig3.ps} \\caption{ The histogram of the number of photometric measurements obtained during 5 years of ASAS life.  A small subset of data extends back to 8 years. } \\end{figure}  The results presented in this paper are not final in any sense. Please note, that all data are public domain.  A different classification of binaries can be readily done by whoever feels like verifying and/or correcting our presentation.  This is an observational paper, but section 2 gives a short introduction to the main ideas about the structure and evolution of contact binaries.  Section 3 gives the information about ASAS data, the classification scheme, and some examples of a diversity of light curves.  Section 4 provides simple statistics of ASAS binaries. Finally, in section 5 we make a somewhat speculative discussion based on this observational paper. ",
        "conclusions": "\\label{sect:discuss}  Our conclusion, based on the distribution of eclipse depths (Fig. 10, 11), is that relaxation oscillations, first proposed by Lucy and Flannery, are real.  There is no shortage of binaries corresponding to no thermal contact, with a very different depth of the two eclipses.  Model calculations of the type recently done by Yakut and Eggleton (2005), when combined with ASAS data, should allow a quantitative verification of the theory.  A very large number of contact or near contact systems reveal features never noticed before: the distribution of eclipse depth ratios as shown in Fig. 10 and Fig. 11 has a distinct break around ${\\rm log P \\approx -0.4 }$ and ${\\rm D_s/D_p \\approx 0.7 }$.  We do not speculate on the origin of this feature, but we bring this break to the attention of our readers.  This break is best seen in systems with the most robust classification.  The most surprising result of this paper is presented in Fig. 6. A traditional view for the origin of W UMa contact binaries is to assume that they come from detached binaries of comparable periods.  For the first time we have approximately complete statistics of binaries of all types with orbital periods shorter than one day, and there are very few detached binaries. Obviously, they are more difficult to find than either contact or semi - detached systems. As times goes on the statistics of ASAS detached binaries will improve, and a statistical analysis will tell us if there is a problem with the origin of W UMa stars. At this time the contact systems seem to appear \"out of nowhere\".  In about a year we shall have I-band data for our eclipsing binaries. This will make it possible to be more quantitative about the distribution properties, including the distribution in our Galaxy. Also, it will be much easier to quantify the impression about a shortage of detached binaries and the space density of systems of different types: EC, ESD, and ED.  At this time it is premature for us to speculate about the outcome of binary statistics while we wait for the I-band data.  While this is an observational paper, written to promote the usefulness of the ASAS catalog of variable stars, we are tempted to speculate about possible interpretation of Fig. 6, which was so surprising to us. We are not consistent with the previous paragraph, but the temptation is hard to resist.  The following is a speculative hypothesis.  There has been a gradual emergence of the notion that contact systems have companions (Ruci\\'nski and Kaluzny 1982, Chambliss 1992, Hendry and Mochnacki 1998, also Tokovinin 2004). Recently. Pribulla and Ruci\\'nski (2005) found that up to 50\\% of W UMa binaries have companions. This opens up a possibility that Kozai (1962) cycle operates in some such triples, as suggested by Kiseleva, Eggleton, and Mikkola (1998), and by Eggleton and Kiseleva-Eggleton (2001). The mechanical three-body orbital evolution, with a large range of inclinations of the two orbits, inner and outer, may induce large variations in the eccentricity of the inner binary.  In time the eccentricity may become large enough to make a contact system out of the inner binary.  In this scenario the inner binary has the initial period that is relatively long, and in most cases no eclipses would be detectable. During a Kozai cycle the inner binary occasionally has a chance to reach a physical contact: it may either merge forming a single star, or it may become a W UMa star.  This process is somewhat similar to the model of formation of close binaries in globular clusters (Pooley et al. 2003, and references therein).  The surprisingly small number of short period detached eclipsing binaries as seen in Fig. 6 may indicate that the Kozai cycle is not just a curiosity, but it may be an important channel for forming W UMa stars. This possibility cannot be rejected without careful analysis.  After all W UMa stars are rare, with the local space density of just 0.2\\% of the main sequence stars (Rucinski 2002, and references therein)."
    },
    "0601/astro-ph0601485.txt": {
        "abstract": "We present a {\\it Spitzer} mid-infrared survey of 42 Fanaroff-Riley class II radio galaxies and quasars from the 3CRR catalog at redshift $z<1$. All of the quasars and $45\\pm 12\\%$ of the narrow-line radio galaxies have a mid-IR luminosity of $\\nu L_\\nu(15 \\mu\\mathrm{m}) > 8 \\times 10^{43}$ erg s$^{-1}$, indicating strong thermal emission from hot dust in the active galactic nucleus. Our results demonstrate the power of {\\it Spitzer} to unveil dust-obscured quasars. The ratio of {\\it mid-IR luminous} narrow-line radio galaxies to quasars indicates a mean dust covering fraction of $0.56 \\pm 0.15 $, assuming relatively isotropic emission. We analyze {\\it Spitzer} spectra of the 14 mid-IR luminous narrow-line radio galaxies thought to host hidden quasar nuclei. Dust temperatures of 210-660 K are estimated from single-temperature blackbody fits to the low and high-frequency ends of the mid-IR bump. Most of the mid-IR luminous radio galaxies have a 9.7 $\\mu$m silicate absorption trough with optical depth $<0.2$, attributed to dust in a molecular torus. Forbidden emission lines from high-ionization oxygen, neon, and sulfur indicate a source of far-UV photons in the hidden nucleus. However, we find that the other $55\\pm 13\\%$ of narrow-line FR II radio galaxies are weak at 15 $\\mu$m, contrary to single-population unification schemes. Most of these galaxies are also weak at 30 $\\mu$m.  Mid-IR weak radio galaxies may constitute a separate population of nonthermal, jet-dominated sources with low accretion power. ",
        "introduction": "The nature of the energy source in active galactic nuclei (AGNs) is a fundamental problem. The basic model attributes the large luminosity of these systems to gravitational energy release in an accretion disk around a supermassive black hole. A jet may be driven by magnetic fields threading the disk \\citep{bp82}. The black hole spin energy may also be tapped and converted into electromagnetic Poynting flux and particles in a relativistic jet \\citep{bz77,pc90,m99,d04}.  Extragalactic radio sources are categorized by their morphology as either of two types \\citep{fr74}. Fanaroff-Riley (FR) type I sources are edge-darkened, while FR IIs are edge-brightened. The different morphology of FR Is indicates that they are not related to FR IIs by orientation. FR Is also have lower radio luminosities than FR IIs for a given host galaxy luminosity \\citep{ol94,b96}  and most have low-ionization nuclear emission region (LINER) spectra \\citep{hl79}. However, not all FR Is can be characterized by low accretion power  \\citep{wa04,cr04}. The present paper focuses on FR IIs, which contain powerful jets with bright terminal hot spots and lobes. Furthermore, we count broad-line radio galaxies (BLRGs) as low-luminosity quasars.  Quasars and narrow-line radio galaxies (NLRGs) may be unified by orientation-dependent obscuration. Radio galaxies are thought to host quasar nuclei that are obscured by circumnuclear dusty tori aligned with the radio jets \\citep{ant84} . Unification of radio galaxies and quasars can therefore explain the lack of quasars viewed at large angles to the radio axis \\citep{b89}. The percentage of high-redshift radio galaxies (60\\% of the 3CRR FR II sample at $z>0.5$) would then indicate a torus covering fraction of $\\sim 0.6$.  However, there appears to be a discrepancy between the redshift distributions of quasars and radio galaxies at $z<0.5$, with a factor of $\\sim 4$  more narrow-line radio galaxies than quasars \\citep{s93}. Furthermore, the median projected linear size of these 'excess' radio galaxies is smaller than expected for quasars seen in the sky plane \\citep{s93,w05}. The unification hypothesis may be modified to include a second population of lower luminosity, low-excitation FR II radio galaxies \\citep{wj97,grw04}. Alternatively, it has been argued that the torus covering fraction may increase with decreasing radio luminosity \\citep{l91}.  The unification hypothesis has been qualitatively confirmed by spectropolarimetry of radio galaxies, many of which have been shown to have highly polarized broad emission lines and blue continuum, scattered from material which has a direct view of the active galactic nucleus \\citep{cdb97, cot99}. Of particular note are the original discovery of highly polarized broad H$\\alpha$ from the hidden quasar nucleus in 3C 234 \\citep{ant84}, and the discovery of highly polarized broad H$\\alpha$ in the spectrum of the powerful radio galaxy Cygnus A \\citep{ocm97}. However, this method of detecting hidden quasars relies on an appropriately placed scattering region to view the otherwise hidden nucleus. Such a region is not guaranteed to exist for all radio galaxies, and thus spectropolarimetry can easily yield false negatives. Polarimetry is also ineffective at determining the luminosity of the hidden nucleus, since the scattering efficiency is usually unknown.  Another way to search for hidden quasar nuclei is to observe radio galaxies in the mid-IR. If the unification hypothesis is correct, the dusty torus should serve as a crude calorimeter of the central engine \\citep{mhm01,sfk04,hmb04,wa04}. Optical, UV, and X-ray  photons from the quasar nucleus are absorbed by dust in the torus and the energy is re-emitted in the thermal infrared. This explains why blue, UV color-selected quasars emit 10-50\\% of their luminosity in the IR \\citep{spn89,hmc00}. There appears to be no connection between the bulk of this IR emission and nonthermal radio emission, except in core-dominated radio sources such as blazars. Observations of matched 3CR quasars and radio galaxies by ISO indicate similar IR luminosities, consistent with the unification picture \\citep{mhm01,hmb04}. However, differences in 24 $\\mu$m/ 70 $\\mu$m color may indicate that mid-IR emission from the torus is anisotropic by a factor of $\\le 3$ \\citep{srh05}.  We present {\\it Spitzer} observations of a sample of 42 FR II radio galaxies and quasars selected from the 3CRR survey. The goals are to search for mid-IR emission from hidden quasar nuclei and test the ubiquity of the unification hypothesis. The {\\it Spitzer} Infrared Spectrograph (IRS) combines the advantages of unprecedented sensitivity from 5-36.5 $\\mu$m to measure the mid-IR continuum and spectral resolution to measure high ionization emission lines powered by hidden AGNs. In the current paper, we present evidence for hidden quasar nuclei based on mid-IR photometry extracted from the IRS spectra. We examine in detail the spectra of the subset of 14 mid-IR luminous radio galaxies which appear to contain hidden quasar nuclei. Spectra of the quasars and mid-IR weak radio galaxies and a statistical study of the complete sample will be presented in separate papers. ",
        "conclusions": "(1.) We report on a large {\\it Spitzer} spectroscopic survey of 3CRR FR II radio sources with $S_{178}>16.4$ Jy at $z<1$. We find strong mid-IR emission from $45 \\pm 12\\%$ (14/31) of NLRGs, which have luminosities comparable to matched BLRGs and quasars. Other indicators including high-ionization mid-IR lines and highly polarized broad emission lines confirm that some of these sources contain hidden quasar or BLRG nuclei. This demonstrates the power of {\\it Spitzer} IRS for unveiling hidden quasars and estimating their luminosities.  (2.) We present {\\it Spitzer} spectra of the 14 mid-IR luminous radio galaxies. In most cases, the mid-IR continuum bump from 3-30 $\\mu$m can be produced by a distribution of hot dust with temperatures in the range 210-$660$ K. These high temperatures are most likely maintained by hidden AGNs. The silicate absorption trough at 9.7 $\\mu$m has an apparent optical depth of $\\tau=0-0.2$ in most cases, consistent with dust temperatures decreasing outward from the center of a dusty torus. Two sources, 3C 55 and 433, have deeper silicate troughs which may be produced by additional cool dust in the host galaxy.  (3.) However, not all FR II radio galaxies emit strongly in the mid-IR. Contrary to single-population unification schemes, the majority of narrow-line radio galaxies in our sample (17/31 or $55 \\pm 13\\%$) have weak or undetected mid-IR emission compared to matched quasars and BLRGs, with $\\nu L_\\nu(15$ $\\mu\\mathrm{m}) < 8 \\times 10^{43}$ erg s$^{-1}$. For a few sources, this may possibly be the result of anisotropic torus emission viewed through a large column density of dust. However, it is likely that most of the weakest sources do not contain a powerful accretion disk. These may be truly nonthermal, jet-dominated AGNs, where the jet is powered by a radiatively inefficient accretion flow or black hole spin-energy rather than energy extracted from an accretion disk."
    },
    "0601/astro-ph0601356_arXiv.txt": {
        "abstract": "Speculations that encounters with interstellar clouds modify the terrestrial climate have appeared in the scientific literature for over 85 years.  The articles in this volume seek to give substance to these speculations by examining the exact mechanisms that link the pressure and composition of the interstellar medium surrounding the Sun to the physical properties of the inner heliosphere at the Earth. ",
        "introduction": "If the solar galactic environment is to have a discernible effect on events on the surface of the Earth, it must be through a subtle and indirect influence on the terrestrial climate.  The scientific and philosophical literature of the 18th, 19th and 20th centuries all include discussions of possible cosmic influences on the terrestrial climate, including the effect of cometary impacts on Earth (\\nolinebreak \\cite{Halley:1694}), and the diminished solar radiation from sunspots, which Herschel attributed to ``holes'' in the luminous fluid on the surface of the Sun\\footnote{In this same paper Herschel commented that ``Whatever fanciful poets might say, in making the sun the abode of blessed spirits, or angry moralists devise, in pointing it out as a fit place for the punishment of the wicked, it does not appear that they had any other foundation for their assertions than mere opinion and vague surmise; but now I think myself authorized, \\emph{upon astronomical principles,} to propose the sun as an inhabitable world, and am persuaded that the foregoing observations, with the conclusions I have drawn from them, are fully sufficient to answer every objection that may be made against it.  ''  These comments show that valuable data are not always interpreted correctly.} (\\nolinebreak \\cite{Herschel:1795}). The discovery of interstellar material in the 20th century led to speculations that encounters with dense clouds initiated the ice ages (\\nolinebreak \\cite{Shapley:1921}), and many papers appeared that explored the implications of such encounters, including the influence of interstellar material (ISM) on the interplanetary medium and planetary atmospheres (e.g. \\cite{Fahr:1968,BegelmanRees:1976,McKayThomas:1978,Thomas:1978,McCrea:1975,TalbotNewman:1977,Willis:1978,ButlerNewmanTalbot:1978}). The ISM-modulated heliosphere was also believed to affect climate stability and astrospheres (e. g. \\cite{Frisch:1993a,Frisch:1997,ZankFrisch:1999}).  Recent advances in our understanding of the solar wind and heliosphere (e. g. \\cite{WangRichardson:2005,Fahr:2004}) justify a new look at this age-old issue.  This book addresses the underlying question: \\begin{verse} \\emph{How does the heliospheric interaction} \\emph{with the interstellar medium affect the heliosphere, interplanetary medium, and Earth?} \\end{verse}  The heliosphere is the cavity in the interstellar medium created by the dynamic ram pressure of the radially expanding solar wind, a halo of plasma around the Sun and planets, dancing like a candle in the wind and regulating the flux of cosmic rays and interstellar material at the Earth.  Neutral interstellar gas and large interstellar dust grains penetrate the heliosphere, but the solar wind acts as a buffer between the Earth and most other interstellar material and low energy galactic cosmic rays (GCR).  Together the solar wind and interstellar medium determine the properties of the heliosphere.  In the present epoch the densities of the solar wind and interstellar neutrals are approximately equal outside of the Jupiter orbit.  Solar activity levels drive the heliosphere from within, and the physical properties of the surrounding interstellar cloud constrain the heliosphere from without, so that the boundary conditions of the heliosphere are set by interstellar material.  Figure \\ref{fig:1} shows the Sun and heliosphere in the setting of the Milky Way Galaxy.   The answer to the question posed above lies in an interdisciplinary study of the coupling between the interstellar medium and the solar wind, and the effects that ISM variations have on the 1 AU environment of the Earth through this coupling.  The articles in this book explore different viewpoints, including \\emph{gedanken} experiments, as well as data-rich summaries of variations in the solar environment and paleoclimate data on cosmic ray flux variations at Earth.  The book begins with the development of theoretical models of the heliosphere that demonstrate the sensitivity of the heliosphere to the variations in boundary conditions caused by the passage of the Sun through interstellar clouds (\\cite{Zanketal:2006jos,PogorelovZank:2006jos}).  A series of \\emph{gedanken} experiments then yield the response of planetary magnetospheres to encounters with denser ISM (\\cite{Parker:2006jos}).  Variations in the galactic environment of the Sun, caused by the motions of the Sun and clouds through the Galaxy, are shown to occur for both long and short timescales (\\cite{Shaviv:2006jos,FrischSlavin:2006jos}).  The heliosphere acts as a buffer between the Earth and interstellar medium, so that dust and particle populations inside of the heliosphere, which have an interstellar origin, vary as the Sun traverses interstellar clouds.  These buffering mechanisms determine the interplanetary medium\\footnote{The buffering processes convert interstellar neutrals into low energy ions, which are convected outwards with the solar wind and accelerated to low cosmic ray energies that have an anomalous composition, including abundant elements with FIP$>$13.6 eV. The high energy galactic cosmic ray population incident on the heliosphere is also modulated.}.  The properties of these buffering interactions are evaluated for heliosphere models that have been developed using boundary conditions appropriate for when the Sun traverses different types of interstellar clouds (\\cite{Landgraf:2006jos,Moebiusetal:2006jos,FlorinskiZank:2006jos,Fahretal:2006jos}).  The consequences of Sun-cloud encounters are then discussed in terms of the accretion of ISM onto the terrestrial atmosphere for dense cloud encounters, and the possibly extreme variations expected for cosmic ray modulation when interstellar densities vary substantially (\\cite{Fahretal:2006jos,YeghikyanFahr:2006jos}). Radioisotope records on Earth extending backwards in time for over $\\sim$0.5 Myrs, together with paleoclimate data, suggest that cosmic ray fluxes are related to climate.  The galactic environment of the Sun must have left an imprint on the geological record through variations in the concentrations of radioactive isotopes (\\cite{KirkbyCarslaw:2006jos}).  The selection of topics in this book is based partly on scientific areas that have already been discussed in the literature.  The authors who were invited to contribute chapters have previously studied the heliosphere or terrestrial response to variable ISM conditions or cosmic rays.    \\begin{figure} \\centering \\parbox[t]{3.10in}{\\centering \\centerline{ \\psfig{figure=frisch_composite.eps,width=5.5in} }} \\caption[The Solar Journey.]{\\label{fig:1} The solar location and vector motion are identified for the kiloparsec scale sizes of the Milky Way Galaxy (large image), and for the $\\sim$500 parsec scale size of the Local Bubble (medium sized image, inset in upper left hand corner).  A schematic drawing of the heliosphere (small image, inset in lower right hand inset, based on Linde,1998\\nocite{Linde:1998}) shows the upwind velocity of the interstellar wind (``ISM'') as observed in the rest frame of the Sun. Coincidently, this direction, which determines the heliosphere nose, is close to the galactic center direction. The orientation of the plane in the small inset differs from the planes of the large and medium figures, since the ecliptic plane is tilted by 60$^\\circ$ with respect to the galactic plane.  The Sun is 8 kpc from the center of the Milky Way Galaxy, and the solar neighborhood moves towards the direction \\glong=90\\deeg\\ at a velocity of 225 \\kms.  The spiral arm positions are drawn from Vallee (2005), except for the Orion spur. \\nocite{ValleeSA:2005} The Local Bubble configuration is based on measurements of starlight reddening by interstellar dust (Chapter 6).  The lowest level of shading corresponds to color excess values \\ebv=0.051 mag, or column densities log $N$(H) (\\cmtwo) =20.40 dex. The dotted region shows the widespread ionized gas associated with the Gum Nebula.  The heliosphere cartoon shows interstellar protons deflected in the plasma flow in the outer heliosheath regions, compared to the interstellar neutrals that penetrate the heliopause.} \\end{figure}  Figure \\ref{fig:1} shows the heliosphere in our setting of the Milky Way Galaxy.  A postscript at the end of this chapter lists basic useful information.  I introduce the term ``paleoheliosphere'' to represent the heliosphere in the past, when the boundary conditions set by the local interstellar material (LISM) may have differed substantially from the boundary conditions for the present-day heliosphere.  The ``paleolism'' is the local ISM that once surrounded the heliosphere.  \\section[Addressing the Query: The Heliosphere for Different Interstellar Environments]{Addressing the Query:  The Heliosphere and Particle Populations for Different Interstellar Environments}   The solar wind drives the heliosphere from the inside, with the properties of the solar wind varying with ecliptic latitude and the phase of the 11-year solar activity cycle.  The global heliosphere is the volume of space occupied by the supersonic and subsonic solar wind.  Interstellar material forms the boundary conditions of the heliosphere, and the windward side of the heliosphere, or the ``upwind direction'', is defined by the interstellar velocity vector with respect to the Sun.  The leeward side of the heliosphere is the ``downwind direction''.  Figure \\ref{fig:1} shows a cartoon of the present-day heliosphere, with labels for the major landmarks such as the termination shock, heliopause, and bow shock.   In the present-day heliosphere, the transition from solar wind to interstellar plasma occurs at a contact discontinuity known as the ``heliopause'', which is formed where the total solar wind and interstellar pressures equilibrate (\\nolinebreak \\cite{Holzer:1989}). For a non-zero interstellar cloud velocity in the solar rest frame, the solar wind turns around at the heliopause and flows around the flanks of the heliosphere and into the downwind heliotail.  Before reaching the heliopause, the supersonic solar wind slows to subsonic velocities at the ``termination shock'', where kinetic energy is converted to thermal energy.  The subsonic solar wind region between the termination shock and heliopause is called the inner ``heliosheath''.  The outer heliosheath lies just beyond the heliopause, where the pristine ISM is distorted by the ram pressure of the heliosphere.  A bow shock, where the interstellar gas becomes subsonic, is expected to form ahead of the present-day heliosphere in the observed upwind direction of the ISM flow through the solar system.  Large interstellar dust grains and interstellar atoms that remain neutral inside of the orbit of Earth, such as He, are gravitationally focused in the downwind direction.  This ``focusing cone'' is traversed by the Earth every year in early December, and extends many AU from the Sun in the leeward direction (e.g. \\cite{Landgraf:2000,Moebiusetal:2004,Frisch:2000amsci}).  The heliotail itself extends $>10^3$ AU from the Sun in the downwind direction, forming a cosmic wake for the solar system.  Of significance when considering the interaction of the heliosphere with an interstellar cloud is that neutral particles enter the heliosphere relatively unimpeded, after which they are ionized and convected outwards with the solar wind.  Ions and small charged dust grains are magnetically deflected in the heliosheath around the flanks of the heliosphere (see Figure \\ref{fig:1}).  Space and astronomical data now confirm the basic milestones of the outer heliosphere.  Voyager 1 crossed the termination shock at 94 AU on 16 December, 2004 (UT), and observed the signature of the termination shock on low-energy particle populations, the solar wind magnetic field, low-energy electrons and protons, and Langmuir radio emission (\\nolinebreak \\cite{Stoneetal:2005,Burlagaetal:2005,Gurnett:2005,Deckeretal:2005}). The present-day termination shock appears to be weak, with a solar wind velocity jump ratio (the ratio of upstream to downstream values) of $\\sim$2.6 and a magnetic field compression ratio of $\\sim$3.  The magnetic wall that is predicted for the heliosphere (\\nolinebreak \\cite{Linde:1998,Ratkiewicz:1998}, Chapter 3 by Pogorelov and Zank) appears to have been detected through observations of magnetically aligned dust grains (\\nolinebreak \\cite{Frisch:2005L}), and the offset between upwind directions of interstellar \\HI\\ and \\HeI\\ (\\nolinebreak \\cite{Lallementetal:2005}). The compressed and heated \\HI\\ in the hydrogen wall region of the outer heliosheath has now been detected around a number of stars (\\nolinebreak \\cite{Woodetal:2005}).  The present-day solar wind is the baseline for evaluating the heliosphere response to ISM variations in the following articles, so a short review of the solar wind is first presented.  The remaining part of $\\S$1.2 introduces the topics in the following articles in terms of the underlying query of the book.  \\subsection{The Present Day Solar Wind }  The solar wind originates in the million degree solar corona that expands radially outwards, with a density $\\sim 1/R^2_\\mathrm{ S }$ where $R_\\mathrm{ S }$ is the distance to the Sun, and contains both features that corotate with the Sun, and transient structures (\\nolinebreak e.~g.~\\cite{Gosling:1996}).  The properties of the solar wind vary with the phase of the solar magnetic activity cycle and with ecliptic latitude.  The best historical indicator of solar magnetic activity levels is the number of sunspots, first detected by Galileo in 1610, which are magnetic storms in the convective zone of the Sun.  Sunspot numbers indicate that the magnetic activity levels fluctuate with a $\\sim$11 year cycle, or the ``solar cycle'', and solar maximum/minimum corresponds to the maximum/minimum of sunspot numbers.  The magnetic polarity of the Sun varies with a $\\sim$22 year cycle.  During solar maximum, a low-speed wind, with velocity $\\sim$300--600 \\kms\\ and density $\\sim 6-10$ particles \\cc\\ at 1 AU, extends over most of the solar disk.  Open magnetic field lines\\footnote{Open magnetic field lines are formed in coronal holes that reconnect in the outer heliosphere and contain low density and very high speed, $\\sim$700 \\kms, solar wind.} are limited to solar pole regions.  A neutral current sheet $\\sim$0.4 AU thick forms between the solar wind containing negative magnetic polarity fields and the solar wind that contains positive magnetic polarity fields. The neutral current sheet reaches its largest inclination ($\\ge 70^\\circ$) during solar maximum. During the conditions of solar minimum, a high speed wind with velocity $\\sim$600--800 \\kms\\ and density $\\sim$5 \\cc\\ is accelerated in the open magnetic flux lines in coronal holes.  During mininum, the high speed wind and open field lines extend from the polar regions down to latitudes of $\\leq 40$\\deeg (\\nolinebreak \\cite{Smithetal:2003,Richardson:1995}).  The higher solar wind momentum flux associated with solar minimum conditions produces an upwind termination shock that is $\\sim$5--40 AU more distant in the upwind direction than during solar maximum conditions (e.g. \\cite{SchererFahr:2003,ZankMueller:2003,Whang:2004}).  The alignment and strength of the solar magnetic multipoles depend on the phase of the solar cycle (\\cite{Bravoetal:1998}).  During solar minimum conditions, the magnetic field is dominated by the dipole and hexapole moments, and the dipole moment is generally aligned with the solar rotation axis. Sunspots migrate from high to low heliographic latitudes. The magnetic poles follow the coronal holes to the solar equator as solar activity increases.  During the solar maximum period, the galactic cosmic rays undergo their maximum modulation, the dipole component of the magnetic field is minimized, the quadrapole moment dominates, and the polarity of the solar magnetic field reverses (\\cite{LockwoodWebber:2005}, Figure \\ref{fig:2}).  Over historic times, the cosmic ray modulation by the heliosphere correlates better with the open magnetic flux line coverage than with sunspot numbers (\\cite{McCrackenetal:2004}).   Variable cosmic ray modulation produced by a variable heliosphere may be a primary factor in both solar and ISM forcing of the terrestrial climate.  The heliosphere modulation of cosmic rays is well established.  John Simpson, to whom this book is dedicated, initiated a program 5 solar cycles ago in 1951 to monitor cosmic ray fluxes on Earth using high-altitude neutron detectors (\\cite{Simpson:2001}). The results show a pronounced anticorrelation between cosmic ray flux levels and solar sunspot numbers, which trace the 11-year Schwabe magnetic activity cycle, and which also show that the polarity of the solar magnetic field affects cosmic ray modulation (see Figure \\ref{fig:2}). The articles in this book show convincingly that the ISM also modulates the heliosphere, and the effect of the solar wind on the heliosphere must be differentiated from the influence of interstellar matter.   Variations in solar activity levels are also seen over $\\sim 100-200$ year timescales, such as the absence of sunspots during the Maunder Minimum in the 17th century.  Modern climate records show that the Maunder Minimum corresponded to extremely cold weather, and radioisotope records show that the flux of cosmic rays was unusually high at this time (see Kirkby and Carslaw, Chapter 12). Similar effects will occur from the modulation of galactic cosmic rays by the passage of the Sun through an interstellar cloud.  These temporal and latitudinal variations in the solar wind momentum flux produce an asymmetric heliosphere, which varies with time.  Any possible historical signature of the ISM on the heliosphere must first be distinguished from variations driven by the solar wind itself.  \\subsection{Present Day Heliosphere and Sensitivity to ISM}  The ISM forms the boundary conditions of the heliosphere, so that encounters with interstellar clouds will affect the global heliosphere, the interplanetary medium, and the inner heliosphere region where the Earth is located.  Today an interstellar wind passes through the solar system at --26.3 \\kms\\ (\\cite{Witte:2004}).  An entering parcel of ISM takes about 20 years to reach the inner heliosphere, so that ISM near the Earth is constantly replenished with new inflowing material.  This warm gas is low density and partially ionized, with temperature $T \\sim$6,300 K, and densities of neutral and ionized matter of \\nHI$\\sim 0.2$ \\cc, and \\nHII$\\sim 0.1$ \\cc.   An elementary perspective of the response of the heliosphere to interstellar pressures is given by an analytical expression for the heliopause distance based on the locus of positions where the solar wind ram pressure, $P_\\mathrm{SW}$, and the total interstellar pressure equilibrate (\\cite{Holzer:1989}).  The solar wind density $\\rho$ falls off as $\\sim 1/R^2$, where $R$ is the distance to the Sun, while the velocity $v$ is relatively constant.  At 1 AU the solar wind ram pressure is $P_\\mathrm{SW,1AU} \\sim \\rho~v^2 $ so the heliosphere distance, $R_\\mathrm{HP} $, is given by:   \\begin{eqnarray*} P_\\mathrm{SW,1AU}/R_\\mathrm{HP}^2  \\sim  P_\\mathrm{B} +P_\\mathrm{Ions, thermal} + P_\\mathrm{Ions, ram} + P_\\mathrm{Dust} + P_\\mathrm{CR} \\end{eqnarray*} The interstellar pressure terms include the magnetic pressure $P_\\mathrm{B}$, the thermal, $P_\\mathrm{Ions, thermal}$, and the ram, $P_\\mathrm{Ions, ram}$, pressures of the charged gas, and the pressures of dust grains, $P_\\mathrm{Dust}$, and cosmic rays, $P_\\mathrm{CR}$, which are excluded by heliosphere magnetic fields and plasma.  Some interstellar neutrals convert to ions through charge exchange with compressed interstellar proton gas in heliosheath regions, adding to the confining pressure.  An important response characteristic is that, for many clouds, the encounter will be ram-pressure dominated, where $P_\\mathrm{ram} \\sim m v^2$ for interstellar cloud mass density $m$ and relative Sun-cloud velocity $v$, so that variations in the cloud velocity perturb the heliosphere even if the thermal pressures remain constant.  The multifluid, magnetohydrodynamic (MHD), hydrodynamic and hybrid approaches used in the following chapters provide much more substantial models for the heliosphere, and include the coupling between neutrals and plasma, and field-particle interactions.  These sophisticated models predict variations in the global heliosphere in the face of changing interstellar boundary conditions, and for a range of different cloud types.  Although impossible to model a solar encounter with every type of interstellar cloud, the following articles include discussions of many of the extremes of the interstellar parameter space, including low density gas with a range of velocities, very tenuous plasma, high velocity clouds, dense ISM, and magnetized material for a range of field orientations and strengths.  The discussions in these chapters extrapolate from our best theoretical understanding of the heliosphere boundary conditions today to values that differ, in some cases dramatically, from the boundary conditions that prevailed at the beginning of the third millennium in the Gregorian calendar.  The Sun has been, and will be, subjected to many different physical environments over its lifetime. Theoretical heliosphere models yield the properties of the solar wind-ISM interaction for these different environments, which in turn determine the nature and properties of interstellar populations inside of the heliosphere for a range of galactic environments. These models form the foundation for understanding the significance of our galactic environment for the Earth.  The interstellar parameter space is explored by Zank et al. (Chapter 2), where 28 sets of boundary conditions are evaluated with computationally efficient multifluid models.  Moebius et al. (Chapter 8), Fahr et al. (Chapter 9), Florinski and Zank (Chapter 10), and Yeghikyan and Fahr (Chapter 11) also develop heliosphere models for a range of interstellar conditions.  Together these models evaluate the heliosphere response to interstellar density, temperature, and velocity variations of factors of $\\sim 10^{9}$, $\\sim 10^{5}$, and $\\sim 10^{2}$, respectively.  The interstellar magnetic field introduces an asymmetric pressure on the heliosphere, affecting the heliosphere current sheet and cosmic ray modulation. Pogorelov and Zank (Chapter 3) use MHD models to probe the heliosphere response to the interstellar magnetic field, including charge exchange between the neutrals and solar wind.  The resulting asymmetry provides a test of the magnetic field direction, and shows strong differences between cases where the interstellar flow is parallel, instead of perpendicular, to the interstellar magnetic field direction. Since the random component of the interstellar magnetic field is stronger, on the average, than the ordered component, particularly in spiral arm regions where active star formation occurs, a range of interstellar magnetic field strengths and orientations are expected over the solar lifetime (Shaviv, Chapter 5, and Frisch and Slavin, Chapter 6).   \\subsection{Planetary Magnetospheres}  The Earth's magnetosphere acts as a buffer between the solar wind and atmosphere, and as such is an ingredient in understanding the effect of our galactic environment on the Earth.  The decreasing solar wind density in the outer heliosphere results in an interplanetary medium around outer planets that is more sensitive to ISM variations than for inner planets, with implications for the magnetospheres of Jupiter, Neptune, and Uranus.  Most topics in this book are already considered in the scientific literature, but questions about magnetosphere variations from an ISM-modulated heliosphere have received scant attention.  In a quintessential \\emph{gedanken} experiment, Parker explores the interaction between magnetospheres and the solar wind for variations in the interstellar density, and for inner versus outer planets (Chapter 4).  \\subsection{Short and Long Term Variations in the Galactic Environment}  There is every reason to expect that the galactic environment of the Sun varies over geological timescales.  The Sun moves through space at a velocity of 13--20 \\kms, and interstellar clouds have velocities ranging up to hundreds of \\kms.  The Arecibo Millennium survey showed that $\\sim$25\\% of the mass contained in interstellar \\HI, including both warm and cold ISM, is in clouds traveling with velocities $\\ge$10 \\kms\\ through the local standard of rest (\\nolinebreak \\cite{HTI}). Thus Sun-cloud encounters with relative velocities exceeding 25 \\kms\\ are quite likely, and for a typical cloud length of $\\sim$1 pc the cloud transit time would be $\\sim$40,000 years.  The many types of ISM traversed by the Sun during the past several million years have affected the heliosphere, the inner solar system, and the flux of anomalous and galactic cosmic rays at Earth (\\cite{FrischYork:1986,Frisch:1997,Frisch:1998}).  For the past $\\sim$3 Myrs the Sun has been in a nearly empty region of space, the ``Local Bubble'', with very low densities of $<$10$^{-26}$ gr \\cc.  Within the past 44,000--140,000 years the Sun entered a flow of tenuous, partly neutral ISM, nick-named the ``Local Fluff'', with density $\\sim$60 times higher (Chapter 6).  This transition was accompanied by the appearance of interstellar dust and neutrals in the heliosphere, along with the pickup ion and anomalous cosmic ray populations. Galactic cosmic ray modulation was affected, providing a possible link between our galactic environment and climate.  Intriguingly, the averaged cosmic ray flux at Earth, as traced by \\Beten\\ records, was lower in the past $\\sim$135 kyrs than for earlier times (Chapter 12). Was the decrease in the galactic cosmic ray flux $\\sim$135 kyrs ago caused by an increase in modulation as the Sun entered the Local Fluff?  The galactic environment of the Sun also varies quite dramatically over long time scales, as discussed by Shaviv (Chapter 5).  Over its 4.5 billion year lifetime, the Sun traverses spiral arm and interarm regions, with atomic densities varying from less than $10^{-26.1}$ g \\cc\\ to over $10^{-20.1}$ g \\cc, and temperatures ranging over 7 orders of magnitude, 10--10$^7$ K.  The Sun is now in low density space between the Perseus and Sagittarius spiral arms, and on the inner edge of what is known as the Orion Spur on the Local Arm.  The Local Arm is not shown in Figure \\ref{fig:1}, as is consistent with the usual Galaxy depictions.  The Local Arm does not appear to be a grand design spiral shock (Bochkarev, 1984\\nocite{Bochkarev:1984}).  The Sun has a systematic motion of 13--20 km s$^{-1}$ with respect to the nearest stars, corresponding to $\\sim$3--4 AU per year.  The Local Interstellar Cloud (LIC) now surrounding the Sun traverses the heliosphere at $\\sim$5.5 AU per year.  The Sun oscillates vertically through the galactic plane once every $\\sim$34 Myrs, and orbits the center of the Milky Way Galaxy once per $\\sim$220 Myrs.  Shaviv evaluates variations in the galactic environment of the Sun over long timescales.  This bold discussion compares various geologic records of cosmic ray flux variations, based on radioisotope data that sample timescales of $\\sim 10^8$ years, with models of the Milky Way Galaxy spiral arm pattern to reconstruct the timing of the Sun's passage through spiral arms.  The chapter concludes that star formation in spiral arms leaves a signature on the radioisotope records of the solar system.  Frisch and Slavin (Chapter 6) reconstruct short-term variations of the galactic environment of the Sun using observations of interstellar matter towards nearby stars and inside of the solar system.  Radiative transfer models of the LIC show that ionization varies across this low density cloud, so that the heliosphere boundary conditions vary from radiative transfer considerations alone as the Sun traverses the LIC. Cloud transitions are predicted during the past $\\sim$3 Myrs, including the departure of the Sun from the Local Bubble interior 44,000--140,000 years ago, and entry into the surrounding cloud 1000--40,000 years ago.  \\subsection{Interstellar Dust}  The particle populations formed by the interactions between the solar wind and interstellar dust, gas, and cosmic rays are emissaries between the cosmos and inner heliosphere, varying as the Sun moves through clouds.  About $\\sim$1\\% of the mass of the cloud surrounding the Sun is contained in interstellar dust grains.  The largest of these charged grains, mass $> 10^{-13} $ g, have large magnetic Larmor radii of $>$500 AU at the heliopause for an interstellar field of $\\sim$1.5 $\\mu$G, and flow into the solar system.  The Earth passes through the gravitational focusing cone formed by these grains early each December.  The smallest charged grains, mass$< 10^{-14.5} $ g and radii$< 0.01 ~\\mu$m, have Larmor radii of $\\sim$20 AU, depending on the magnetic field strength and radiation field, and are deflected around the heliosheath (\\cite{Frischetal:1999}).  Interstellar dust grains are measured in the inner heliosphere within $\\sim$5 AU of the Sun, and over the solar poles, by satellites such as Ulysses, Galileo and Cassini.  Landgraf (Chapter 7) reviews the properties of the interaction between interstellar dust and the solar wind, and speculates on the changes that might be expected from an encounter with a dense interstellar cloud.  Should it some day be possible to compare the ratio of large to small interstellar dust grains on the surfaces of the inner versus outer planets, it would become possible to disentangle cloud encounters from solar activity effects.  At the very large end of the dust population mass spectrum we find interstellar micrometeorites, with masses $\\sim$3 $\\times$ 10$^{-7}$ g, open orbits, and inflow velocities greater than the 42 \\kms\\ escape velocity from the solar system at 1 AU.  These interstellar objects, detected by radar as they impact the atmosphere, evidently originate in circumstellar disks such as that around $\\beta$ Pictoris, and in the interior of the Local Bubble (\\cite{Baggaley:2000,Meiseletal:2002}). These objects do not collisionally couple to the interstellar gas (\\cite{GruenLandgraf:2000}), and should not vary with the type of ISM surrounding the Sun.  \\subsection{Particle Populations in the Inner and Outer Heliosphere}  Presently, low energy interstellar neutrals, high energy galactic cosmic rays, and interstellar dust all enter the heliosphere. The characteristics of each of these populations and their secondary products are modified as the Sun transits the ISM, or the cloud ionization changes. The first ionization potential (FIP) of \\HI\\ is 13.6 eV.  Neutral interstellar atoms with FIP$<$13.6 eV are ionized in nearly all interstellar clouds because the main source of interstellar opacity is \\HI.  Interstellar ions are deflected around the heliosheath, so the result is that only interstellar atoms with FIP$>$13.6 eV enter the heliosphere where they are then destroyed, primarily by charge exchange with solar wind ions.  The density of interstellar neutrals in the inner heliosphere depends on the density and ionization of the surrounding cloud, the ionization (or ``filtration'') of those neutrals by the heliosheath, and the subsequent interactions with the solar wind inside of the heliosphere. Secondary products produced by solar wind interactions with interstellar neutrals inside of the heliosphere include pickup ions \\footnote{The pickup ions are interstellar neutrals formed by charge exchange with the solar wind.  Energetic neutral atoms are formed by energetic ions that capture an electron from a low energy neutral by charge exchange.  The gravitational focusing cone contains heavy elements (mainly He) that are predominantly ionized inside of 1 AU and therefore gravitationally focused downwind of the Sun (Chapter 8). Large interstellar dust grains are also gravitationally focused (Chapter 7).  The anomalous cosmic ray population is formed from pickup ions accelerated to low cosmic ray energies, $<$\\nolinebreak 1 GeV, in the solar wind and at the termination shock, and then subjected to the same modulation and propagation processes as galactic cosmic rays (\\cite{Jokipii:2004}).}, energetic neutral atoms, the gravitational focusing cone formed by helium (also seen in dust), and the anomalous cosmic ray population with energies $<$\\nolinebreak 1 GeV.  Interstellar neutrals inside of the heliosphere, and the heliosphere itself, form a coupled system that together respond to variations in the heliosphere boundary conditions.  Moebius et al. (Chapter 8) model the heliosphere for several different conditions, and then probe the response of the inner heliosphere to the density of interstellar neutrals flowing into this ISM-modified heliosphere.  At 1 AU, the neutral densities, particle populations derived from interstellar neutrals,  and characteristics of the helium focusing cone all respond to variations in the interstellar boundary conditions. For some cases, increased neutral fluxes fall on the atmosphere of Earth (also see Yeghikyan and Fahr, Chapter 11).  The velocity structure of the ISM appears to vary on subparsec scale lengths (Frisch and Slavin, Chapter 6), and these variations may in some cases result in significant modifications of the inner heliosphere, particularly the gravitational focusing cone, when all other interstellar parameters such as thermal pressure are invariant (Zank et al., Chapter 2, Moebius et al., Chapter 8).  The most readily available diagnostics of the paleoheliosphere are radioisotopes, formed by cosmic ray spallation on the atmosphere, interplanetary and interstellar dust, and meteorites.  Thus, the evaluation of cosmic ray modulation for various types of interstellar cloud boundary conditions is a key part of understanding the paleoclimate records that might trace the solar journey through the Milky Way Galaxy. Fahr et al. (Chapter 9) and Florinski and Zank (Chapter 10) use our understanding of galactic cosmic ray modulation in the modern-day heliosphere as a basis for making detailed calculations of the response of the paleoheliosphere, or the heliosphere as it once was, to the paleolism, or the local interstellar medium that once surrounded the Sun.  The predictions of these calculations are quite intriguing.  Both the termination shock compression ratio and the solar wind turbulence spectrum may vary dramatically with different environments, as mass-loading by pickup ions and the heliosphere properties vary.  The problem of galactic cosmic ray modulation in an ISM-forced heliosphere is extremely important to understanding the paleoheliosphere signature in the terrestrial isotope record.  Today, galactic cosmic rays (GCR) with energies $\\ge 0.25$ GeV penetrate the solar system, and anomalous cosmic rays (energies $<1$ GeV) are formed from accelerated pickup ions. The cosmic ray flux at Earth is sampled by geological radioisotope records, as reviewed Kirkby and Carslaw (Chapter 12, also see Florinski and Zank, Chapter 10). Astronomical data indicates that the Sun has emerged from a region of space with virtually no neutral ISM within the past $\\sim$0.4--1.5 10$^5$ years, and entered the Local Fluff (Chapter 6).  The GCR modulation discontinuity that accompanied this transition may be in the geologic record, which show lower cosmic ray fluxes at Earth, on the average, for the past 135 kyrs years than the 135 kyrs before that (\\cite{Christl:2004}).  \\subsection{Atmosphere Accretion from Dense Cloud Encounters}  Harlow Shapley (1921) suggested \\nocite{Shapley:1921} that an encounter between the Sun and giant dust clouds in Orion may have perturbed the terrestrial climate and caused ice ages.  The discovery of interstellar \\HI\\ and \\HeI\\ inside the heliosphere was soon followed by studies of the ISM influence on the atmosphere for dense cloud conditions (\\cite{Fahr:1968,BegelmanRees:1976,McKayThomas:1978,Thomas:1978,McCrea:1975,TalbotNewman:1977,Willis:1978,ButlerNewmanTalbot:1978}). Yeghikyan and Fahr (Chapter 11), evaluate the density of ISM at the Earth based on models describing the heliosphere inside of an dense cloud, and the interactions between the solar wind and ISM for these dense cloud conditions (also see Chapter 9, by Fahr et al.).  These models then yield the concentration of interstellar hydrogen at the Earth, and the flow of water downward towards the Earth's surface, as a function of the dense cloud density. Significant atmosphere modifications are predicted in some cases. Enhanced neutral populations at 1 AU for a somewhat lower interstellar cloud density regime are discussed in Chapter 8, by Moebius et al.  \\subsection{Possible Effects of Cosmic Rays}  Both solar activity cycles (Figure \\ref{fig:2}) and ISM variations modulate the cosmic ray flux in the heliosphere, and Kirkby and Carslaw (Chapter 12) compare galactic cosmic ray records with paleoclimate archives.  They examine sources of climate forcing such as solar irradiance and cosmic ray fluxes, and conclude that arguments in favor of cosmic ray climate forcing are strong although the mechanism is uncertain.  This relation between cosmic ray flux levels and the climate is shown by radioisotope records and climate archives, such as ice cores, stalagmites, and ice-rafted debris, and for modern times, by historical records. Paleoclimate archives include terrestrial records of cosmic ray spallation in the atmosphere, as traced by radioisotopes with short half-lives (\\tauhalf), e.~g.~ \\Cfourteen\\ (\\tauhalf=5,730 yrs) and \\Beten\\ (\\tauhalf=1.6 Myrs). Possible mechanisms linking the cosmic ray flux at 1 AU and the climate include cloud nucleation by cosmic rays, and the global electrical circuit (see Chapter 12 and \\cite{RobleHaysII:1979}).  The discussion in Chapter 12 provides persuasive evidence linking the surface temperature to cosmic ray fluxes at Earth.  The anticorrelation between sunspot number and cosmic ray fluxes in Figure \\ref{fig:2} shows the heliosphere role in cosmic ray modulation; this mechanism must have also been a prominent mechanism for relating the ISM-modulated heliosphere with the climate.  Fortunately this hypothesis is also verifiable by comparing paleoclimate data with astronomical data on the timing of cloud transitions.  The radioisotope records also indicate that cosmic ray fluctuations have occurred over longer timescales of many $10 ^8$ years. Shaviv compares the \\Clthirtysix\\ (\\tauhalf $\\sim$0.3 Myrs) and \\Kforty\\ (\\tauhalf $\\sim$1.3 Gyr) cosmic ray exposure records in iron meteorites (Chapter 5), but in this case to obtain cosmic ray flux increases due to the Sun's location in spiral arms where active star formation occurs.  A number of studies, none convincing, have invoked the geological \\Beten\\ record, as a proxy for cosmic ray fluxes at Earth, to infer historical encounters with interstellar clouds.  As a way of dating the Loop I supernova remnant, it was suggested that the relative constancy of \\Beten\\ in sea sediments precluded a strong nearby X-ray source within the past $\\sim$2 Myrs (\\cite{Frisch:1981}).  Sonett (1992) \\nocite{Sonett:1992} suggested that peaks in \\Beten\\ layers 35,000 and 65,000 years ago resulted from a compressed heliosphere caused by the passage of a high-velocity interstellar shock.  This extreme heliosphere compression expected for a rapidly moving cloud is supported by heliosphere models (Chapter 2).  Structure in the \\Beten\\ peaks has also been related to spatial structure in the local ISM (\\cite{Frisch:1997}), and solar wind turbulence caused by mass-loading of interstellar neutrals may supply the required mechanism. Global geomagnetic excursions such as the events $\\sim$32 kyr and $\\sim$40 kyr ago also affect the \\Beten\\ record, and can not be ignored (\\nolinebreak \\cite{Christl:2004}).  Indeed, Figure \\ref{fig:2} shows the sensitivity of galactic cosmic ray fluxes on Earth to geomagnetic latitude. ",
        "conclusions": ""
    },
    "0601/astro-ph0601160_arXiv.txt": {
        "abstract": "We study the self-consistent, linear response of a galactic disc to vertical perturbations, as induced say by a tidal interaction. We calculate the self-gravitational potential corresponding to a non-axisymmetric, self-consistent density response of the disc using the Green's function method. The response potential is shown to oppose the perturbation potential because the self-gravity of the disc resists the imposed potential,  and this resistence is stronger in the inner parts of a galactic disc.  For the $m=1$ azimuthal wavenumber, the disc response opposes the imposed perturbation upto a radius that spans a range of 4-6 disc scale-lengths, so that the disc shows a net warp only beyond this region. This physically explains the well-known but so far unexplained observation (Briggs 1990) that warps typically set in beyond this range of radii. We show that the inclusion of a dark matter halo in the calculation only marginally changes (by $\\sim 10\\%$) the radius for the onset of warps. For perturbations with higher azimuthal wavenumbers, the net signature of the vertical perturbations can only be seen at larger radii - for example beyond 7 exponential disc scale-lengths for $m=10$. Also, for high $m$ cases, the magnitude of the negative disc response due to the disc self-gravity is much smaller. This is shown to result in corrugations of the mid-plane density, which explains the puzzling scalloping with $m=10$ detected in HI in the outermost regions $\\sim 30 kpc$ in the Galaxy by Kulkarni et al. (1982). ",
        "introduction": "Galaxy interactions including interactions with satellite galaxies are now known to be common. As a consequence of such a tidal interaction, vertical distortions in the galactic discs can be easily set in. The most common kind of vertical distortions in discs are warps, corresponding to the azimuthal wavenumber $m=1$ . Tidal interactions with neighbouring galaxies have often been suggested as the mechanism for the origin of warps (Schwarz 1985, Zaritsky \\& Rix 1997). The generation of the warp of our Galaxy due to the tidal interaction with the Large Magellanic Clouds(LMC) was studied by Hunter \\& Toomre (1969), and more recently by Weinberg (1995) in the presence of an active dark matter halo. Recent statistical studies by Sanchez-Saavedra et al. (1990) and Reshetnikov \\& Combes (1998) conclude that at least $50\\%$ of all the observed galaxies show warping in their discs. Their study on environmental effects, along with another recent analysis of a large number of edge-on spiral galaxies in the R-band (Schwarzkopf \\& Dettmar 2001), suggests that tidal perturbation (the $m=1$ component) plays an important role in generating and influencing large-scale warps.   While it is clear that a warp is more likely to set in in the outer parts of a galaxy where the tidal force is stronger, it is not yet known  what decides the actual radius at which a warp starts in the disc.  It is a well-known observational fact that warps are seen in the outer regions of galaxies, typically starting at $\\sim 4-6$ exponential disc scale-lengths, as was first seen in HI observations by Briggs (1990). However there exists no clear physical explanation for the radius at which a warp sets in. In fact, little theoretical attention has been devoted to this question.   Other kinds of vertical distortions that are seen are small-scale corrugations in HI in our Galaxy (Quiroga 1974)  or in stars in external galaxies (Florido et al. 1991), and also scalloping in HI in the outer Galaxy (Kulkarni et. al. 1982). Again, the  physical origin of these non-axisymmetric features is not understood.  In this paper, we address these questions and show that the self-gravity of the disc resists any vertical distortion, especially in the inner regions. Thus using very general physical principles, we derive semi-analytically a minimum radius for the onset of each of the above vertical perturbations of a galactic disc. Each Fourier component of a perturbing potential will generate its own signature in the galactic disc. Observations of these signatures in the disc depend on the strength of the perturbation component and the disc response corresponding to it. The behaviour of a disc subjected to an external perturbation can entirely be studied by its response funtion.  We study here the response of an axisymmetric disc, perturbed by a tidal potential and show that the disc self-gravity plays a significant role in sketching a finger-print of the Fourier component of the perturbation onto the disc. We obtain the self-gravitational potential corresponding to a non-axisymmetric, self-consistent, density response of the disc induced by the perturbation component and show that it opposes the perturbing potential. The resulting negative disc response (see Jog 1999 for the planar case) decreases the strength of the perturbation in the disc. The vertical linear disc response obtained here is shown to oppose the perturbation upto a certain radius and for the $m=1$ component, in particular, it is upto about $4$ disc scale-lengths. Thus the net disc response for $m=1$ or warps can only be seen beyond this radius, which is in a fairly good agreement with observations. We also add a rigid dark matter halo and show that it only marginally affects the location of the onset of warps.  For perturbations with higher azimuthal wavenumbers ($m$) the disc is shown to be well resistant upto higher radii. In particular we show that any signature of the $m=10$ component of the perturbation can only be observed in the very outer region of the disc beyond $\\sim 7$ disc scale-length for an exponential disc.  In Section 2, we describe the treatment for the self-consistent linear response of the disc. The results and the comparison with observations are given in Section 3, and the conclusions are summarised in Section 4.  \\bigskip ",
        "conclusions": ""
    },
    "0601/astro-ph0601483_arXiv.txt": {
        "abstract": "Inversion techniques (ITs) allow us to infer the magnetic, dyna\\-mic, and thermal properties of the solar atmosphere from polarization line profiles. In recent years, major progress has come from the application of ITs to state-of-the-art observations. This paper summarizes the main results achieved both in the photosphere and in the chromosphere. It also discusses the challenges facing ITs in the near future.  Understanding the limitations of spectral lines, implementing more complex atmospheric models, and devising efficient strategies of data analysis for upcoming ground-based and space-borne instruments, are among the most important issues that need to be addressed. It is argued that proper interpretations of diffraction-limited Stokes profiles will not be possible without accounting for gradients of the atmospheric parameters along the line of sight. The feasibility of determining gradients in real time from space-borne observations is examined. ",
        "introduction": "We derive information on the physical properties of the solar atmosphere by interpreting the polarization profiles of spectral lines. Extracting this information is not easy, because the observed profiles depend on the atmospheric parameters in a highly non-linear manner through the absorption matrix and the source function vector. To solve the problem, least-squares inversion techniques (ITs) based on analytical or numerical solutions of the radiative transfer equation were developed in the past. These methods compare the observed Stokes profiles with synthetic profiles emerging from an initial guess model atmosphere.  The misfit is used to modify the atmospheric parameters until the synthetic profiles match the observed ones. This yields a model atmosphere capable of explaining the measurements, within the assumptions and limitations of the model.  The first ITs were proposed by \\citet{b3 Ha72} and \\citet{b3 Au77}. Many other codes have been developed since then \\citep[cf.\\ Table 1 in][]{b3 TI03}. Today we have an IT for almost any application we may be interested in: LTE or non-LTE line formation, one-component or multi-component model atmospheres, photospheric or chromospheric lines, etc. These codes have been used intensively to study the magnetism of the solar atmosphere, and now they are essential tools for the analysis of spectro-polarimetric measurements.  This paper concentrates on recent achievements and future challenges of ITs. Additional information on ITs can be found in the reviews by del Toro Iniesta \\& Ruiz Cobo (1996), Socas-Navarro (2001), and del Toro Iniesta (2003). ",
        "conclusions": "Inversion techniques (ITs) have become essential tools to investigate the magnetism of the solar atmosphere. Nowadays, they represent the best option to extract the information contained in high precision polarimetric measurements. The reliability and robustness of least-squares Stokes inversions have been confirmed many times with the help of numerical tests. Part of the community, however, is still concerned with uniqueness issues. These concerns will hopefully disappear with the implementation of more realistic model atmospheres.  During the last years, major progress in the field has resulted from the application of ITs to state-of-the-art observations. The advent of spectro-polarimeters for the near IR has represented a breakthrough, allowing the observation of atomic and molecular lines that provide increased magnetic and thermal sensitivity, and extended chromospheric coverage. The potential of simultaneous observations of visible and IR lines for precise diagnostics of solar magnetic fields has just started to be exploited. Visible and IR lines constrain the range of acceptable solutions, which is especially useful for the investigation of complex structures with different magnetic components and/or discontinuities along the LOS. Finally, we have begun to invert spectro-polarimetric observations at very high spatial resolution, with the aim of reaching the diffraction limit of current solar telescopes (0\\farcs1--0\\farcs2). High spatial resolution allows to separate different magnetic components that might coexist side by side, thus facilitating the determination of their properties.  The application of ITs to these observations is casting doubts on the capabilities of certain lines for investigating particular aspects of solar magnetism. A detailed study of the limitations of spectral lines, in particular the often used \\ion{Fe}{i} pair at $630\\,$nm, seems necessary to clarify their range of usability. An obvious cure for any problem that might affect the observables is to invert visible and IR lines simultaneously. This will require modifications of current ITs to account for different instrumental effects in the different spectral ranges.  Perhaps the most important challenge facing ITs in the next years is the analysis of the enormous data sets expected from upcoming space-borne polarimeters. So far, the efforts have concentrated on the development of fast PCA-methods and ANNs for real-time inversions of the data. However, the unprecedented quality of these observations in terms of spectral and spatial resolution makes it necessary to explore the feasibility of more complex inversions capable of determining gradients of field strength and velocity along the LOS. Tests with numerical simulations demonstrate the importance of gradients to reproduce the very large asymmetries of the Stokes profiles expected at a resolution of 0\\farcs1--0\\farcs2. Current computational resources allow us to determine gradients from full line profiles observed with grating spectro-polarimeters such as the one onboard Solar-B, at a very reasonable cost. Gradients might also be recovered from high-resolution filtergraph observations if sufficient wavelength points are available.  Interestingly, real-time ME inversions of data with limited wavelength sampling seem possible using FPGAs.  The feasibility of such electronic ME inversions needs to be assessed. At the same time, it is important to continue the development of PCA and ANN methods to provide the more complex ME inversions with good initial guesses."
    },
    "0601/astro-ph0601430_arXiv.txt": {
        "abstract": "{The presence of superstarclusters is a characteristic feature of starburst galaxies. We examine the properties of star forming regions and young star clusters in various environments, ranging from common to extreme. We then discuss new high spatial resolution mid-infrared imaging and spectroscopy of extreme superstarclusters in the obscured region of the Antennae ($\\NGC{4038{-}4039}$). We find that the PAH emission in this region is not dominated by the superstarclusters but is mostly diffuse. Emission line ratios found in our high spatial resolution data differ significantly from those in larger apertures, strongly affecting the derived results. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601576_arXiv.txt": {
        "abstract": "Cosmic microwave background (CMB) anisotropy is our richest source of cosmological information; the standard cosmological model was largely established thanks to study of the temperature anisotropies.  By the end of the decade, the Planck satellite will close this important chapter and move us deeper into the new frontier of polarization measurements.  Numerous ground--based and balloon--borne experiments are already forging into this new territory.  Besides providing new and independent information on the primordial density perturbations and cosmological parameters, polarization measurements offer the potential to detect primordial gravity waves, constrain dark energy and measure the neutrino mass scale.  A vigorous experimental program is underway worldwide and heading towards a new satellite mission dedicated to CMB polarization. ",
        "introduction": "Observations of the cosmic microwave background (CMB) anisotropy have driven the remarkable advance of cosmology over the past decade~\\cite{cmb_t}. They tell us that we live in a spatially flat universe where structures form by the gravitational evolution of nearly scale invariant, adiabatic perturbations in a predominant form of non--baryonic cold dark matter. Together with either results from supernovae Ia (SNIa) distance measurements~\\cite{sn}, the determination of the Hubble constant~\\cite{freedman01} or measures of large scale structure~\\cite{lss}, they furthermore demonstrate that a mysterious dark energy (cosmological constant, vacuum energy, quintessence...) dominates the total energy density of our Universe. These observations have established what is routinely called the standard cosmological model: $\\OmM\\approx 0.3=1-\\OmL$, $\\OmB h^2\\approx 0.024$ and $\\Ho\\approx 70$~km/s/Mpc~\\cite{stand_model}. Because the observations in fact over--constrain the model, they test its coherence and its foundations, marking a new era in cosmology.  The CMB results are remarkable for several reasons. They show us density perturbations on superhorizon scales at decoupling and therefore evidence for new physics (inflation or other) working in the early universe.  The observed peaks in the power spectrum confirm the key idea that coherent density perturbations enter the horizon and begin to oscillate as acoustic waves in the primordial plasma prior to recombination; their position justifies the long--standing theoretical preference for flat space with zero curvature. Their heights measure both the total matter and baryonic matter densities, and thereby provide direct evidence that most of the matter is non--baryonic; and, in a scientific {\\em tour de force}, the CMB--determined baryon density broadly agrees with the totally independent estimation from Big Bang Nucleosynthesis~\\cite{bbn}.  These milestones were all obtained from study of the temperature, or total intensity, anisotropies.  The Planck mission\\footnote{\\tt http://www.esa.int/science/planck} (launch 2007/2008) will largely complete this work by decade's end with foreground--limited temperature maps down to $\\sim 5$~arcmin resolution, leaving only the smallest scales unexplored.  With this in mind, the field is already turning to CMB polarization measurements and their wealth of new information. ",
        "conclusions": ""
    },
    "0601/astro-ph0601095_arXiv.txt": {
        "abstract": "We search for signatures of Lorentz and $CPT$ violations in the cosmic microwave background (CMB) temperature and polarization anisotropies by using the WMAP and the 2003 flight of BOOMERANG (B03) data. We note that if the Lorentz and $CPT$ symmetries are broken by a Chern-Simons term in the effective lagrangian, which couples the dual electromagnetic field strength tensor to an external four-vector, the polarization vectors of propagating CMB photons will get rotated. Using the WMAP data alone, one could put an interesting constraint on the size of such a term. Combined with the B03 data, we found that a nonzero rotation angle of the photons is mildly favored: $\\Delta \\alpha= -6.0^{+4.0}_{-4.0}$ $^{+3.9}_{-3.7}$ deg (1, 2 $\\sigma$ ). ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601081_arXiv.txt": {
        "abstract": "\\input abstract.tex ",
        "introduction": "\\label{sec_intro}  \\input sec1.tex ",
        "conclusions": ""
    },
    "0601/astro-ph0601562_arXiv.txt": {
        "abstract": "The last decade has seen applications of Adaptive Mesh Refinement (AMR) methods for a wide range of problems from space physics to cosmology. With the advent of these methods, in which space is discretized into a mesh of many individual cubic elements, the contemporary analog of the extensively studied line radiative transfer (RT) in a semi-infinite slab  is that of RT in a cube. In this study we  provide an approximate solution of the RT equation, as well as analytic expressions for the probability distribution functions (pdfs) of the properties  of photons emerging from a cube,  and compare them with the corresponding slab problem. These pdfs can be used to perform fast resonant-line RT  in optically thick AMR cells where, otherwise, it could take unrealistically long times to transfer even a handful of photons. ",
        "introduction": "\\label{sec:intro} The classic problem of resonant radiative transfer (RT) in a semi-infinite slab has been extensively studied in literature \\citep[e.g.,][]{harrington73,neufeld90}. On the other hand, the last decade has seen applications of Adaptive Mesh Refinement (AMR) methods to problems as diverse as solar physics, supernovae and nucleosynthesis, interstellar medium physics, star formation, astrophysical jets, cosmology, etc. \\cite[for a summary see, e.g.,][]{norman}. In mesh-based methods the continuous domain of interest is discretized into a grid of many individual cubic elements. With the advent of  AMR codes, the contemporary analogue -- at least in terms of usefulness and applicability range -- of the extensively studied problem of resonant-line RT in a slab is the relatively unexplored problem of RT in a cube.  Understanding  resonant RT in optically thick cubes is useful in particular because to  perform RT in  AMR simulations one must solve numerous  cube RT problems, since each time a photon enters an AMR cell one has a new cube RT problem. Furthermore, as is the case, e.g., for \\lya line RT in cosmological simulations of galaxy formation, the high resolution achieved with AMR codes (along with the cooling of the gas), leads to very high scattering optical depths~\\cite[see][]{tasitsiomi05b}. To obtain results within realistic times, we need a very fast RT algorithm, faster than the standard direct Monte Carlo approach in which one follows the detailed scattering of photons in each one of the cells. A way to obtain this  very fast algorithm is to study a priori the RT in cubes of various physical conditions. Using the results of such a study, we can avoid following the detailed photon scattering in each cell. Instead, we can immediately take the photon out of the cell treating thus the problem on a per cell rather than on a per scattering basis, thus accelerating the RT scheme considerably.   In this  letter we  discuss  results on the  resonant RT in  a cube, and in comparison with resonant RT in a slab of similar physical conditions. ",
        "conclusions": ""
    },
    "0601/astro-ph0601048_arXiv.txt": {
        "abstract": "A comparative study of near--maximum--light optical spectra of 24 Type~Ia supernovae (SNe~Ia) is presented.  The spectra are quantified in two ways, and assigned to four groups.  Seven ``core--normal'' SNe~Ia have very similar spectra, except for strong high--velocity Ca~II absorption in SN~2001el.  Seven SNe~Ia are assigned to a ``broad--line'' group, the most extreme of which is SN~1984A.  Five SNe~Ia, including SN~1991bg, are assigned to a ``cool'' group. Five SNe~Ia, including SN~1991T, are assigned to a ``shallow--silicon'' group.  Comparisons with \\syn\\ synthetic spectra provide a basis for discussion of line identifications, and an internally consistent quantification of the maximum--light spectroscopic diversity among SNe~Ia.  The extent to which SN~Ia maximum--light spectra appear to have a continuous distribution of properties, rather than consisting of discrete subtypes, is discussed. ",
        "introduction": "In the first paper of this series (Branch et~al. 2005; hereafter Paper~I), the parameterized supernova synthetic--spectrum code \\syn\\ was used to fit and interpret a time series of spectra of the well--observed, spectroscopically normal, Type~Ia (SN~Ia) SN~1994D. In this paper we concentrate on near--maximum--light spectra of 24 SNe~Ia.  Maximum--light spectra are of particular interest because they are the ones most likely to be obtained, especially for high--redshift SNe~Ia.  In \\S2 the spectra are quantified and arranged to illustrate relationships between them.  In \\S3 to \\S6 four groups of spectra are discussed, and selected comparisons with \\syn\\ synthetic spectra are shown.  Results are discussed in \\S7.  As in Paper~I, we confine our attention to optical spectra, from the Ca~II H\\&K feature in the blue ($\\sim3700$\\ang) to the Ca~II infrared triplet (Ca~II IR3) in the red ($\\sim9000$\\ang).  The 24 spectra selected for study were obtained within three days of $B$--band maximum light [which occurs about 18~days after explosion (Vacca \\& Leibungut (1996)], a time interval short enough that spectroscopic evolution is not strong, yet long enough to allow some events of special interest to be included.  All spectra have been corrected for the redshifts of the parent galaxies.  The original spectra have a range of slopes, owing to intrinsic differences, differences in the amount of reddening by interstellar dust, and observational error.  In this paper, because we are interested only in the spectral features, and not in the underlying continuum slopes, the spectra are tilted by multiplying the flux by $\\lambda^\\alpha$, with $\\alpha$ chosen to make the flux peaks near 4600\\ang\\ and 6300\\ang\\ equally high.  The tilting makes it much easier to compare the spectral features.  The events of the sample are listed in Table~1, together with parameters to be discussed below. ",
        "conclusions": "The high degree of spectral homogeneity among the core--normals suggests that they result from a standard, repeatable SN~Ia explosion mechanism that does not produce spatially large composition inhomogeneities near the characteristic photospheric velocity of 12,000~\\kms\\ (Thomas et~al. 2004).  The strong HV~Ca~II of SN~2001el may involve a spatially large structure near 20,000~\\kms\\ (Kasen et~al. 2003; Kasen \\& Plewa 2005).  On the other hand, the features attributed to HV~Fe~II are only mildly enhanced in SN~2001el and rather homogeneous in the others, suggesting that weak HV~Fe~II is ``natural'' in SNe~Ia, in the sense that it does not require unusual high--velocity structures (H\\\"oflich, Wheeler, \\& Thielemann 1998, Lentz et~al. 2001b).  Weak HV~Ca~II also is ubiquitous (Mazzali et ~al. 2005b).  Based on previous work (Branch 1987; Benetti et~al. 2004) we expected to fit the broad--line SNe~Ia with high values of \\vphot, but better fits were obtained by holding \\vphot\\ constant at 12,000~\\kms, a typical value for the core--normals, and using increased \\ve\\ values. This may be a consequence of nuclear burning extending to higher velocity in the broads than in the core--normals (Lentz et~al. 2001a; Benetti et al. 2004; Stehle et~al. 2005).  Although the composition structures of the cools surely differ from those of the core--normals (e.g., the low luminosity requires a low mass of $^{56}$Ni), the immediate reason for the spectroscopic differences appears to be the lower temperature.  The blue trough in cools, usually referred to as the Ti~II trough, is produced not only by Ti~II, but also by other ions that appear at low temperature. Ti~II is not significant for the 5750\\ang\\ absorption.  The 5750\\ang\\ absorption of the cools can be accounted for in terms of a high Si~II optical depth together with a maximum Si~II velocity.  This is a plausible way to simultaneously fit these two absorptions while explaining why \\rsi\\ and the ratio of \\575\\ to \\610\\ go the wrong way with temperature (i.e., increasing with falling temperature), considering that the lower level of \\lam5972 is the upper level of \\lam6355. If correct, this maximum velocity would indicate that in the cools nuclear burning extends only to unusually low velocities, near 15,000~\\kms.  In most of the shallow--silicon SNe~Ia the spectral features appear to be produced by the same ions as in the core--normals, although the optical depths are quite different, especially in SN~1991T and SN~2002cx.  Most of the differences from core--normals may be caused by a higher temperature.  Like SN~2001el among the core--normals, SN~1999ee provides support for the HV~Fe~II identification.  The cause of the depression of the spectra of SN~2000cx in the blue from about 3800\\ang\\ to 4200\\ang\\ appears to be extra HV line blocking.  If HV Ti~II is present, then HV~Cr~II is a more plausible identification than HV H$\\beta$ (Branch et~al. 2004a) for the 4530\\ang\\ absorption in SN~2000cx.  As noted in \\S3, a less distinct absorption appears near 4550\\ang\\ even in the core--normals; HV Cr~II is a possibility.  One of the most interesting issues that can be addressed by studies of this kind is whether SNe~Ia have a continuous distribution of properties, or consist of discrete subgroups.  On the whole, we have a sense of continuity.  SN~1992A and SN~1981B may be small steps from the core in the direction of SN~1984A; SN~1989B may be a step in the direction of SN~1986G and SN~1991bg; SN~1999ee may be a step in the direction of SN~1991T.  If the distribution is continuous, then exactly how the boundaries of the core--normal group are defined is not important.  The cools stand apart from the others in Figure~1, but this in itself does not establish that they differ in kind from the others.  As noted by Hatano et~al. (2002) and H\\\"oflich et~al. (2002), there is a temperature threshold below which, owing to abrupt changes in key ionization ratios, line optical depths change abruptly (Hatano et~al. 1999).  Therefore, the temperature interval required to go from SN~1989B through SN~1986G to SN~1991bg--like events may be narrow. SN~1991bg--like events also are quite subluminous, but especially in the $B$ band, which is strongly blocked by the blue trough. Similarly, there is a temperature threshold above which Fe~III features become conspicuous (Hatano et~al. 2002).  Most SNe~Ia may be variations on the basic core--normal theme, with the diversity reflecting differences of degree, rather than differences of kind.  Evidence recently has been presented that in terms of stellar population age there are two SN~Ia groups of roughly equal size, one from a young ($\\sim10^8$ years) population and another from another from an older ($\\sim3$~Gyr) population (Mannucci et~al. 2005a; Scannapieco \\& Bildsten 2005; Mannucci, Della Valle, \\& Panagia 2005b).  However, we see no basis for splitting the SNe~Ia of our sample into two groups of roughly equal size.  It may be that there are two different binary--star evolution paths to SNe~Ia, one requiring a short time and the other a long time, but with the two paths producing similar outcomes, i.e., carbon--oxygen white dwarfs that explode as they approach the Chandrasekhar mass.  The one event of our sample that does seem to be different in kind is SN~2002cx, which we have placed in the shallow--silicon group. Like SN~1991T it had conspicuous Fe~III lines, but unlike SN~1991T it did not have a slow light--curve decline rate and it was subluminous (Li~et~al. 2003).  Its late--time spectra were unlike those of normal SNe~Ia, being dominated by very low velocity ($\\sim$700~\\kms) permitted lines of Fe~II, Ca~II, Na~I, and perhaps O~I.  Several other SN~2002cx--like events have been discovered (S.~Jha et~al., in preparation).  On the basis of the first--season observations by Li~et~al. (2003), Branch et~al. (2004b) suggested that SN~2002cx may have been a pure deflagration while most other SNe~Ia, even SN~1991bg--likes (H\\\"oflich et~al. 2002), may be delayed detonations or gravitationally confined detonations (Plewa, Calder, \\& Lamb 2004; Kasen \\& Plewa 2005). However, the late--time spectra of SN~2002cx (S.~Jha et~al., in preparation) do not resemble the spectra calculated by Kozma et~al. (2005) for one particular 3--D deflagration model, so the nature of SN~2002cx remains uncertain.  The normality of PV features in events such as the core--normal SN~2001el and the shallow--silicon SN~2000cx, both of which have strong HV features, suggests that HV and PV are somewhat independent (with HV~Fe~II and weak HV~Ca~II being ubiquitous and natural, as discussed above).  An interesting issue for further study is how much of the photometric diversity of SNe~Ia is caused by HV features (which may not correlate with luminosity).  Another issue is whether the strong HV blue line blocking of SN~2000cx occurs only in the shallow--silicon group.  It is clear that strong HV~Ca~II IR3 absorption occurs not only in the core--normal SN~2001el but also in the shallows.  Our study is complementary to that of Benetti et~al. (2005; hereafter Be05), who did not consider the strenths of spectral features. Sixteen events of our sample were included in the study of Be05, who assigned SNe~Ia to three groups on the basis of absolute magnitude ($M_B$); light--curve decline rate (\\dm15; Phillips 1993); the blueshift of the 6100\\ang\\ absorption 10 days after maximum [\\v10; Branch \\& van~den~Bergh 1993]; the rate at which the 6100\\ang\\ absorption drifts to the red [\\vdot; Be05]; and \\rsi.  With a few borderline exceptions (SN~1989B and SN~1992A), our cools corresponds to the Be05 FAINT group, which is to be expected since both temperature and luminosity are controlled mainly by the $^{56}$Ni mass.  Our broad--line group corresponds to the Be05 HVG (high temporal--velocity gradient) group. This also makes sense because broad 6100\\ang\\ absorption requires high Si~II optical depth over a substantial velocity range, which makes it possible for the absorption minimum to shift appreciably with time. The broads appear to have thicker silicon layers than the core--normals.  Our core normals and shallows correspond to the Be05 LVG (low temporal--velocity gradient) group.  The lower velocity range over which Si~II has a high optical depth permits only a smaller shift in the absorption minimum with time.  Polarization observations (Wang et~al. 2001; Howell et~al. 2001; Wang et~al. 2003; Leonard et~al. 2005) show that some spectral features are affected by asymmetric structures.  A straightforward way to rule out asymmetry as the sole cause of the differences between groups of SNe~Ia would be to demonstrate a correlation with parent galaxy type (Be05) or other indicators of the parent--galaxy population (Gallagher et~al. 2005).  If asymmetry were the sole cause of SN~Ia differences, SN~Ia subtypes would not correlate with galaxy types.  But, in fact, there are correlations and emerging correlations.  SN~1991bg--like events tend to occur in older populations (Howell 2001).  In the Be05 sample the mean host galaxy morphological type (the T~type) was found to increase from the FAINT to HVG to LVG groups.  Similarly, for our sample the mean T~type increases from cools to core--normals to broads to shallows, but the statistics are too poor to allow definite conclusions.  The expected rapid increase in the number of well observed SNe~Ia in the near future (e.g., Wood--Vasey et~al. 2004) will allow the relationships among SNe~Ia to be clarified.  We expect that our Paper~III (in preparation), on premaximum spectra, also will clarify the view presented here.  We are grateful to all observers who have provided spectra.  This work has been supported by NSF grants AST-0204771 and AST-0506028, and NASA LTSA grant NNG04GD36G.   \\appendix"
    },
    "0601/astro-ph0601424_arXiv.txt": {
        "abstract": "We compute the power spectrum of the Cosmic Microwave Background temperature anisotropies generated by the Intergalactic Medium. To estimate the electron pressure along the line of sight and its contribution to the Sunyaev-Zeldovich component of the CMB anisotropies, we assume the non-linear baryonic density contrast is well described by a log-normal distribution. For model parameters in agreement with observations and for an experiment operating in the Rayleigh-Jeans regime, the largest IGM contribution corresponds to scales $l\\approx 2000$. The amplitude is rather uncertain and  could be as large as $100-200\\mu$K$^2$, comparable to the contribution of galaxy clusters. The actual value is strongly dependent on the gas polytropic index $\\gamma$, the amplitude of the matter power spectrum $\\sigma_8$, namely $C_l^{IGM}\\sim(\\gamma^2\\sigma_8)^{12}$. At all redshifts, the largest contribution comes from scales very close to the baryon Jeans length. This scale is not resolved in numerical simulations that follow the evolution of gas on cosmological scales. The anisotropy generated by the Intergalactic Medium could make compatible the excess of power measured by Cosmic Background Imager (CBI) on scales of $l\\ge 2000$ with $\\sigma_8=0.9$. Taking the CBI result as an upper limit, the polytropic index can be constraint to $\\gamma < 1.5$ at $2\\sigma$ level at redshifts $z\\sim 0.1-0.4$. With its large frequency coverage, the PLANCK satellite will be able to measure  the secondary anisotropies coming from hot gas. Cluster and Intergalactic Medium contributions could be separated by cross correlating galaxy/cluster catalogs with CMB maps. This measurement will determine the state of the gas at low and intermediate redshifts. ",
        "introduction": "In the last decade, numerical simulations (Cen \\& Ostriker 1999; Dave et al. 1999, 2001) and observational evidence (Rauch 1998; Stocke, Shull \\& Penton 2004) indicate that the highly ionized intergalactic gas has evolved from the initial density perturbations into a complex network of mildly non-linear structures in the redshift interval $0<z<6$, called cosmic web. This structure could contain most of the baryons in the Universe (Rauch et al. 1997; Schaye 2001; Fukugita \\& Peebles, 2004). With cosmic evolution, the fraction of baryons in these structures decreases as more matter is concentrated within compact virialized objects. The Ly$\\alpha$ forest absorbers at low redshifts are filaments (with low HI column densities) containing about 30\\%  of all baryons (Stocke et al. 2004). Hydrodynamical simulations predict that another large fraction of all baryons resides within mildly-nonlinear structures which are partly shock-confined gas filaments heated up to temperatures of $10^5-10^7 K$, called  Warm-Hot Intergalactic Medium (WHIM). The amount of baryons within this WHIM is estimated to reach 20 - 40 \\% (Cen \\& Ostriker 1999; Dave et al. 1999, 2001).  In this article we compute the electron pressure along the line of sight of both the ionized gas in the Ly$\\alpha$ forest and the hot gas in the WHIM. Even if the temperature of the gas in these absorbing filaments is relatively low (usually less than $10^5$K) the total amount of ionized gas is large. This electron pressure induces distortions and temperature anisotropies on the Cosmic Microwave Background (CMB) spectrum by the Thermal Sunyaev-Zeldovich (TSZ) effect (Sunyaev \\& Zeldovich, 1972, 1980) as do clusters of galaxies. Early calculations of the TSZ power spectrum used analytical approaches (Atrio-Barandela \\& M\\\"ucket 1999; Komatsu \\& Kitayama 1999, Hern\\'andez-Monteagudo, Atrio-Barandela \\& M\\\"ucket 2000; Molnar \\& Birkinshaw 2000). Numerical simulations were soon used to make more accurate predictions (Refregier et al. 2000, Refregier \\& Teyssier 2002, Zhang, Peng \\& Wang 2002). Da Silva et al. (2001) studied the relative contribution from high and low density gas and found that in all models considered, the high density gas in halos dominated the TSZ signal. White et al. (2002) noticed that at $l=2000$ about 25\\% of the signal came from gas in regions with density smaller than one hundred times the cosmological mean. At $l=6000$ the contribution from the diffuse gas was less than 2\\%.  Numerical simulations currently agree with the predictions based on the halo model and it is reasonable to expect that the physics of baryons on dense environments has been well described in cosmological simulations. What has not been proven is whether numerical simulations have reached convergence, i.e, if further increments in resolution will not increase the signal (Bond et al. 2005). To study numerically the evolution of the complex network of filaments that develops when the evolution of the gas in low density regions becomes non-linear, requires computer resources that are at present unavailable. Soft Particle Hydrodynamics (SPH) codes do not have enough resolution and Adaptive Refinement Methods (ARM) are controlled by density and follow the evolution of high dense regions. To study the low dense regions with similar accuracy is computationally very expensive.   In this article we study the contribution coming from baryons located in the low dense regions that constitute the IGM. The web of cosmic filaments corresponds to scales in between large scale fluctuations and strongly nonlinear scales inside halos. The halo model is likely to fail to account for its complicated geometry and an analytical treatment would require a different approach. We shall use a log-normal probability distribution function (PDF) to describe the non-linear evolution of baryons in low density environments. Briefly, in Sec 2  we describe the model and in Sec 3 we derive the expressions that give the contribution of the IGM to the average y-parameter and CMB temperature anisotropies.  In Sec. 4 we discuss our results and their dependence with cosmological and physical parameters. In Sec. 5 we present our conclusions and the observational prospects to measure the anisotropies generated by the IGM. ",
        "conclusions": "In this article we have shown that ionized gas, in the deep potential wells of clusters of galaxies and in the IGM, have a significative effect on the CMB, generating both temperature anisotropies and spectral distortions. While the latter is proportional to the electron pressure along the line of sight, the former depends on the clustering properties of the hot  gas  (Hern\\'andez-Monteagudo et al. 2000). The high amplitude of the radiation power spectrum, similar in magnitude to the contribution of clusters, is due to the non-linear evolution and high degree of clustering of the baryonic matter at low redshift. It is the lognormal distribution of the ionized gas that drives the strong dependence of the amplitude of the radiation power spectrum with $\\gamma$ and $\\sigma_8$. This extra component can make compatible the CBI measured power excess at $l\\sim 2000$ with $\\sigma_8=0.9$. At all redshifts, the largest contribution comes from scales around the baryon Jeans length. In cosmological simulations set to compute the TSZ contribution to the Cosmic Microwave Background temperature anisotropies, this scale is not well resolved, what explains why this contribution has not yet been identified in numerical simulations.  The log-normal model is remarkable for predicting a strong dependence of the radiation power spectrum with two parameters: $C_{l,max}\\sim (\\gamma^2\\sigma_8)^{12}$. Even if the range of scales, densities and redshifts to which the model can be applied is rather uncertain, those parameters produce minimal variations in shape of the radiation power spectrum and the variations in the amplitude are much smaller than that of $\\gamma$ and $\\sigma_8$. Remarkably, the anisotropy due to the IGM is generated in a narrow redshift interval.  Assuming $\\sigma_8=0.9$, one can obtain a firm upper limit of $\\gamma \\le 1.5$ at the 2$\\sigma$ confidence level in the redshift range [0.1-0.4]. As different redshift intervals dominate the anisotropy at different angular scales, measurements of the power spectrum at different $l$ will allow to determine the polytropic index $\\gamma$ and the state of the ionized gas at different redshifts.  By hypothesis, the gas obeys a polytropic equation of state. This cannot account for the effect of the shock heated gas of the WHIM leading to high temperatures ($T \\approx 10^6 - 10^7 K$). Thus, this baryon component might provide an extra TSZ contribution at redshifts close to zero. To distinguish the IGM TSZ effect from the one coming from clusters of galaxies, it is necessary to use their different statistical properties. Since the TSZ effect is independent of redshift, the cluster signal will correlate with cluster positions on the sky. Cross-correlation of cluster catalogs with CMB maps  opens the possibility of determining the cluster contribution and to separate the IGM component. Hern\\'andez-Monteagudo, G\\'enova-Santos \\& Atrio-Barandela (2004) and Afshordi, Lin \\& Sanderson (2005) have recently carried out such an analysis using WMAP data. They found strong evidence (at the 5 and 8$\\sigma$  levels, respectively) of a TSZ contribution to the radiation power spectrum due to clusters but the data was not sensitive enough to yield the radiation power spectrum. The PLANCK satellite, with its large frequency coverage, will be well suited for measuring the TSZ power spectrum. It is worth to explore correlation techniques that will permit to separate the cluster from the IGM component. Measurements of the IGM power spectrum at different multipoles will provide a measurement of the state of the gas (temperature and polytropic index) at different redshifts."
    },
    "0601/astro-ph0601338_arXiv.txt": {
        "abstract": "We study the evolution of the stellar mass density for the separate families of bulge--dominated  and disk--dominated galaxies over the redshift range  $0.25\\le z\\le1.15$. We derive quantitative morphology for a statistically significant galaxy sample of 1645 objects selected from the FORS Deep and the GOODS--S Fields. We find that the morphological mix evolves monotonically with time: the higher the redshift, the more disk systems dominate the total mass content. At $z\\sim1$, massive objects (M$_*\\ge7\\times10^{10}$M$_\\odot$) host about half of the mass contained in objects of similar mass in the local universe. The contribution from early and late type galaxies to the mass budget at $z\\sim1$ is nearly equal. We show that {\\it in situ} star formation is not sufficient to explain the changing mass budget. Moreover we find that the star formation rate per unit stellar mass of massive galaxies increases with redshift only for the intermediate and early morphological types, while it stays nearly constant for late-type objects. This suggests that merging and/or frequent accretion of small mass objects has a key role in the shaping of the Hubble sequence as we observe it now, and also in decreasing the star formation activity of the bulge--dominated descendants of massive disk galaxies. ",
        "introduction": "\\label{sec_introduction}  \\citet{lilly1996}, and soon after \\citet{madau1996}, produced the first estimates of the cosmic star formation history. Since then, the growing number of deep extragalactic surveys has allowed galaxy evolution to be tackled in more and more detail, to higher and higher redshifts \\citep[e.g.][]{giavalisco2004, gabasch:sfr, bouw2004,juneau2005}. All these studies point towards a substantial amount of star formation at early cosmic epochs.  More recently, a number of studies, mainly based on NIR selected surveys like  2MASS \\citep{cole2001}, MUNICS \\citep{drory2004a}, K20 \\citep{fontana2004}, FIRES \\citep{rudnick2003}, HDF \\citep{dickinson2003} and GOODS+FDF \\citep{drory2005}, have been able to measure directly the stellar mass density up to high redshifts. The two independent approaches are in a remarkably good agreement, and suggest that about half of the present--day stars was already in place at z $\\approx$ 1, when the Universe was half of its present age.  The assembly of the stellar mass through cosmic time is a crucial test for models of galaxy formation and evolution \\citep{KC1998}. Such models aim at linking the hierarchical growth of dark matter structures to the observed galaxy properties, by means of simplified prescriptions for the formation of baryonic systems in dark matter haloes \\citep[e.g.][]{cole2000,gd2004}. %  An effective way to put sensitive constraints on the models is to understand where the stars were located, specifically what was the morphology of their host galaxies, at different look-back times. In fact, this permits to directly probe how galaxies assembled their stars, and how their morphology (i.e., at least with some approximation, their dynamical status) evolves. Since merging is driving both mass assembly and dynamical evolution in a hierarchical scenario, these studies offer a direct probe of the role of this process in galaxy evolution.  Because of the lack of sufficient angular resolution, galaxies have been often classified, both at low and high redshift, as early or late types based on their broad band colors \\citep{baldry2004,fontana2004}. However this kind of approach suffers, especially at high redshift, from the obvious drawback that a disk galaxy populated by an old stellar population would be classified as an early--type object (and vice versa). \\citet{bell2003} used the concentration parameter to discriminate between early and late type objects in a complete sample extracted from local surveys.  They found that in the local Universe the {\\em transition mass}, i.e. the mass at which disks become dominant in the relative contribution to the total stellar mass function, is $\\approx5\\times10^{10}$M$_\\odot$. %   In this work we rely on a deep, complete and automatically morphologically classified sample to study the contribution of galaxies of different morphologies to the redshift evolution of the stellar mass density. \\citet{bundy2005}, and also \\citet{brinchmanneellis2000}, have carried out a work qualitatively similar to the one we present here, but based on shallower samples and visual morphological classification.%   This Letter is organized as follows: in \\S \\ref{sec_data} we discuss the dataset on which this work is based, in \\S \\ref{sec_sersic} we present the surface brightness profiles modeling, and in \\S \\ref{sec_mass} our results on the evolution of mass functions, mass densities and specific star formation rates (SSFR, i.e. star formation rate per unit stellar mass) split by morphological type. Finally, in \\S \\ref{sec_conclusions} we draw our conclusions.  We use AB magnitudes and adopt a cosmology with $\\Omega_M$=0.3, $\\Omega_\\Lambda$=0.7, and H$_0$=70 kms$^{-1}$Mpc$^{-1}$. ",
        "conclusions": "\\label{sec_conclusions} In this Letter we have studied the contribution of disks and bulges to the evolution of the stellar mass function up to $z\\sim1$.  We agree with \\citet{bundy2005} that the transition mass, i.e. the mass at which disks become dominant in the relative contribution to the total mass function, increases with redshift. We estimate that at $z\\sim1$ the transition mass is up to a factor 2 larger than its measured local value.  We show that the morphological mix evolves with redshift. At $z\\sim1$ early and late type galaxies contribute nearly equally to the total mass budget in massive objects (i.e. with M$_*\\ge7\\times10^{10}$M$_\\odot$). We suggest that merging events must play a key role in the {\\it mass pouring} from disks to bulges.  We find a different behavior of the SSFR, i.e. the star formation rate per unit stellar mass, as a function of redshift for the different morphological types. The median SSFR of late-type objects shows almost no evolution up to $z\\sim1$. Conversely, median SSFRs for early and intermediate types increase systematically. Since these latter morphological types are dominating in numbers, they drive the total SFR evolution of massive galaxies. This suggests a scenario where the morphological evolution of massive disk galaxies through merging is followed by the decrease of the star formation in their bulge--dominated descendants, maybe after a burst of star formation that exhausts the available gas."
    },
    "0601/astro-ph0601497_arXiv.txt": {
        "abstract": "{ We review the most recent observational results on the formation and evolution of early-type galaxies and their mass assembly by focusing on: the existence, properties and role of distant old, massive, passive systems to $z\\sim2$, the stellar mass function evolution, the ``downsizing'' scenario, and the high-$z$ precursors of massive early-type galaxies. ",
        "introduction": "Early-type galaxies (ETGs) play a crucial role in cosmology. They represent the most massive galaxies in the local Universe, contain most of the stellar mass (${\\cal M}$) and are the primary probes to investigate the cosmic history of galaxy mass assembly. As ETGs are the most clustered galaxies, they are also fundamental in tracing the evolution of the large scale structures. Because of the correlation between the black hole and galaxy bulge masses, ETGs are also important to investigate the co--evolution of normal and active galaxies.  Although ETGs at $z\\sim0$ are rather ``simple'' and homogeneous systems in terms of morphology, colors, stellar population content and scaling relations \\cite{alvio}, their formation and evolution is still a debated question. In the so called \"monolithic'' collapse scenario \\cite{egg, tinsley}, ETGs converted most of the gas mass into stars at high redshifts (e.g. $z>3$) through an intense burst of star formation followed by passive evolution of the stellar population. However, in the modern scenario of CDM cosmology, the evolution of the gravitational perturbations implies that galaxies assembled their mass more gradually through hierarchical merging of CDM halos and make the monolithic collapse unlikely or even unphysical. For instance, the most recent $\\Lambda$CDM hierarchical models predict that most of the ETG stellar mass is assembled at $0<z<1$ \\cite{delucia}. The modern generation of galaxy surveys can test the different predictions on the evolution of the number density, colors, merger types and rates, luminosity and mass functions, and provide a crucial feedback to improve the theoretical simulations.  Despite the intense observational activity on ETG evolution, a general consensus has not been reached yet. For instance, recent results based on optically-selected, moderately deep ($R<24$) samples of ``red sequence'' ETGs (e.g. COMBO-17, DEEP2) \\cite{bell04,faber} suggest a strong evolution of the number density with $\\phi^*(z)$ displaying an increase by a factor of $\\sim$6 from $z\\sim1$ to $z\\sim0$, and a corresponding increase by a factor of 2 of their stellar mass density since $z\\sim1$. Dissipationless ETG--ETG merging (also called \"dry\" merging) is advocated to explain how ETGs assembled at $0<z<1$ \\cite{bell05,vd05}. HST optical, moderately deep ($I<24$) samples of morphologically--selected ETGs also suggest a substantial decrease of the number density with $n \\propto (1+z)^{-2.5}$ \\cite{ferre}, and spatially resolved color maps show that $\\sim$30-40\\% of ETGs at $0<z<1.2$ display variations in their internal color properties (e.g. blue cores and inverse color gradients) suggestive of recent star formation activity in a fraction of the ETG population \\cite{men}. However, the optically-selected ($I<24$) VIMOS VLT Deep Survey (VVDS) does not confirm the above findings, and shows that the rest-frame $B$-band luminosity function of ETGs (selected based on the SED) is consistent with passive evolution up to $z\\sim1.1$, while the number of bright ETGs seems to decrease only by $\\sim$40\\% from $z\\sim0.3$ to $z\\sim1.1$ \\cite{zucca}. The overall picture is made more controversial by recent results indicating a weak evolution of the stellar mass function \\cite{fon04,caputi,bundy} and that ``dry mergers'' cannot represent the major contributors to the assembly of ETGs at $0<z<1$ \\cite{bundy,masjedi}. ",
        "conclusions": ""
    },
    "0601/astro-ph0601668_arXiv.txt": {
        "abstract": "We present numerical simulations of convective overshoot in a two-dimensional model of the solar equatorial plane.  The simulated domain extends from 0.001 R$_{\\odot}$ to 0.93 R$_{\\odot}$, spanning both convective and radiative regions.  We show that convective penetration leads to a slightly extended, mildly subadiabatic temperature gradient beneath the convection zone below which there is a rapid transition to a strongly subadiabatic region.  A slightly higher temperature is maintained in the overshoot region by adiabatic heating from overshooting plumes.  This enhanced temperature may partially account for the sound speed discrepancy between the standard solar model and helioseismology.  Simulations conducted with tracer particles suggest that a fully mixed region exists down to at least 0.687 R$_{\\odot}$. ",
        "introduction": "One of the main unsolved problems in all of stellar evolution theory is the treatment of convective-radiative boundaries.  In the Sun, understanding convective overshoot is crucial in determining properties of the solar tachocline and the solar dynamo.  Helioseismic observations have shown that the differential rotation observed at the solar surface persists throughout the convection zone (Thompson et al. 1996).  More surprisingly, the p-mode splittings showed that over a very thin radial extent the rotation profile changes from differential in the convection zone, to solid body in the radiative interior.  The unresolved radius over which this transition occurs has been named the tachocline (Spiegel \\& Zahn 1992).  Whether this region is quiescent and mostly devoid of overshooting motions or is violent and constantly bombarded by plumes is an unanswered question, the solution to which may help constrain theoretical models for the tachocline.  Understanding the tachocline is not only important for comprehending the internal rotation of the Sun; it is crucial for understanding the dynamo process.  In classic mean field theory, the dynamo process is explained in two steps: poloidal field is sheared into toroidal field by differential rotation (the $\\Omega$ effect), toroidal field then buoyantly rises and, because of Coriolis forces is twisted back into a poloidal field (the $\\alpha$ effect).  These processes were initially postulated to occur in the convection zone.  However, it became clear some time ago (Parker 1975) that magnetic field would become buoyant and be quickly shredded within the turbulent convection zone.  During the ensuing years it was proposed that the storage and amplification of field could occur in the stably-stratified overshoot region (Spiegel \\& Weiss 1980).  The tachocline, with its strong differential rotation would provide the ideal site for storage and amplification of the field.  To explain how the $\\alpha\\Omega$ dynamo would work in this scenario Parker (1993) proposed the interface dynamo, which places the $\\Omega$ effect in the stable region beneath the convection zone, while keeping the alpha effect in the bulk of the convection zone.  If the interface dynamo is to work, the overshoot region must play two crucial roles in the dynamo cycle: (1) the stable stratification allows field, which is transported into the region by overshooting plumes (Tobias et al. 1998), to be stored there on the solar cycle timescale and (2) the strong shear in this region must convert poloidal field into toroidal field.  Understanding the precise nature of this overshoot region, the amplitude of the subadiabaticity and its depth, is crucial for understanding the efficiency with which field can be pumped into the overshoot region, the magnitude of the field capable of being stored there and the timescale on which it could be stored.  The problem of overshoot is not confined to the Sun.  Most stages of stellar evolution are affected.  The Lithium depletion in some main sequence stars may be explained by convective overshoot.  In more massive stars with convective cores, overshoot can lead to increased central hydrogen abundance, and therefore longer main sequence lifetimes, which in turn affects isochrone fitting and age predictions for stellar clusters.  Dredge up during the Asymptotic Giant Branch (AGB) phase can lead to the surface enrichment of material from the core and to the presence of $^{13}C$ necessary for s-process nucleosynthesis (Herwig 2000).  Convective overshoot can affect the energy production and nucleosynthesis of classical novae by enriching the accreted layer with underlying white dwarf material (Woosley 1986).  Clearly, overshoot is an important physical process which must be explained if stellar evolution is to be better understood.  Many analytic and numerical theories have been proposed, but no clear consensus has been reached.  Early analytic models (Shaviv \\& Salpeter 1973, van Ballegoojen 1982, Schmitt, Rosner \\& Bohn 1984) predicted extended adiabatic regions and early numerical experiments concurred (Hurlburt et al. 1994).  More recent analytic results (Zahn 1991, Rieutord \\& Zahn 1995, Rempel 2004) model the penetration using plume dynamics and agree that the extent and nature of the overshoot depends critically on the filling factor of the plumes at the base of the convection zone as well as the flux.  It also appears that the inclusion of upflow-downflow interaction, and the prescription for it, affects the penetration depth (Rieutord \\& Zahn 1995, Rempel 2004).  More recent numerical experiments at higher spatial resolution in two- (Rogers \\& Glatzmaier 2005a) and three dimensions (Brummell et al. 2002) have found subadiabatic overshoot but no extended adiabatic region.  All analytic predictions make some crude assumptions, while all numerical simulations employ far from realistic solar parameters.  The numerical simulations presented in this work still can not reach the level of turbulence that surely exists in the Sun.  However, in this work we have taken a few steps forward by using a realistic thermal profile, including rotation, employing a more realistic geometry and including most of the convective and radiative regions.  In section 2 we describe our numerical model and equations; in section 3 we discuss some of the features of basic convective penetration.  In section 4 we present estimates for the depth of the penetrative convection before our reference state model is evolved.  Section 5 details the evolution of the mean thermal profile and section 6 compares these results with previous analytic and numerical models.  In section 7 we discuss mixing of species as modeled by tracer particles. ",
        "conclusions": "We have presented self-consistent, numerical simulations of convective overshoot in a 2D model of the solar equatorial plane.  This model employs a realistic thermal reference state and evolves in response to convection and overshoot.  We find that the convective penetration leads to a thermal profile which is mildly subadiabatic in a shallow region at the base of the convection zone.  Beneath this mildly subadiabatic region the profile drops precipitously to the extreme subadiabaticity of the radiative interior.  Overshooting plumes cause a mild heating in the overshoot region, leading to a region which has a slightly higher temperature and a greater subadiabaticity than the standard solar model.  The region of increased temperature is at the location and in the sense of the helioseismology-standard solar model sound speed discrepancy and may provide an alternative explanation for that difference.  Passive tracer particles mix only slightly below the convective-radiative interface over the two turnover times followed.  Over longer times, a fully mixed region can be expected at least down to a radius of 0.687 R$_{\\odot}$.  At this radius, the Lithium burning timescale is 8 billion years, a factor of 4 to 5 too long.  However, only two turnover times were simulated with tracer particles and its possible that given longer integration time particles initiated in the convection zone could plunge deeper, thus leading to a shorter burning time.  In order to study this further, longer time integrations are needed.  Furthermore, studies varying the number of particles to determine the minimum radius at which there cease to be particles should be conducted.  We should bare in mind that these results are 2D and far from the turbulent nature of the Sun.  It is unclear what three dimensional effects would be.  While it seems clear that 3D models have smaller filling factors, it is unclear what this affect will be and it depends not only on the filling factor but the number of plumes accross which that filling factor is distributed.  Furthermore, it is not clear what the effect of meridional circulation or the full Coriolis force would be on the penetration depth and mixing.  Studies on the effect of Rayleigh number (turbulence level) on the penetration depth indicate that increasing the Rayleigh number decreases the penetration depth \\citep{bct02, rg05a}.  However, it is unclear how these parametrized results come into play in more realistic models.  The jury is still out and awaits more sophisticated, 3D numerical simulations."
    },
    "0601/astro-ph0601518_arXiv.txt": {
        "abstract": "{We have surveyed a sample of 165 solar-type spectroscopic binaries (SB) with periods from 1 to 30 days for higher-order multiplicity. A subsample of 62 targets were observed with the NACO adaptive optics system and 13 new physical tertiary companions were detected.  An additional 12 new wide companions (5 still tentative) were found using the 2MASS all-sky survey.  The binaries belong to 161 stellar systems; of these 64 are triple, 11 quadruple and 7 quintuple. After correction for incompleteness, the fraction of SBs with additional companions is found to be 63\\% $\\pm$ 5\\%.  We find that this fraction is a strong function of the SB period $P$, reaching 96\\% for $P<3^d$ and dropping to 34\\% for $P>12^d$. Period distributions of SBs with and without tertiaries are significantly different, but their mass ratio distributions are identical. The statistical data on the multiplicity of close SBs presented in this paper indicates that the periods and mass ratios of SBs were established very early, but the periods of SB systems with triples were further shortened by angular momentum exchange with companions. ",
        "introduction": "\\label{sec:intro}  The formation of binary stars remains a  subject of active research and debate (Zinnecker \\& Mathieu \\cite{IAU200}).  Close binaries are particularly difficult to explain because at orbital periods of a few days, the separations are less than the size of the individual components during  the protostellar phase.  Fission of a rapidly rotating pre-main sequence stellar configuration appears unlikely according to the most recent theoretical results (Tohline \\cite{Toh02}).  Thus, some mechanism(s) of orbit shrinkage must be present in order to extract the angular momentum from a proto-binary. This momentum could be deposited in wider stellar companions -- a hypothesis we test in this paper.   Stars form in groups and interact dynamically. Detailed studies of dynamical decay of small unstable clusters, e.g. by Sterzik \\& Durisen (\\cite{SD98}), were successful in explaining such properties as multiplicity rate and mass ratio distribution.  But dynamical interactions within star clusters alone cannot sufficiently broaden an initially narrow separation distribution (Kroupa \\& Burkert \\cite{KB01}).  Additional dissipative processes and even higher stellar densities are required (Sterzik, Durisen \\& Zinnecker \\cite{SDZ04}).  Larson (\\cite{Larson02}) argues that tidal torques excited by stellar companions in an accretion disk are the most likely and efficient mechanism to extract angular momentum from accreting stars. In the absence of angular momentum transport, accreting matter accumulates in a disk or torus that fragments rapidly, creating the companion.  Indeed, hydrodynamical simulations of Bate et al. (\\cite{Bate02}) show how such companions are formed and how accreting binaries become tighter by interacting with their distant companions or with fly-by members of  a  nascent cluster.  Reipurth (\\cite{R2000}) gives convincing evidence that strong accretion and jet activity is actually observed in embedded young multiple systems.  He relates periodic knots in Herbig-Haro jets to periastron passages in inner (unresolved) binaries and interprets  the decreasing separation between the knots as a fossil record of orbit shrinkage.  However, the triple systems examined by Reipurth and Bate et al.  are relatively wide systems ($ \\sim 10$ A.U.), and cannot account for the characteristics found in short period pre-main-sequence (PMS) spectroscopic binaries (SB).  Several PMS binaries with periods shorter than 2 days are known (Melo et al. \\cite{Melo01}), and the SB fraction among PMS stars may be as high as among field stars.  Many PMS SBs are actually members of higher-order multiple systems (Sterzik et al.  \\cite{SMTB05}).  A comprehensive star formation theory must be able to explain these observational facts.  The discovery of massive extrasolar planets with predominantly short orbital periods revived the interest in orbital decay via interactions within disks. Planet migration in a disk is now a generally accepted theory.  A distant (stellar) companion enhances planet growth and migration (Zucker \\& Mazeh \\cite{ZM02}; Udry et al. \\cite{Udry03}; Eggenberger et al.  \\cite{Egg04}), possibly in an analogous way to multiple-star formation as envisioned by Larson (\\cite{Larson02}) and others.   Present-day parameters of close binaries are not identical to their parameters at birth (i.e. when their masses were assembled) because of subsequent evolution. One such evolutionary mechanism is called KCTF -- Kozai cycles with tidal friction (Eggleton \\& Kisseleva-Eggleton \\cite{KCTF}).  A distant companion in a stable (hierarchical) triple system causes periodic modulation of the inner-binary eccentricity by Kozai cycles.  These cycles are periodic, but only as long as the inner system does not interact tidally at periastron.  In this case the Kozai cycles are gradually modified, the inner system gets ``locked'' in a high-eccentricity state and its orbit then slowly decays to a circular SB with a period of few days (Kisseleva et al.  \\cite{Kisseleva98}). The period distribution of solar-type close binaries within higher-order multiple systems seems to match this scenario, showing a sharp drop in the number of systems with $P>7^d$ (Tokovinin \\& Smekhov \\cite{TS02}).  A competing evolutionary effect is the disruption of multiple systems by dynamical instabilities or by perturbations from other stars passing nearby.  Thus, SBs formed within higher order multiples can also lose tertiary companions (TCs) at later times.  Interactions with field stars can be neglected in the present study because  they disrupt only very wide tertiaries with separations above 0.1~pc (see Close et al. \\cite{Close90} for a review of observational data and theory).  Much closer tertiaries with periods as short as $\\sim$ 300~yr could be ``ionizied'' during early evolutionary stages in  very young clusters, explaining the difference in wide-binary frequency between young T~Tau associations and the field (Kroupa \\cite{Kroupa01}).   Observers have noticed the high multiplicity of close binaries for some time.  Mayor \\& Mazeh (\\cite{MM87}) looked for orbital precession (presumably caused by TCs) in a sample of 25 solar-type SBs and found that some 25\\% of those binaries are triple.  Isobe et al. (\\cite{Isobe}) searched for visual companions to SBs by means of speckle interferometry.  Tokovinin (\\cite{MSC}, MSC) showed that 43\\% of nearby solar-type stars with periods under 10 days have known tertiaries.  All five such systems in the Duquennoy \\& Mayor (\\cite{DM91}, DM91) G-dwarf sample are triple.  Are {\\it all} close binaries triple?  Is the angular momentum of the close binary system always extracted by a tertiary component (via the KCTF mechanism or during accretion)?  In the present paper, we attempt an observational  investigation of the question and prove that not all close SBs are triple.  The paper is organized as follows.  We describe our sample of close solar-type binaries in Sect.~\\ref{sec:sample}.  New components were discovered with adaptive optics (Sect.~\\ref{sec:AO}) and the 2MASS sky survey (Sect.~\\ref{sec:2MASS}).  We estimate the massses and periods of all tertiaries in Sect.~\\ref{sec:par} and describe the detection limits of new and existing techniques in Sect.~\\ref{sec:det}.    We study the statistics of tertiary components in Sect.~\\ref{sec:stat}. Section~\\ref{sec:sum} summarizes the results, and the implications for close-binary formation theories are discussed in Sect.~\\ref{sec:disc}. ",
        "conclusions": "\\label{sec:sum}   \\begin{enumerate} \\item The period distributions of SBs with and without tertiary companions (TCs) are significantly different. This proves the existence of pure SBs without TCs.  \\item The mass ratio distributions of SBs with and without TCs are identical.  \\item The frequency of TCs is 63\\% for the whole  SB sample. However, it is a strong function of the SB period, reaching 96\\% for the close ($P_1<3^d$) SBs and decreasing to 34\\% for SBs with $P_1>12^d$ (Fig.~\\ref{fig:f3p1}).  The last number is less than the companion frequency  of $\\sim$50\\% for G-dwarfs in the field.   These results are  robust because the TC detection does not depend on the SB period.  \\item The  periods of   most TCs  in  our sample  range  from 2~yr  to $10^5$~yr.  There is no correlation between $P_3$ and $P_1$.  The triple systems with $P_1<30^d$ have large period ratios  and are very stable dynamically (Fig.~\\ref{fig:p1p3}).   \\item The TCs are more massive  than spectroscopic primaries in  a small fraction (17\\% $\\pm$ 4\\% ) of multiple systems, suggesting that there is a tendency of the most massive component to be found preferentially in shortest-period sub-systems. \\end{enumerate}    The  knowledge of  the  detection  limits  of various  observing techniques is essential to  characterize the completeness of the TC statistics   (Sect.~6)   and   to   correct  the   final   results (Sect.~7.4). The  limits of each  technique are discussed  and modeled below.   \\subsection{Adaptive optics}  We studied the detection limits  for triple companions in our NACO observations of spectroscopic binaries. At each radial distance from the  spatially unresolved SB the rms fluctuation $\\sigma$ in the flux over the corresponding circle was computed, and the detection limit was assumed to be $5 \\sigma$.  This procedure has been checked by simulating artificial companions and then detecting them visually. The resulting detection limit is described by a model   \\begin{eqnarray} \\Delta K  \\leq & 7 r, \\;\\;  & \\rho  \\leq  0\\farcs1  \\nonumber \\\\ \\Delta K  \\leq & 5 r +0.9, \\;\\; & 0\\farcs1 <  \\rho   < 1\\farcs5 \\nonumber \\\\ \\Delta K  \\leq & 9.05, \\;\\;   & \\rho   \\ge 1\\farcs5 , \\label{eq:rect} \\end{eqnarray} where  $r  =  \\log   (\\rho/0\\farcs035)$.   The  diffraction  limit  at $\\lambda=2.12$~$\\mu$m  is $\\lambda/D =  0\\farcs054$.  A  similar model was  used by Shatsky  \\& Tokovinin  (\\cite{ST02}). Our  NACO detection limits are corroborated by Brandeker (\\cite{Brandeker}, his Fig.~4).  \\subsection{Visual and CPM companions}  The detectable  magnitude differences  of  known  visual companions increases  with their  separation $\\rho$.  The upper  envelope corresponds to  \\begin{eqnarray} \\Delta V & \\leq & 4.5 \\log (\\rho/0\\farcs05). \\label{eq:detvis} \\end{eqnarray} The components  lying at this  limit were all discovered  with speckle interferometry (cf.   INT4).  Thus, (\\ref{eq:detvis})  may be somewhat optimistic for 1/2  of our systems that have  never been observed with speckle techniques according  to INT4.  No new companions  are given in the   CHARM   catalog   of   lunar   occultaions   and   long-baseline interferometry  by  Richichi  et  al. (\\cite{CHARM}).   An  additional constraint  that  the apparent  magnitude  of  a  companion should  be brighter  than  $V=15^m$  is  introduced  because  fainter  stars  are generally not detected as common-proper-motion (CPM) companions.    \\subsection{Radial velocities}  A TC modulates the center-of-mass velocity  of the SB with a longer  period.  The probability of detecting  such modulation depends on  the  precision  of  radial-velocity observations  $\\sigma$,  their number $N$ and the time  span $\\Delta T$.  Studies of detection limits with simulations  (Halbwachs et al. \\cite{H03}) established that the relevant  parameter is  the radial velocity  amplitude $A_0$  of a circular-orbit binary with the same period and $90^\\circ$ inclination,  \\begin{equation} A_0 = \\frac{213 {\\cal M}_3 P_3^{-1/3} }{{\\cal M}^{2/3} }, \\label{eq:A0} \\end{equation} where $A_0$ is in km/s, tertiary mass ${\\cal M}_3$ and the SB system mass ${\\cal M}$ are in solar masses, and the tertiary's period $P_3$ is in days.  Following Halbwachs et al., we model the spectroscopic detection limit conservatively as $A_0/\\sigma >3$ for all periods shorter than the data span $\\Delta T$, and as $ A_0/\\sigma > 3 + 20 \\log (P_3/\\Delta T)$ for periods shorter than $2 \\Delta T$.   The actual detection limit may extend to periods $P_3$ much longer than $\\Delta T$ if the data are analyzed carefully, but we prefer to be conservative here by assuming the sharp loss of detection capacity at $P_3 > 2 \\Delta T$.  For  each  SB, we  determined  the  relevant  parameters $\\sigma$  and $\\Delta T$  either  from the data  given in SB9  (Pourbaix et al. \\cite{SB9}) or  from the original publications. The quality of spectroscopic orbits varies greatly, from crude orbital solutions  made some 80  years ago  (e.g. HIP  5689, 72524,  86263) to high-precision  velocities  (HIP  80686)  or long-term  coverage  (HIP 107095).   Moreover, orbits  are  often computed  by combining  radial velocities of different quality  or with different systematic offsets. Thus, any  model of  spectroscopic detection limit  is necessarily very crude and $A_0/\\sigma$ is only an indicative parameter, at best.  Despite uncertainties in the spectroscopic  detection limits, we note that half of our  sample has $\\Delta T >5.5$~yr, for 78\\%  of SBs $\\Delta T >2$~yr. The median  error  is $\\sigma =  1$~km/s. The TCs with periods  of  a few  years  would  have a good  chance of being  discovered spectroscopically,   hence the paucity of such  companions in our sample must be real.   \\subsection{Astrometry}  Astrometry is a poweful technique to detect low-mass companions with orbital periods over several years.  Makarov \\& Kaplan (\\cite{MK05}) recently published a catalog of astrometric binaries revealed either by a significant difference between ``instantaneous'' proper motion (PM) measured by Hipparcos and long-term PM in the Tycho-2 catalog ($\\Delta \\mu$ binaries) or by the acceleration measured directly by Hipparcos ($\\dot{\\mu}$ binaries, also known as G-solutions in Hipparcos). There are 11 astrometric binaries in common with our sample, of which six are already listed here with TCs discovered by other techniques (including HIP 19248 and 64219 resolved for the first time here).  We did not resolve two $\\dot{\\mu}$ binaries, HIP 35487 and 114639, with NACO, possibly because their astrometric companions are white dwarfs (cf. discussion in Notes to Table~5).  Gontcharov  et  al.   (\\cite{G01})  published a  similar  study:  they brought historical astrometric catalogs  into the Hipparcos system and identified several stars as astrometric binaries.  However, only a few astrometric orbits of these stars  are published to date (none for our sample),  and Makarov  \\& Kaplan  (\\cite{MK05}) did  not independently confirm the PM variations of many stars. We observed with NACO but did not  resolve HIP  67153, 80686,  85365,  86263 listed  as binaries  by Gontcharov  et al.   Here we  consider astrometric  detections without published orbits only  as hints of the existence of  TCs.  Hence we do not take into account astrometric companions in our statistics."
    },
    "0601/astro-ph0601296_arXiv.txt": {
        "abstract": "Photoelectric emission from dust plays an important role in grain charging and gas heating.  To date, detailed models of these processes have focused primarily on grains exposed to soft radiation fields. We provide new estimates of the photoelectric yield for neutral and charged carbonaceous and silicate grains, for photon energies exceeding $20 \\eV$.  We include the ejection of electrons from both the band structure of the material and the inner shells of the constituent atoms, as well as Auger and secondary electron emission.  We apply the model to estimate gas heating rates in planetary nebulae and grain charges in the outflows of broad absorption line quasars. For these applications, secondary emission can be neglected; the combined effect of inner shell and Auger emission is small, though not always negligible. Finally, we investigate the survivability of dust entrained in quasar outflows.  The lack of nuclear reddening in broad absorption line quasars may be explained by sputtering of grains in the outflows. ",
        "introduction": "Introduction}  Cosmic grains are subjected to electrical charging processes and the resulting non-zero charge can have important astrophysical consequences. In Galactic \\ion{H}{1} regions, the grain charge is primarily determined by a balance between starlight-induced photoelectric emission and the collisional capture of electrons from the gas.  Furthermore, in these regions the gas heating is dominated by photoelectric emission from dust. Thus, there have been numerous detailed studies of UV-induced photoelectric emission (Spitzer 1948; Watson 1972; de Jong 1977; Draine 1978; Tielens \\& Hollenbach 1985; Bakes \\& Tielens 1994; Weingartner \\& Draine 2001b).  Photoelectric emission from grains exposed to extreme ultraviolet (EUV) and X-ray radiation has received much less attention.  Dwek \\& Smith (1996) modelled, in detail, the dust and gas heating for neutral grains exposed to high-energy radiation, but did not address the grain charging.  Charging by high-energy photons can be very important for grains immersed in hot plasma and exposed to a hard radiation field.  The grain potential is limited by the highest photon energy present in the incident radiation; as this increases, the grain can reach higher potentials.  The efficiency of the gas heating decreases as the grain charge increases, since the photoelectrons have to climb out of the potential well.  Hard radiation fields can be found, e.g., in planetary nebulae and active galactic nuclei (AGNs).  If the gas contains dust, then a fraction of the incident radiation will be processed into infrared (IR) radiation, i.e., thermal dust emission.  Dwek \\& Smith (1996) were motivated by the prospect of learning about these environments through the analysis of the IR emission.  Ferland et al.~(2002) stressed that the X-ray and IR spectra in AGNs may be correlated as a result of the interaction between hot grains and gas.  Detailed modelling of these regions will only be possible once the grain charging is properly modelled, since the grain and gas heating rates depend on the charging.  In a previous paper (Weingartner \\& Draine 2001b, hereafter WD01), we modelled in detail the photoelectric emission from grains exposed to UV radiation; this is briefly summarized in \\S \\ref{sec:summ}. In \\S \\S \\ref{sec:pe} through \\ref{sec:htg}, we extend the WD01 model to include EUV and X-ray radiation. Applications to planetary nebulae and quasar outflows are described in \\S \\ref{sec:PN} and \\S \\ref{sec:quasar}, respectively. ",
        "conclusions": "In this paper, we have extended the WD01 photoelectric emission model to higher photon energies, enabling the treatment of dust exposed to hard radiation fields.  We provide revised yields for photoelectric emission from the band structure when $h\\nu > 20 \\eV$ and treat the emission of Auger electrons and secondary electrons excited by primary and Auger electrons.  For H$\\,$II regions, ionized by OB stars, the radiation field is soft enough that the WD01 model can be employed without modification; the new treatment yields  a $\\ltsim 10\\%$ difference in the gas heating rate.  For planetary nebulae, the gas heating rates found with the extended model are lower, by a factor 2 or less, than those estimated with the WD01 model, since the revised yields for photoelectric emission from the band structure are lower. Inner shell, Auger, and secondary emission are not important. For quasar outflows, the new treatment yields significantly lower grain potentials.  In this case, combined inner shell and Auger emission contribute at as high as the 30\\% level.  This is less than the uncertainty associated with photoelectric emission from only the band structure.  Secondary emission is unimportant.  It appears that thermal sputtering can destroy dust in BALQSO outflows in less than a dynamical time, for a wide range of plausible physical conditions in the outflow.  However, better constraints on the outflow conditions are needed to definitively settle the issue of dust survivability.  In environments with harder radiation fields (e.g. supernova remnants and the central regions of galaxy clusters), inner shell, Auger, and secondary emission may be more important relative to photoelectric emission from the band structure.  However, in these cases the gas is so hot that secondary emission induced by gas-phase electrons and ions incident on the grain is expected to dominate over photoelectric emission in grain charging (Draine \\& Salpeter 1979)."
    },
    "0601/astro-ph0601069_arXiv.txt": {
        "abstract": "Emission lines formed in the circumstellar envelopes of several type of stars can be modeled using first principles of line formation. We present simple ways of calculating line emission profiles formed in circumstellar envelopes having different geometrical configurations. The fit of the observed line profiles with the calculated ones may give first order estimates of the physical parameters characterizing the line formation regions: opacity, size, particle density distribution, velocity fields, excitation temperature. ",
        "introduction": "In optically thick regions, more than 90\\% of the emitted energy in a given line can be considered produced in a limited region of the circumstellar disc. This is the case of Fe\\,{\\sc ii} in Be stars (Arias et al. 2005), but also hydrogen Balmer, Paschen series if their formation region has the aspect of an expanding ring. In both circumstances, the emitted radiation can be estimated using an equivalent isothermal ring with uniform semi-height $H$, whose surface density equals the radial column density. If the particle density has a distribution $N(R)\\!\\sim$ $R^{-\\beta}$, the radius of the ring is then $R_r/R_*\\!=\\!$ $[(1-\\beta)/(2-\\beta)][1-(R_i/R_e)^{\\beta -2}]/[1-(R_i/R_e)^{\\beta-1}]$, where $R_{i,e}$ are the internal and external radii of the line forming region. The ring can be considered having expansion $V_{exp}$ and rotation $V_{\\Omega}$ velocities, both represent averages of the respective velocity fields (Arias 2004). The emitted flux is simply $F_{\\lambda}=$ $\\int_{{\\cal S}(i)}I_{\\lambda}(x,y){\\rm d}x{\\rm d}y$ where ${\\cal S}(i)$ is the aspect-angle effective emitting surface projected on the sky and $I_{\\lambda}(x,y)$~is: \\begin{equation} \\left.\\begin{array}{lcl} I^a(x,y,v-v_r) & = & I^*(x,y)e^{-{\\tau^f(v-v_r)\\over\\mu(x,y)}}+ S[1-e^{-{\\tau^f(v-v_r)\\over\\mu(x,y)}}] \\\\ &  & ({\\rm stellar\\ emission\\ absorbed\\ by\\ the\\ front\\ side}) +\\cr &  & ({\\rm front\\ side\\ emission})\\\\ I^b(x,y,v-v_r) & = & S[1-e^{-{\\tau^t(v-v_r)\\over\\mu(x,y) }}]e^{-{\\tau^f(v-v_r)\\over\\mu(x,y)}}+ S[1-e^{-{\\tau^f(v-v_r)\\over\\mu(x,y)}}] \\\\ &  & ({\\rm rear\\ emission\\ absorbed\\ by\\ the\\ front\\ side}) +\\\\ &  & ({\\rm front\\ side\\ emission}) \\\\ \\end{array} \\right\\} \\end{equation} \\noindent with $\\tau_v\\!=$ $\\tau_o\\Phi(v-v_r)$ and $v_r(x,y)\\!=$ $\\pm \\{V_{exp}(R)[1-({x\\over R})^2]^{1\\over2}\\pm V_{\\Omega}(R){x\\over R}\\}\\sin i$, where $(f,r)$ stand for `front' and `rear' sides of the ring. The signs in the radial velocity $v_r$ are chosen according to the quadrant; $v$ is the Doppler displacement in the observed emission line profile, $\\mu(x,y)\\!=$ $\\cos$(ring-normal;observer), $i$ inclination of the ring-star system, $\\tau_o$ is the radial opacity of the ring in the line center and $\\Phi(v)$ is the intrinsic absorption line profile. The line source function is: $S_{\\lambda}(\\tau_{\\rm o})\\!=$ $S_o$ for $\\tau_o\\leq1$; $S_o\\tau_{\\rm o}^p$ for $\\tau_o\\!>\\!1$ [$p\\!\\simeq\\!1/2$ for Gaussian $\\Phi(v)$] and $B_{\\lambda_o}(T_{\\rm e})$ for $\\tau_o\\!\\geq$ $(B_{\\lambda_o}/S_o)^2$ (Mihalas 1978, Chap.11). $S_o$ depends on the nature of the line transition: collision-, photoionization-, mixed-dominated. Using $S_o$ as done in Cidale \\& Ringuelet (1989) we can determine the excitation temperature (Arias et al. 2005). Thus, the free parameters to fit the observed emission line profiles are: $S_o/F_*$, $\\tau_o$, $R_r$, $H$, $i$, $V_{\\Omega}$, $V_{exp}$ and density distribution $\\beta$ (or $\\alpha\\!=$ $2-\\beta$). The shape of the line profiles reduce strongly the space of free parameters, mainly those of the velocity field and the ratio between $\\tau_o$ and $H$. Line intensities are sensitive to $(S_o/F_*)R^2_r$ and $R_r$ to the temperature, which in turn determines $S_o/F_*$. Figure~\\ref{f1}a shows line profiles obtained for the indicated parameters of the ring with $V_{exp}\\!=\\!0$, which are typical for {\\it `shell'} lines. The same ring parameters are used for Fig.~\\ref{f1}b, but $V_{exp}\\!=$ 150 km/s. The line profiles are of the {\\it`steeple'} type seen frequently in Fe\\,{\\sc ii} lines. Flaring discs can be treated in the same way as cylindrical ones, except that the surface density, and hence $\\tau_o$ depends on the coordinate perpendicular to the equator (Vinicius et al. 2005). A fit of the H$\\alpha$ line of $\\alpha$~Eri in 1994 with a flaring disc of opening angle $\\phi=15^o$ is shown in Fig.~\\ref{f1}c, where are also shown line profiles for different $\\beta$-values. \\begin{figure}[t] \\centerline{\\psfig{file=marias_fig1.eps,width=12.5truecm,height=4.3truecm}} \\caption[]{a) Line profiles for a ring with $V_{exp}\\!=\\!0$ and the indicated parameters. b) 'Steeple' line profiles produced by rings with same parameters as in a) but $V_{\\exp}\\!\\neq\\!0$. c) Line profiles due to flaring discs treated as rings, for several $\\beta$ and opening angle $\\phi\\!=\\!15^o$. $\\beta\\!=$ 0.1 closely fits H$\\alpha$ of $\\alpha$~Eri in 1994 (a flat ring would require $H\\!=\\!3.8$; $\\beta\\!=\\!-0.1$)} \\label{f1} \\end{figure} \\begin{figure}[] \\centerline{\\psfig{file=marias_fig2.eps,width=12.5truecm,height=4truecm}} \\caption[]{Three-peak line profiles produced by rotating+expanding rings} \\label{f2} \\end{figure} Several three-peak emission line profile are shown in Fig.~\\ref{f2} produced in regions treated as equivalent expanding rings. Due to the Doppler shifts, the central emission components are produced by the front and rear sides of ring sectors where the radial velocity is $v_r\\simeq0$. ",
        "conclusions": ""
    },
    "0601/astro-ph0601543_arXiv.txt": {
        "abstract": "We present new optical emission-line images of the young supernova remnant (SNR) 1E~0102--7219  (hereafter E0102)  in the Small Magellanic Cloud obtained with the {\\it Hubble Space Telescope} ({\\it HST}) Advanced Camera for Surveys (ACS).  E0102 is  a member  of the  oxygen-rich class of SNRs showing strong oxygen, neon, and other metal-line emissions in its optical and X-ray spectra, and an absence of hydrogen and helium. The progenitor of E0102 may have been a Wolf-Rayet star that underwent considerable mass loss prior to exploding as a Type Ib/c or IIL/b supernova.  The ejecta in this SNR are generally fast-moving (V $>$ 1000 \\kms) and emit as they are compressed and heated in the reverse shock. In 2003, we obtained optical [O~III], H$\\alpha$, and continuum images with the ACS Wide Field Camera. The [O~III] image through the F475W filter captures the full velocity range of the ejecta, and shows considerable high-velocity emission projected in the middle of the SNR that was Doppler-shifted out of the narrow F502N bandpass of a previous Wide Field and Planetary Camera 2 image from 1995.  Using these two  epochs separated by $\\sim$8.5  years, we measure the transverse expansion of the ejecta around the outer rim in this  SNR  for the  first time  at  visible wavelengths.  From proper-motion measurements of 12  ejecta filaments, we estimate a mean expansion velocity for the bright ejecta of $\\sim$2000 \\kms\\ and an inferred kinematic age for the SNR of $\\sim2050 \\pm 600$ years.  The age we derive from {\\it HST} data is about twice that inferred by Hughes et al.\\ (2000, ApJ, 543, L61) from X-ray data, though our 1-$\\sigma$ error bars overlap. Our proper-motion age is consistent with an independent optical kinematic age derived by Eriksen et al.\\ (2003, AIPC, 565, 419) using spatially resolved [O~III] radial-velocity data.  We derive an expansion center that lies very close to conspicuous X-ray and radio hotspots, which could indicate the presence of a compact remnant (neutron star or black hole). ",
        "introduction": "\\label{introduction}  Supernova remnants (SNRs) reveal how stars synthesize  and distribute the elements heavier than H and He.  Young SNRs with uncontaminated debris are among our best natural laboratories for understanding how nucleosynthesis operates in massive stars.  Studying the SNR ejecta provides clues about the dynamics of supernova (SN) explosions, the composition and spectral type of progenitors, and ultimately about the chemical enrichment process in galaxies.  The oxygen-rich SNR 1E~0102--7219 (E0102) in the Small Magellanic Cloud (SMC) offers an exceptional opportunity for multi-wavelength, high resolution studies of a young SNR. E0102 is a vital object in the small class of oxygen-rich SNRs that contain filaments of uncontaminated debris from core-collapse supernovae.  Out of the few ($\\sim$10) known objects in the oxygen-rich class, E0102 is one of the most interesting because of its similarities to the prototype O-rich SNR, Cassiopeia~A (Cas~A) in the Galaxy.  However, unlike Cas~A, the sight-line to E0102 in the SMC has low interstellar extinction that allows us to gather data from the X-rays to the radio; yet it is close enough that we can still discern physically important scales.  E0102 was discovered by \\cite{seward81} in an {\\it Einstein} X-ray observatory survey of the Magellanic Clouds. It is located in the SMC just northeast of the high-excitation H~II region N76A (\\citealt{hen56}). Dopita, Tuohy, \\& Mathewson (1981) recognized E0102 as an O-rich SNR by virtue of its strong optical [O~III]$\\lambda$$\\lambda$4959,5007 emission but complete lack of hydrogen emission from a network of filaments. Subsequent spectra by \\cite{td83} confirmed the presence of fast-moving, highly processed ejecta remarkable for showing emission from O and Ne only. They also estimated a kinematic age for the SNR of $\\sim$1000 years based on the velocity range they observed and its angular size.  E0102 exhibits the textbook forward/reverse shock structure that forms during the evolution of SNRs (\\citealt{chev82}). Emissions from the faint $\\sim$44$^{\\prime\\prime}$ diameter blast wave and bright $\\sim$30$^{\\prime\\prime}$ diameter shell of shocked ejecta have now been observed in exquisite detail across the electromagnetic spectrum from the ground and space (e.g., \\citealt{amyball,blair89,hay94,hughes00,gaetz00,stan05}). The combined X-ray, UV, and optical emission spectra of the ejecta contains lines from O, Ne, C, Si, and Mg over a wide range of ionization states, from neutral to H-like species (\\citealt{blair00,ras}). The comprehensive study by \\cite{flan04} of the high-resolution {\\it Chandra} X-ray spectrum shows that the highly ionized species are stratified; the higher ionization states occur at larger radii as the reverse shock propagates into the expanding ejecta. These authors also estimated a progenitor zero-age-main-sequence (ZAMS) mass of $\\sim$32M\\sol\\ based on their derived oxygen mass of 6M\\sol\\ and heavy element abundance ratios, and comparing to theoretical predictions.  In this paper we present new optical images from the Advanced Camera for Surveys (ACS) aboard the {\\it Hubble Space Telescope} ({\\it HST}). We describe the observations in Sec.\\ 2 and present the imaging results in Sec.\\ 3. The new images capture the full velocity range of the emitting ejecta and reveal fresh details on the fastest filaments projected against the interior of the remnant.  We compare our new [O~III] image to one obtained previously with the Wide Field and Planetary Camera 2 (WFPC2) and measure proper motions of several outer filaments captured in both data sets. In Sec.\\ 4 we discuss some implications of our proper-motion measurements and examine additional clues about the possible progenitor of E0102. Our conclusions are presented in Sec.\\ 5. We assume a distance to the SMC of 59 kpc (e.g., van den Bergh 1999), where $0.1'' \\approx 9 \\times 10^{16}$ cm. ",
        "conclusions": "\\label{conc}  A new imaging data set from the {\\it Hubble Space Telescope} Advanced Camera for Surveys reveals spectacular details in the O-rich SNR 1E~0102--7219 in the SMC. We used the broad-band SDSS $g$ (F475W) filter to capture the full [O~III]$\\lambda\\lambda4959,5007$ velocity range of the ejecta. Many high-velocity knots and filaments projected through the middle of the SNR are seen for the first time at $\\sim0.1''$ resolution. We discuss their morphology, kinematics, and relationship to soft X-ray emission observed in a {\\it Chandra} ACIS-S image. The wide field of view of the ACS images also allows us to examine the circumstellar environment of E0102. A narrowband H$\\alpha$+[N~II] (F658N) image shows considerable inhomogeneity in the surrounding medium and a possible physical association between some of the gas in the halo of E0102 and the nearby N76A nebula.  We compared the 2003 ACS [O~III] image to a 1995 narrowband WFPC2 Planetary Camera [O~III] image and measured tangential expansion velocities in 12 regions around the periphery of the SNR. We find a mean expansion velocity for the optical ejecta of $\\sim 2000$ \\kms\\ and a kinematic age for E0102 of $\\sim2050 \\pm 580$ years. Our kinematic age agrees with a recent estimate of \\cite{erik01} based on spatially resolved [O~III] radial-velocity data, but is a factor of two longer than found by \\cite{hughes00} using multi-epoch, low-resolution X-ray images (although it is within 2$\\sigma$).  The ejecta proper motions we measured trace back to an expansion center $\\sim2.4''$ south of the geometric center of the {\\it Chandra} X-ray image. The 1$\\sigma$ error circle of our expansion center encompasses the conspicuous radio and X-ray blobs noted in previous studies (e.g., \\citealt{amyball,gaetz00,flan04}). Our kinematic data provide additional evidence for believing that a compact remnant (neutron star or black hole) may be associated with the central hotspot emissions. E0102's proximity to the N76A Wolf-Rayet nebula and Hodge 53 OB association leads us to speculate that the E0102 progenitor may have been a very massive ($>50$ M$_{\\odot}$) main-sequence star that underwent considerable mass loss as it evolved through a WN phase before exploding as a Type Ib/c supernova.  We thank the referee for carefully reviewing the manuscript and recommending changes that significantly improved the fidelity of the results. This  work  has  been  supported  by  NASA  grant  NAG5-12279  to  the University of Colorado and the COS Science Team."
    },
    "0601/astro-ph0601405_arXiv.txt": {
        "abstract": "Cepheus B (Cep B) molecular cloud and a portion of the nearby Cep OB3b OB association, one of the most active regions of star formation within 1 kpc, has been observed with the ACIS detector on board the {\\it Chandra X-ray Observatory}. We detect 431 X-ray sources, of which $89\\%$ are confidently identified as clustered pre-main sequence (PMS) stars. Two main results are obtained. First, we provide the best census to date for the stellar population of the region.  We identify many members of two rich stellar clusters: the lightly obscured Cep OB3b association, and the deeply embedded cluster in Cep B whose existence was previously traced only by a handful of radio sources and T Tauri stars. Second, we find a discrepancy between the X-ray Luminosity Functions (XLFs) of the Cep OB3b and the Orion Nebula Cluster (ONC). This may be due to different Initial Mass Functions (IMFs) of two regions (excess of $\\simeq 0.3$ M$_\\odot$ stars), or different age distributions.  Several other results are obtained. A diffuse X-ray component seen in the field is attributed to the integrated emission of unresolved low mass PMS stars. The X-ray emission from HD 217086 (O7n), the principle ionizing source of the region, follows the standard model involving many small shocks in an unmagnetized radiatively accelerated wind. The X-ray source \\#294 joins a number of similar superflare PMS stars where long magnetic structures may connect the protoplanetary disk to the stellar surface. ",
        "introduction": "Cepheus B (Cep B), the hottest CO component of the Cepheus molecular cloud \\citep{Sargent77,Sargent79}, is located near the northwestern edge of the cloud at a distance of 725 pc from the Sun \\citep{Blaauw59,Crawford70}. Outside Cep B, to the north and west of the cloud, lies the stellar Cep OB3 association, composed of two subgroups of stars of different ages with the youngest (Cep OB3b) lying closer to the molecular cloud \\citep{Blaauw64,Jordi96} (see Figure \\ref{fig_intro_fig}).  \\citet{Sargent79} considered Cep OB3  to be a good example of sequential star formation, following the model of \\citet{Elmegreen77}. The interface between the molecular cloud and the OB star association is clearly delineated by the optically visible H{\\sc II} region Sharpless 155 (S 155), where cloud material is ionized and heated by the radiation field of the youngest generation of the Cep OB association \\citep{Panagia81}. The photodissociation region (PDR) at S 155 is favorably oriented to reveal the progression of star formation. As the surface of the cloud is being eroded by the early-type stars, the cloud edge moves eastward across the observer's field of view with new stars emerging from the obscuring molecular cloud. Radio continuum \\citep{Felli78, Testi95}, far-infrared, CO \\citep{Minchin92}, and near-infrared studies \\citep{Moreno-Corral93, Testi95} reveal a few young stars embedded in the Cep B cloud behind the PDR. Sources located on the edge of the molecular cloud and exhibiting high extinction were suggested to represent a third generation of star formation triggered by the H{\\sc II} region shock propagating into the cloud.  The low mass pre-main sequence (PMS) stellar population of Cep OB3b can be effectively discriminated using X-ray emission due to the heavy contamination in optical and near-infrared (ONIR) surveys by both foreground and background Galactic field stars. X-ray surveys are complementary to ONIR surveys because they trace magnetic activity (mainly plasma heated in violent magnetic reconnection flares) rather than photospheric or circumstellar disk blackbody emission, and PMS X-ray emission is elevated $10^{1}-10^{4}$ above main sequence (MS) levels \\citep{Preibisch05}. A {\\it ROSAT} X-ray study of Cep OB3b identified over 50 likely PMS stars with $A_V < 3$ mag and $L_X \\ga 10^{30}$ ergs s$^{-1}$ \\citep{Naylor99}.  Ten of these {\\it ROSAT} sources are coincident with optical sources, and four of these are spectroscopically confirmed classical and weak-lined T Tauri stars \\citep{Pozzo03}.  Much more can be achieved with the current generation of X-ray telescopes, particularly the {\\it Chandra X-ray Observatory} with its excellent high-resolution mirrors and low-noise detectors. Studies of the Orion Nebula region demonstrate that {\\it Chandra} images can penetrate up to $A_V \\simeq 500$ mag into the cloud \\citep{Getman05, Grosso05}, considerably deeper than near-infrared surveys in the $JHKL$ bands. $Chandra$ is also effective in resolving crowded fields down to $\\simeq 0.7$\\arcsec\\/ scales \\citep{Getman05b}.  In X-rays, OB stars are often not much brighter than PMS stars, so multiple systems associated with OB stars can be identified \\citep{Stelzer05}.  $Chandra$ X-ray studies are particularly effective in uncovering heavily obscured low-mass cloud populations and in discriminating PMS populations from unrelated older stars.  We present here a {\\it Chandra} study of the region using its Advanced CCD Imaging Spectrometer (ACIS) detector.  Over 380 PMS stars in two rich clusters are found, concentrated in the lightly obscured Cep OB3b cluster and a previously unknown embedded cluster, and we infer the existence of an additional $700-800$ low-mass members. Only $\\simeq 7-9$\\% of the detected X-ray sources are attributable to contaminating Galactic and extragalactic populations.  The observation and source lists are described in \\S \\ref{data_reduction_section}. Identification with 2MASS counterparts and their NIR properties are considered in \\S \\ref{counterparts_and_properties_section}. X-ray fluxes and their errors are determined in \\S \\ref{flux_det_section}. X-ray sources unrelated to the cloud are discussed in \\S \\ref{unidentified_main_section}. The discovery of two rich clusters is described in \\S \\ref{new_cluster_section}.  Their stellar populations are discussed in \\S \\ref{imf_xlf_section} using measured X-ray luminosity functions (XLFs) and Initial Mass Functions (IMFs). The origin of the soft diffuse X-ray emission discovered here is discussed in \\S \\ref{diffuse_emission_section}. Finally, we address two individual stars of particular interest: HD 217086, the most massive star in Cep OB3b (\\S \\ref{OB_section}), and {\\it Chandra} source \\#294 which exhibited an unusually luminous and hot X-ray flare (\\S \\ref{interesting_sources_section}). ",
        "conclusions": "\\label{summary_section}  We present a 30 ks (net exposure) $Chandra$ ACIS observation of the Cepheus B molecular cloud and a portion of the nearby Cep OB3b OB association. The main conclusions of our study are as follows:  1. We detect 431 X-ray sources with a limiting luminosity for the region outside of the molecular cloud of $\\sim 10^{29}$ ergs/s. Three hundred eighty four of them are unambiguously identified with 2MASS NIR sources. Upon visual inspection of 2MASS atlases we find NIR counterparts for at least 10 (2) more soft (hard) X-ray sources, which have not been registered in the 2MASS catalog. On the NIR color-magnitude diagram most of the detected X-ray sources occupy the locus of low-mass PMS stars with a typical absorption of $A_V \\sim 2.5$ mag (assuming age of $\\sim 1$ Myr), but as many as 60 sources with NIR counterparts, residing close to/in the molecular cloud, have even much higher absorptions up to $A_V \\sim 28$ mag.  2. Twenty four unidentified hard X-ray sources are classified as possible extragalactic contaminants (Table \\ref{tbl_spe}, `a'). Though some of them still may be embedded members of the cloud, about 70\\% percent of them are probably true AGNs, as argued by a detailed simulation of the extragalactic background population, including instrumental background and local (Cepheus) plus large scale Galactic absorptions. Through careful simulations of X-ray emission from the Galactic foreground and background populations, employing a stellar population synthesis model of the Galactic disk, and through the inspection of NIR colors of X-ray detected sources we expect to have no more than a few Galactic background stars and we suggest 13 field star candidates (Table \\ref{tbl_spe}, `f').  3. In order to obtain a reliable X-ray luminosity function of the detected PMS stars we determine source X-ray fluxes based on knowledge of their NIR absorptions and X-ray median energies. We then perform careful MARX simulations to derive errors on median energy, combine them with known errors on count rate and absorption, and propagate them to the flux errors through Monte-Carlo simulations.  4. We discover more than 380 PMS stars in the ACIS-I field of the Cepheus region, which provide the best census to date for the stellar population of the cloud. Based on $N_H$ and $K$-band stratified spatial distributions of the detected PMS stars we separate them into two rich clusters: the lightly obscured cluster (= ``Cep OB3b within ACIS-I field'') (321 stars, Table \\ref{tbl_spe}, `l') and the embedded cluster (64 stars, Table \\ref{tbl_spe}, `e').  We estimate that another $400-450$ ($300-340$) low mass PMS stars are present in the unobscured (obscured) clusters within the $Chandra$ field but are too X-ray-faint to be individually identified.  5. We find that the XLF of the lightly-obscured Cep OB3b cluster region within the ACIS-I field shows an excess of stars around $\\log L_X \\sim 29.7$ erg s$^{-1}$ ($0.5-8$ keV, after individual correction for absorption) compared to the well-studied Orion Nebula Cluster. This may be caused by a nonstandard IMF in Cep OB3b with an excess of $\\simeq 0.3$ M$_\\odot$ stars if the cluster age is $\\sim 1$ Myr. This excess also appears in the independently derived KLF based on 2MASS source counts (after an uncertain subtraction of background counts from two control fields) which show more $K \\sim 13-14$ stars than expected from a standard IMF. However, the discrepancy may also be attributed to different ages and star formation histories for the Cep OB3b and ONC clusters.  6. We show that the HD 217086 (O7n) star with its known wind terminal velocity and mass-loss rate is unable to produce the bright soft X-ray diffuse emission seen in the field. An alternative explanation --- the X-ray diffuse emission originates from the unresolved PMS point sources in the region --- is suggested by (1) the spatial coincidence of the diffuse emission with the detected point source population, (2) the similarity of the plasma temperature of the diffuse emission with that of the combined weak PMS Orion-COUP sources, (3) the equivalence of the amount of the observed X-ray luminosity for the diffuse emission with that of the estimated undetected members of the unobscured population.  7. We establish that the X-ray emission from the HD 217086 (O7n) star is a constant soft-spectrum emission at the level $\\log (L_{X}/L_{bol}) \\sim -7$ that is expected from the standard model involving many small shocks in an un-magnetized radiatively accelerated wind.  8. From the evolution of an adaptively smoothed median energy during the unusually hot and strong flare in source \\# 294 we infer the slope of the trajectory in the temperature-density diagram, estimate the lower limit on the size of the flaring structure, and suggest that the source \\# 294 may join a number of similar super-flare sources seen in Orion clouds where long magnetic structures may connect the protoplanetory disk to the stellar surface.  We thank the referee, Tim Naylor, for his time and many useful comments that improved this work. This work was supported by the ACIS Team (G. Garmire, PI) through NASA contract NAS8-38252 and from {\\it Chandra} contract SV4-74018 issued by the {\\it Chandra} X-ray Center operated by SAO under NASA contract NAS8-03060. M.T. acknowledges financial support from Japan Society for the Promotion of Science. 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."
    },
    "0601/astro-ph0601680_arXiv.txt": {
        "abstract": "The bulk of the molecular component in galaxies is made of cold H$_2$, which is not observed directly, but which abundance is derived from indirect tracers such as CO emission. The CO to H$_2$ conversion ratio remains uncertain, and may vary by large factors in special environments with different excitation or metallicity. Recent cold gas discoveries (through $\\gamma$-rays or cold dust emission) are reviewed and the most promising tracers in the future are discussed, such as the primordial molecules HD and LiH, or the pure rotational lines of excited H$_2^*$. ",
        "introduction": "The H$_2$ molecule has no electric dipole because of its symmetry, and therefore cannot emit line radiation in the radio domain, the only ones that could be excited at the low temperature of the interstellar medium (10-15K). The first rotational line comes from the $J=2$ level, through quadrupole radiation, and is 512K above the ground state, S(0) at 28$\\mu$m in wavelength. The second rotational line of the para H$_2$ is S(2) at 12$\\mu$m, and the two first lines of the ortho H$_2$ are S(1) and S(3) at 17$\\mu$m and 9$\\mu$m respectively. ISO observations of PDR have revealed a large number of H$_2$ pure rotational lines (Timmermann et al 1996), and higher resolution observations have confirmed an unexpectedly large amounts of warm H$_2$ gas in PDR (Allers et al 2005).  The only other way to detect cold H$_2$ directly is through the UV lines in absorption, unfortunately this concerns only the untypical very diffuse regions, since higher H$_2$ column densities are associated to high extinction, that prevent to detect the background sources. This was confirmed by the FUSE survey towards galactic sources (Shull et al 2004). The average molecular fraction of the diffuse ISM is of the order of 10\\%, and the temperature is warm ($\\sim$ 100K). Even in \"translucent line of sights\", there is only evidence of diffuse clouds (Rachford et al 2002).  We first review our current knowledge of the CO to H$_2$ conversion ratio, which is still the main method to trace cold H$_2$ molecules in galaxies. After presenting possible new tracers, we then describe recent discoveries of cold gas, locally in our Galaxy, and in external galaxies or clusters. We conclude with future experiments, which could bring new insights. ",
        "conclusions": ""
    },
    "0601/astro-ph0601363_arXiv.txt": {
        "abstract": "We present an analysis of interstellar Na\\,{\\sc i} ($\\lambda_{\\rm air}$=3302.37\\AA, 3302.98\\AA), Ti\\,{\\sc ii} ($\\lambda_{\\rm air}$=3383.76\\AA) and Ca\\,{\\sc ii} K ($\\lambda_{\\rm air}$=3933.66\\AA) absorption features for 74 sightlines towards O- and B-type stars in the Galactic disc. The data were obtained from the UVES Paranal Observatory Project, at a spectral resolution of 3.75~km~s$^{-1}$ and with mean signal to noise ratios per pixel of 260, 300 and 430 for the Na\\,{\\sc i}, Ti\\,{\\sc ii} and Ca\\,{\\sc ii} observations, respectively. Interstellar features were detected in all but one of the Ti\\,{\\sc ii} sightlines and all of the Ca\\,{\\sc ii} sightlines. The dependence of the column density of these three species with distance, height relative to the Galactic plane, H\\,{\\sc i} column density, reddening and depletion relative to the solar abundance has been investigated. We also examine the accuracy of using the Na\\,{\\sc i} column density as an indicator of that for H\\,{\\sc i}. In general we find similar strong correlations for both Ti and Ca, and weaker correlations for Na. Our results confirm the general belief that Ti and Ca occur in the same regions of the interstellar medium and also that the Ti\\,{\\sc ii}/Ca\\,{\\sc ii} ratio is constant over all parameters. We hence conclude that the absorption properties of Ti and Ca are essentially constant under the general interstellar medium conditions of the Galactic disc. ",
        "introduction": "\\label{s_intro}  The study of the interstellar medium in our Galaxy has a long history, dating back a century when interstellar Ca\\,{\\sc ii} absorption was first detected (Hartmann 1904). Historically, ground-based work has concentrated on the Ca\\,{\\sc ii} K and Na\\,{\\sc i} D lines at around 3933~\\AA\\ and 5980~\\AA. Studies such as Albert et al. (1993), Sembach, Danks \\& Savage (1993), Welty, Hobbs \\& Kulkarni (1994), Welty, Morton \\& Hobbs (1996) and Smoker et al. (2003) have observed the Ca\\,{\\sc ii} K and Na\\,{\\sc i} D lines in an attempt to provide information on the structure and physical characteristics of interstellar clouds, in particular their spatial distribution (e.g. Lallement et al. 2003), kinematics (e.g. Sembach \\& Danks 1994) and depletion patterns of these elements (e.g. Crinklaw, Federman \\& Joseph 1994; Wakker \\& Mathis 2000). The general situation can be summarised as follows; Ca\\,{\\sc ii}, with its ionisation potential of $\\sim$ 11.9 eV, is typically a trace ionisation stage in diffuse, neutral clouds, plus a more extended intercloud medium. However Na\\,{\\sc i} is found to trace the neutral clouds only. Ca\\,{\\sc ii} is thought to be heavily depleted onto dust in denser clouds, with Na\\,{\\sc i} a little-depleted element (Welty et al. 1996).  Other interstellar lines which have been studied much less frequently are the UV Na\\,{\\sc i} doublet at $\\lambda_{\\rm air}$=3302.37\\AA, 3302.98\\AA, and Ti\\,{\\sc ii} at $\\lambda_{\\rm air}$=3383.76\\AA. The relative paucity of observations is for two main reasons. First, the oscillator strengths of $f$=0.0092, 0.0046 and 0.358 for the Na\\,{\\sc i} doublet and Ti\\,{\\sc ii}, respectively, are somewhat smaller than for the Ca\\,{\\sc ii} K ($f$=0.627) and Na\\,{\\sc i} D lines ($f$=0.318, 0.631). Second, both the Na\\,{\\sc i} UV doublet and Ti\\,{\\sc ii} lines lie toward the blue end of the optical range, where detectors are less efficient. Although the weakness of the lines clearly makes them more difficult to observe, it has the advantage that (especially in the case of Na\\,{\\sc i} UV doublet) saturation effects are much less of an issue. Hence these transitions can be used to probe quite dense regions of the ISM, without having to apply large corrections for the effects of line saturation. In the case of Ti\\,{\\sc ii}, the ionisation potential of 13.57 eV is nearly the same as that of H\\,{\\sc i}, so that this species is the dominant ionisation stage in H\\,{\\sc i} regions and no ionisation corrections are necessary to obtain the total abundance of Ti in these regions. This is in contrast to Ca\\,{\\sc ii} K in warm neutral or weakly ionized gas, where the species is not the dominant ionisation stage (e.g. Sembach et al. 2000).  Previous observations of the Na\\,{\\sc i} UV lines have been published by only a handful of authors including Hobbs (1978), Vallerga et al. (1993) and Welsh et al. (1997); corresponding results for Ti\\,{\\sc ii} appear in Stokes (1978), Hobbs (1984) and Welsh et al. (1997). In total approximately 110 Ti\\,{\\sc ii} sightlines have been studied to date, although many of these were non-detections.  The results for the Na\\,{\\sc i} UV lines are complementary to previous studies of Na\\,{\\sc i} D. Data for Ti\\,{\\sc ii} thus far indicate that this element appears to be associated with a smoothly distributed neutral medium. Although, unlike Ca\\,{\\sc ii}, Ti\\,{\\sc ii} is the dominant ionisation stage in neutral gas, the Ti\\,{\\sc ii} abundance shows a similar amount of scatter when plotted against $N_{\\rm HI}$ (Wakker \\& Mathis 2000), with the Ti\\,{\\sc ii}/Ca\\,{\\sc ii} abundance ratio being constant over a range of interstellar density conditions (Welsh et al. 1997).  The current paper concerns observations of the Na\\,{\\sc i} UV doublet, Ti\\,{\\sc ii} and Ca\\,{\\sc ii} K towards a sample of 74 early-type stars in the Southern Hemisphere. The wavelengths and transition strengths for the studied lines are shown in Table~\\ref{t_wavef}.   \\begin{table} \\begin{center} \\caption[]{Main transitions studied.} \\label{t_wavef} \\begin{tabular}{lcr} \\hline Transition & $\\lambda_{\\rm air}$ & $f$-value \\\\ &  (\\AA)              &           \\\\ \\hline Na\\,{\\sc i}    & 3302.368  & 0.00921         \\\\ Na\\,{\\sc i}    & 3302.978  & 0.00460         \\\\ Ti\\,{\\sc i}    & 3383.759  & 0.358           \\\\ Ca\\.{\\sc ii}   & 3933.661  & 0.627           \\\\ \\hline \\end{tabular} \\end{center} \\end{table}  In contrast to previous surveys which generally have signal-to-noise (S/N) ratios in the Na\\,{\\sc i} UV lines and Ti\\,{\\sc ii} of around 40, the current paper utilises data that typically have S/N ratios exceeding 250. This enables features with equivalent widths of $<$ 1~m\\AA\\ to be detected. This contrasts with the only previously existing Southern Hemisphere survey of Ti\\,{\\sc ii} by Welsh et al. (1997; hereafter W97) in which, with a S/N of 40, 18 out of 42 of the observed sightlines in Ti\\,{\\sc ii} were non-detections.  In Sect. 2 we describe the sample of stars observed, while Sect. 3 discusses the data analysis and methods used to obtain equivalent widths, $b$-values, velocities of components and column densities, including an estimate of the errors in the derived parameters and comparison with previous work. Sect. 4 presents correlations between the derived column densities and parameters such as reddening and height above the Galactic disc ($z$). Finally, in Sect. 5 we list the main conclusions. A future paper will deal in more detail with the search for intermediate and high velocity gas along the current sightlines. ",
        "conclusions": "\\label{s_conclusions}  We have obtained spectra towards 74 early-type stellar targets from the UVES POP survey, at a velocity resolution of 3.75\\,km~s$^{-1}$. We have studied the interstellar Na\\,{\\sc i} UV doublet at 3302\\,\\AA, Ti\\,{\\sc ii} at 3383\\,\\AA\\ and Ca\\,{\\sc ii} at 3933\\,\\AA. The average S/N ratio per pixel for the spectra are 260, 300 and 430, respectively, which far exceeds that in previous work in a sample of this size. Column densities were measured for each component along each line-of-sight, with the individual components being summed to obtain the total. This quantity was also calculated by the AOD method, which is in excellent agreement with the total generated by the fitting of each individual component. There is also good agreement between the total column densities derived from each line of the Na\\,{\\sc i} UV doublet.  The median of the errors in individual components are 0.06~dex, 0.06~dex and 0.02~dex for Na, Ti and Ca components respectively. In general, we have not studied cases where the stellar line is of a similar width and velocity to the interstellar profile, and so we have avoided stellar contamination of our derived column densities.  \\subsection{General correlations}                             \\label{s_con_corr}  We have investigated the dependence of the total column density of each species along the line-of-sight to a star with distance, reddening and total H\\,{\\sc i} column density. Our stars span distances from 52\\,pc to 6000\\,pc, with the maximum height from to the Galactic plane, $|$z$|$, being 1793\\,pc. We find a strong positive correlation between Ti and Ca column densities and distance, and a similar correlation for Na although with much more scatter. We also find a good correlation between the column densities and $|z|$, although again the correlation is weaker in Na. Our sightlines have values of reddening $E(B-V)$ of up to 1.0, and some correlation between the column density per parsec of each species and the reddening per kiloparsec has been found. We find similar correlations between the H\\,{\\sc i} column density per parsec and the column density per parsec of Ca and Ti. There is evidence of a stronger correlation between Na and H\\,{\\sc i} but given the scatter of the points this may not be significant. Wakker \\& Mathis (2000) have previously found this large dispersion, and have thus cautioned against for example using non-detections of Na to set limits on the distances of High Velocity Clouds. The observed scatter in Na and H\\,{\\sc i} is attributed by these authors to the fact that it is not the dominant ionisation stage in the ISM, so ionisation effects may play a more important role. This certainly calls into question the accuracy of using the Na column density as an indicator of the value for H\\,{\\sc i}. We also find that the column density derived from the Na\\,{\\sc i} D lines, generally used to predict the H\\,{\\sc i} values, appears to underestimate the column density relative to the UV lines that we have studied. Even when the saturation corrections of Savage \\& Sembach \\cite{sav91} are applied to the column density derived from the Na\\,{\\sc i} D lines it still appears to underestimate the Na column density. For Na column densities greater than 13\\,dex the corrected values underestimate the column density by as much as 0.3\\,dex relative to column density derived from the UV lines.  \\subsection{Depletions}                                        \\label{s_con_dep}  The depletion of each species relative to the solar system value has been correlated with $|z|$ and $E(B-V)$. We find some evidence for decreasing depletion with increasing $z$ for both Ti and Ca, with no correlation with depletion and reddening for these species. For Na, we find decreasing depletion with increasing E(B--V), some of which may be caused by the conversion of H\\,{\\sc i} to H$_{2}$.  \\subsection{$N$(Ti\\,{\\sc ii})/$N$(Ca\\,{\\sc ii}) ratio}           \\label{s_con_Ti/Ca}  We find a good correlation between the Ti\\,{\\sc ii} and Ca\\,{\\sc ii} column densities, both in the total values and those from individual components that are observed in both species. The ratio from the total is 0.30 compared to that from the individual components of 0.26. Our ratios compare very well with the value of 0.29~$\\pm$~0.03 calculated by W97. The similarity between our plots of the correlation of Ti and Ca with each of the different parameters reinforces the possibility that the species show good correspondence if the Ca ionisation balance is controlled by the average interstellar radiation field, and the relative Ca and Ti depletions are relatively constant in the clouds (Welty et al. 1996, Wakker \\& Matthis 2000).  \\subsection{Future work}  As discussed in Section 4.4.1, many of our sightlines show strong IVC and HVC features. A future paper will involve a more detailed analysis of these components, using the Villa-Elissa H\\,{\\sc i} survey and the current observations to provide improved distance limits to IVCs and HVCs."
    },
    "0601/astro-ph0601649_arXiv.txt": {
        "abstract": "{ A search for the near-infrared water-ice absorption band was made in a number of very red OH/IR stars which are known to exhibit the 10$\\mu$m silicate absorption. % As a by-product, accurate positions of these highly reddened objects are obtained. We derived a dust mass loss rate for each object by modelling the spectral energy distribution and the gas mass loss rate by solving the equation of motion for the dust drag wind. The derived  mass loss rates show a strong correlation with the silicate optical depth as well as that of the water-ice. The stars have a high mass loss rate ($> 10^{-4}$ M$_{\\odot}$ yr$^{-1}$) with an average gas-to-dust mass ratio of 110. In objects  which show the 3.1$\\mu$m water-ice absorption, the near-IR slope is much steeper than those with no water-ice. Comparison between our calculated mass loss rates and those derived from OH and CO observations indicates that these stars have recently increased their mass loss rates. ",
        "introduction": "In order to estimate the dust mass loss rate in O-rich Asymptotic Giant Branch (AGB) stars, it is common to use the strength of the 10$\\mu$m silicate emission. There seems to be a direct correlation between the 10$\\mu$m excess emission and the derived mass loss rate estimated by CO (Skinner \\& Whitmore \\cite{skinner}). However, this does not apply when the silicate feature becomes self-absorbed or in absorption as mass loss rate increases. On the contrary, it has been known that extreme OH/IR stars with strong silicate absorption show a very weak CO line strength which indicates low derived gas mass loss rate (Heske et al. \\cite{heske}). Justtanont et al. (\\cite{kay96}) investigated one such star, OH26.5+0.6  and found that the discrepancy can be explained by the fact that the star has recently undergone a superwind where the mass loss rate increases by two orders of magnitude. The low rotational transitions of CO mainly trace the cool gas mass loss rate in the outer part of the envelope, which has not been affected by this mass loss increase. The interferometric CO J=1-0 map using BIMA confirms the envelope size is marginally resolved (Fong et al. \\cite{fong}), consistent with a compact CO envelope due to the superwind.  Accompanied by the deep 10$\\mu$m silicate absorption, the near-infrared spectrum of some extreme OH/IR stars shows the absorption at  3.1$\\mu$m due to water-ice. Meyer et al. (\\cite{meyer}) suggested a correlation between the mass loss rate and the column density of water-ice in a few sample of stars observed. So far, there has not been a systematic search for the 3.1$\\mu$m absorption in these stars as a long integration is needed in order to detect the absorption band when the infrared flux is already very low. A medium resolution spectrum is needed in order to discriminate between the type of ice condensed, i.e., amorphous or crystalline. Maldoni et al. (\\cite{maldoni}) suggested a significant fraction of water-ice in crystalline form condensed in the outflow of OH32.8-0.3. The crystalline ice has sharp features at 3.1, 44 and 62$\\mu$m, the latter two were first detected by the observation by the Kuiper Airborne Observatory (Omont et al. \\cite{omont}). The libration band at 12$\\mu$m is more difficult to discern due to the fact that it blends with the strong silicate absorption. Observations of OH/IR stars using the Infrared Space Observatory (ISO) Short- and Long Wavelength Spectrometers (SWS and LWS) show that these stars exhibit signatures of crystalline silicate dust (e.g., Cami et al. \\cite{cami}), as well as water-ice (Sylvester et al. \\cite{roger}). ",
        "conclusions": "We search the circumstellar envelope of extreme OH/IR stars for the presence of water-ice in an attempt to check the reported correlation between the water-ice column density and the mass loss rate (Meyer et al \\cite{meyer}). We found that such a correlation exists in our sample stars which display the 3.1$\\mu$m water-ice absorption. However, not all the stars in our study show the water-ice band. Furthermore, from the depth of the silicate 9.7$\\mu$m band which is an indication of the dust mass loss rate, it is not possible to predict which of the circumstellar shells should display the water-ice band. Stars which show a distinct water-ice absorption band also shows a near-IR deficiency compared to stars with no water-ice detection. A plausible explanation is another dust species which condensed at the same time as the water-ice, e.g., iron (Kemper et al. \\cite{kemper}). It remains a puzzle how the two species with different condensation temperatures are linked. It is also possible that the water-ice is present in a disk around the central star and when viewed along the line-of-sight, gives higher extinction than at an inclined angle.  We found a simple estimate of the mass loss rate from the ratio of the depth of the silicate absorption feature at 9.7$\\mu$m and the continuum flux. The main uncertainty in our estimated dust mass loss rate is the distance to the stars which in many of the cases here is unknown.  Comparison of our derived dust mass loss rate with the expected gas mass loss rate from the OH luminosity and CO rotational observations indicates that these extreme OH/IR stars have recently undergone an increase in the mass loss rate by a factor of $\\geq$ 10. The maximum timescale for the superwind is less than the time it takes for the mass loss to reach the H$_{2}$O photodissociation radius, i.e., $\\sim$ 2\\,000 yrs. In some cases, this can be much smaller as the dynamical mass loss rate is higher than the mass loss rate derived from the OH observations. High angular resolution observations ($\\leq 0.1^{\\prime\\prime}$) will be able to confirm the extent of the superwind."
    },
    "0601/astro-ph0601155_arXiv.txt": {
        "abstract": "We discuss the detectability of high-redshift galaxies via \\CII 158$\\mu$m line emission by coupling an analytic model with cosmological Smoothed Particle Hydrodynamics (SPH) simulations that are based on the concordance $\\Lambda$ cold dark matter (CDM) model. Our analytic model describes a multiphase interstellar medium (ISM) irradiated by the far ultra-violet (FUV) radiation from local star-forming regions, and it calculates thermal and ionization equilibrium between cooling and heating.  The model allows us to predict the mass fraction of a cold neutral medium (CNM) embedded in a warm neutral medium (WNM). Our cosmological SPH simulations include a treatment of radiative cooling/heating, star formation, and feedback effects from supernovae and galactic winds. Using our method, we make predictions for the \\CII luminosity from high-redshift galaxies which can be directly compared with upcoming observations by the {\\it Atacama Large Millimeter Array} (ALMA) and the {\\it Space Infrared Telescope for Cosmology and Astrophysics} (SPICA). We find that the number density of high-redshift galaxies detectable by ALMA and SPICA via \\CII emission depends significantly on the amount of neutral gas which is highly uncertain. Our calculations suggest that, in a CDM universe, most \\CII sources at $z=3$ are faint objects with $\\Snu < 0.01$\\,mJy. Lyman-break galaxies (LBGs) brighter than $\\Rab=23.5$ mag are expected to have flux densities $\\Snu = 1-3$\\,mJy depending on the strength of galactic wind feedback. The recommended observing strategy for ALMA and SPICA is to aim at very bright LBGs or star-forming DRG/BzK galaxies. ",
        "introduction": "Up to now, the high-redshift galaxies detected in large numbers are observed by the radiation from stars in rest-frame UV to optical wavelengths. Such studies have been fruitful in delineating the properties of star-forming galaxies at $z\\gtsim 3$ that are bright in UV; i.e., the Lyman-break galaxies (LBGs) \\citep[e.g.,][]{Steidel99, Ade00, Shapley01, Iwata03, Steidel03, Ouchi04a, Ouchi04b}. Within the hierarchical galaxy formation paradigm based on the cold dark matter (CDM) model \\citep{Blumenthal84, Davis85}, it has been shown that the spectro-photometric properties of LBGs can be accounted for if we associate them with relatively massive galaxies ($M_\\star \\gtsim 10^{10} \\Msun$) situated in large dark matter halos ($M_{\\rm halo}\\gtsim 10^{12} \\Msun$) \\citep[e.g.][]{Mo96, Dave99, Nag02, Weinberg02, NSHM, Nachos2, Nachos3, Night06, Finlator06}.   However, observing the UV light emitted by the young stars does not directly tell us about how much gas there is in the galaxy. In order to obtain a full picture of galaxy formation, we need to find answers to the questions such as: 1) How much neutral gas is available for star formation in high-redshift galaxies? 2) What is the baryonic mass (in particular, the neutral hydrogen mass) fraction as a function of halo mass, and how does it evolve as a function of redshift? 3) How does the volume-averaged neutral gas mass density evolve as a function of redshift? 4) What are the main cooling and heating processes of the ISM in high-redshift galaxies?  One of the ways to detect the neutral hydrogen in high-redshift galaxies is to search for damped Lyman-$\\alpha$ absorbers \\citep[DLAs,][for a review]{Wolfe05} in quasar absorption lines. DLAs are defined as quasar absorption systems with a neutral hydrogen column density $\\NHI>2\\times 10^{20} \\cm^{-2}$, a threshold column density that effectively guarantees gas neutrality at high redshifts \\citep{Wolfe86, Wolfe05}. Since DLAs are known to dominate the neutral hydrogen mass density at $z\\gtsim 3$ \\citep[e.g.,][]{Lan95, Storr00}, it is expected that they represent a significant reservoir of neutral gas for star formation. A large number of DLAs have been discovered at $z\\gtsim 3$, and they have proven to be one of the best probes of neutral gas in high-redshift galaxies, complementary to the optical to infra-red (IR) observations. A recent search for DLAs in the Sloan Digital Sky Survey (SDSS) data archive yielded a large sample of over 500 DLAs at $z>2.2$ \\citep*{Pro04, Pro05}. These observational results are beginning to put stringent constraints on theoretical/numerical models of galaxy formation. Furthermore, \\citet{Wolfe03a} opened a new window for probing star formation in high-redshift DLA galaxies utilizing the C{\\sc ii}$^*$ absorption lines, and this paper is motivated by their work.  The C{\\sc ii}$^*$ absorption lines originate from the excited $^2P_{3/2}$ state in the ground 2s$^2$\\,2p term of C$^+$. This is the same state that gives rise to \\CII emission at $\\lambda=157.74\\,\\mu$m through the $^2P_{3/2} \\rightarrow ^2P_{1/2}$ fine structure transition. Because this is the dominant coolant of diffuse ISM at temperatures $T \\le 5000$ K \\citep{Dalgarno72, Tielens85, Wolfire95, Lehner04}, detection of \\CII emission is potentially another method for finding cold neutral gas at high redshifts \\citep{Petrosian69, Loeb93}.  \\CII emission was first detected towards the inner regions of gas-rich spiral galaxies and starburst galaxies; i.e., dense star-forming gas irradiated by UV radiation from young star-forming regions near galactic centers \\citep{Russell80, Crawford85, Stacey91, Carral94}. In this case, the brightness of the emission line suggests that it is produced in warm ($T > 200$\\,K), dense ($n_{\\rm H} > 10^3$\\,cm$^{-3}$) photodissociated regions (PDRs) with pressures $P/k_B=10^4 - 10^7$\\, K\\,cm$^{-3}$ at the interface between giant molecular clouds and fully ionized media. The \\CII luminosity accounts for $0.1 - 1$\\% of the total far-IR luminosity of the nuclear regions.  However, later space-based observations with the Long Wavelength Spectrometer (LWS) on-board ESA's Infrared Space Observatory (ISO) have shown that, for quiescent late-type galaxies like NGC6946, the \\CII emission from extended diffuse gas over the entire disk could also be significant compared to that from dense compact star-forming gas \\citep[e.g.,][]{Madden93, Malhotra97, Leech99, Malhotra01, Contursi02}. The diffuse cold components of the ISM seems to be at least as important sources of the \\CII emission as are compact regions \\citep[e.g.,][]{Sauty98, Pierini99, Pierini01, Contursi02}.  More recently, \\CII emission has also been detected in the highest redshift SDSS quasar at $z=6.42$ \\citep{Maiolino05}. Quasars are considered to host significant amounts of molecular gas ($M_{\\rm gas}\\approx 10^{11}\\Msun$) and dust ($M_{\\rm dust}\\approx 10^8\\Msun$), and are therefore forming stars at a significant rate of $SFR \\gtrsim 1000\\,\\Msun\\, \\yr^{-1}$ \\citep[e.g.,][]{Omont96, Omont01, Omont03, Carilli01, Walter03, Beelen04, Carilli04}. Given this large amount of gas contained in them, they should be very bright in \\CII luminosity, allowing them to be easily detected even at very high redshift. However, bright quasars have much lower comoving space density than LBGs and are not representative of the entire galaxy population; they are special sources hosting supermassive black holes, and may be in unusual dynamical states if their activity is triggered by mergers \\citep[e.g.,][]{Hopk05}.  The detection of \\CII emission from normal (i.e., not quasars or active galactic nuclei [AGN]) high-redshift star-forming galaxies would be a more {\\it direct} way than the absorption line technique to find out about the amount of neutral gas available for star formation. Such detections of \\CII emission from high-redshift galaxies, for example, by ALMA and SPICA, would set an important constraint on the theory of galaxy formation. In particular interferometric maps of \\CII emission would reveal the spatial dimensions of the DLAs, a property that so far has eluded detection. By combining these measurements with the velocity widths of the \\CII emission lines, it will be possible to infer the dark matter masses of these objects, and the luminosity of the \\CII emission lines would tell us the heating rate of the gas.  In this paper we study the detectability of high-redshift DLA galaxies by coupling an analytic model with cosmological SPH simulations that are based on the concordance $\\Lambda$CDM model, and make predictions for upcoming observations by ALMA and SPICA. We also compare the computed \\CII luminosity function with the observed LBG luminosity function, and discuss the connection between DLA galaxies and LBGs. \\citet{Stark97} discussed the potential measurement of \\CII emission from high-redshift galaxies, but he made various assumptions for the \\CII luminosity function and the evolution of the characteristic luminosity based on the local observations and numerical simulations, and focused more on the observing time and conditions rather than the intrinsic nature of high-redshift galaxies. The present study extends his original work, and is intended to be more physically motivated by using full cosmological hydrodynamic simulations which simulate structure formation from early epochs to the present time. We note that \\citet{Suginohara99} studied the detectability of atomic emission lines of carbon, nitrogen, and oxygen, but they focused on the intensity fluctuations owing to sources at $z\\sim 10$ rather than the detection of individual objects. ",
        "conclusions": "\\label{sec:conclusion}  We have coupled state-of-the-art cosmological SPH simulations with an analytic model of a multiphase ISM in order to compute the \\CII emission from galaxies at $z=3$. We find that, in a $\\Lam$CDM universe, the majority of the sources are very faint with $\\Snu < 0.1$\\,mJy. This is presumably a generic prediction of the CDM model owing to an increasing number of low-mass dark matter halos with decreasing halo mass. If our model prediction on the faintness of the \\CII sources is indeed correct, then it will be difficult for ALMA and SPICA to detect normal (i.e., not quasars, active AGNs, or strong submillimeter sources) high-redshift galaxies via \\CII emission in large number, as the sensitivity limit of both telescopes are expected to be $\\Snu \\sim 0.1 - 1$\\,mJy. The recommended observing strategy is therefore to focus on very bright LBGs with $\\Rab \\ltsim 24$ mag that are pre-selected through optical imaging and spectroscopically known redshifts. Since there is a large body of spectroscopic data on LBGs at $z\\gtsim 3$, it should not be a problem to come up with observing targets.  Our calculation shows that the brightest LBGs with $\\Rab \\sim 23.5$ mag could have flux densities $\\Snu = 1-3$\\,mJy depending on the strength of galactic wind feedback.  Some caveats to note for the present work is that we have assumed that the physical conditions of neutral gas in DLAs and LBGs are the same, and that the dust-to-gas ratios of the gas are similar. If these assumptions are inappropriate for the LBGs at $z\\sim 3$, then the above conclusions might not be entirely accurate. Our multiphase ISM model also assumed the SMC-type (silicates) dust composition which is less efficient for heating than other models such as the Galactic (carbonaceous) or the LMC-type models of dust grains. The true nature of dust in DLAs could be different from what we assumed here, and we expect a factor of two uncertainty in the dust-to-gas ratio. Note, however, that the metallicity of LBGs approaches solar, which is significantly higher than for DLAs. As a result, it is natural to expect that gas surrounding LBGs would have a dust-to-gas ratio higher than for the average DLA. Since the heating rate of the gas is proportional to dust-to-gas ratio, the \\CII emission from LBGs is likely to be at the upper end of our calculation. In passing we note that star-forming DRG/BzK galaxies \\citep{Dokkum04, Daddi04b} would also be good candidates for \\CII observation as they are considered to have higher star formation rates and dust content than LBGs and hence more luminous in \\CII luminosity. But their space density is lower by a factor of $\\sim 10$ than that of LBGs.  We have some confidence on the reliability of the properties of bright galaxies in the simulation, because we know from our previous work \\citep{NSHM, Night06, Finlator06} that the simulations can reproduce the observed properties of LBGs at $z=3-6$ relatively well. However the neutral gas is not well resolved near the resolution limit of each simulation, and this introduces some uncertainty for low-mass galaxies (except for the Q5 run in which low-mass galaxies are well modeled). For this reason our scaling laws were mainly determined using the data for the massive galaxies in the simulation. There is also an intrinsic scatter in the CNM mass as a function of halo mass around the scaling relation (Figure~\\ref{fig:nhmass}). As a result, we expect 30\\% uncertainty in our predictions of the CNM mass fractions for high-mass halos, and somewhat larger uncertainties for lower mass halos.  Even if the observation of normal high-redshift galaxies is difficult, the reward of \\CII detection from such galaxies would be quite significant, because such measurements would directly constrain the amount of neutral gas and the properties of ambient ISM in high-redshift galaxies which is otherwise difficult to do. The observations of DLAs \\citep[e.g.][]{Pro05, Wolfe05} have given us tremendous insights on the properties of neutral gas in high-redshift galaxies already, but \\CII measurements will provide complementary information to further constrain theoretical/numerical models of galaxy formation. The combination of interferometric maps and the velocity widths of the \\CII emission could in principle tell us about the physical sizes of DLAs and the mass of hosting dark matter halos, as well as the physical properties of the gas such as the heating rate.  This paper is just a first step towards such a goal, and a number of improvements in the analysis are needed in the future. For example, in the current simulations, an optically thin approximation was assumed throughout when solving for ionization equilibrium, and a simple radiative transfer calculation assuming a disk geometry was used to approximate the effect of radiative transfer from local star-forming regions. As computing power increases, a more accurate treatment of radiative transfer \\citep[e.g.,][]{Razoumov05} and ISM physics on small-scales will become possible using direct information from simulations, such as star formation rates and metallicity of the gas."
    },
    "0601/astro-ph0601225_arXiv.txt": {
        "abstract": "We utilize nonlinear pulsation models to reproduce the observed light and color curves for two samples of bump Cepheid variables, 19 from the Large Magellanic Cloud and 9 from the Small Magellanic Cloud.  This analysis determines the fundamental parameters mass, luminosity, effective temperature, metallicity, distance and reddening for the sample of stars.  The use of light curve shape alone to determine metallicity is a new modelling technique introduced here.  The metallicity, distance and reddening distributions for the two samples are in agreement with those of similar stellar populations in the literature. The distance modulus of the Large Magellanic Cloud is determined to be 18.54$\\pm$0.018 and the distance modulus of the Small Magellanic Cloud is determined to be 18.93$\\pm$0.024.  The mean Cepheid metallicities are $Z = 0.0091\\pm0.0007$ and $0.0050\\pm0.0005$ for the LMC and SMC, respectively.  The masses derived from pulsation analysis are significantly less than those predicted by stellar evolutionary models with no or mild convective core overshoot. We show that this discrepancy can not be accounted for by uncertainties in our input opacities or in mass-loss physics. We interpret the observed mass discrepancy in terms of enhanced internal mixing in the vicinity of the convective core during the main-sequence lifetime and find that the overshoot parameter $\\Lambda_{c}$ rises from 0.688$\\pm$0.009H$_{\\rm p}$ at the mean LMC metallicity to 0.746$\\pm$0.009H$_{\\rm p}$ in the SMC. ",
        "introduction": "Wide-field photometric studies such as the OGLE and MACHO projects have left a legacy of well-sampled highly uniform photometry for the Cepheid populations of the Magellanic Clouds. In the present paper we compare these data to the predictions of stellar pulsation models. A number of recent studies have shown that nonlinear pulsation models are able to reproduce both the gross features of Cepheids light curves such as periods and amplitudes as well as much smaller morphological features such as resonance features seen in bump Cepheids \\citep{woo97,bon02}.  The morphology of the light curve proves highly sensitive to the physical parameters of mass, luminosity and effective temperature. By matching the model light and color curves to those observed we can constrain the physical parameters of the star to greater precision than available with other techniques \\citep{kel02}. The distance and reddening to each Cepheid can also be determined. The resulting distances offer an independent method of distance determination to the Magellanic Clouds: one which is solely based on the physics of stellar interiors.  Mass estimates derived from bump Cepheids have caused much debate since \\cite{sto69} first showed that pulsation masses for bump Cepheids are significantly lower than those predicted from stellar evolution. Largely in response to this mass discrepancy, more detailed modelling of opacities was undertaken by the OPAL \\citep{rog92} and Opacity Projects \\citep{sea94}. The resultant enhanced metal opacities resolved much of the mass discrepancy \\citep{mos92}.  However a series of studies of galactic \\citep{cap05} and LMC \\citep{woo97,kel02,bon02} bump Cepheids have shown that the masses of bump Cepheids remain smaller than predicted by evolution. This discrepancy amounts to 15\\% in mass or, seen another way, bump Cepheids appear $\\sim20\\%$ more luminous for a given mass \\citep{kel02}.  The luminosity of a Cepheid is critically dependent on the mass of the He core established during the star's main-sequence life. The mass of the He-core is dependent on the extent of the convective core during core-H burning. Classical models define the limit to convection via the Schwarzschild criterion. This places the boundary to convection at the radius at which the buoyant force drops to zero. However, the temperature and density regime in the vicinity of the convective boundary of the main-sequence Cepheid progenitor are such that restorative forces in the region formally stable to convection are mild and some significant level of overshoot of the classical boundary is expected \\citep{zahn91}.  It is typical to parameterize the level of convective core overshoot using the mixing-length formulation where gas packets progress a distance given by $\\Lambda_{c}$ pressure-scale heights ($H_{\\rm p}$) into the classically stable region. A star with larger core overshoot is able to draw upon more H and hence lives longer on the main-sequence, develops a more massive He core and is more luminous than classical models during the post-main-sequence evolution.  Our description of convection is the weakest point in our understanding of the physics of massive stars. Numerical modelling of core convection requires a description of the turbulence field at all scales. Three-dimensional hydrodynamical calculations capable of adequate resolution have only recently become a possibility and are in their infancy \\citep{djehuty}. At present we must rely on observation for constraint.  Observational determinations of the magnitude of convective core overshoot (CCO) have focused on populations of intermediate mass stars ($M=5-12$ M$_{\\odot}$) where the signature of overshoot is expected to be most clearly seen. The study of \\cite{mer86} examined a range of galactic clusters with main-sequence turn-offs in mass range 9-15M$_{\\odot}$ and derived overshoot parameter $\\Lambda_{c} \\sim 0.6$$H_{\\rm p}$.\\footnote{We quantify CCO using the formalism of \\cite{bre81} who use an overshoot parameter $\\Lambda_{c}$ pressure scale heights.  Note that $\\Lambda_{c}$ is a factor of 2.0 times the overshoot parameter $d_{\\rm over}/H_{\\rm p}$ in the formalism of the Geneva group \\citep{schaller92}.} \\cite{chiosiARAA92} summarize a number of similar studies converging on a mild efficiency of overshoot $\\Lambda_{c} \\sim 0.5$. This is the basis for the level of CCO included in the stellar evolutionary models for $M > 1.3$M$_{\\odot}$ by \\cite{gir00}.  A number of subsequent studies have focused on the intermediate age LMC cluster NGC 1866 whose large population of post-main-sequence stars was proposed as proof both for \\citep{bar02} and against \\citep{testa99} mild overshoot. \\cite{kel01} examined the populations of a series of young populous clusters in the Magellanic Clouds and found $\\Lambda_{c}=0.67\\pm0.20$$H_{\\rm p}$ for the SMC cluster NGC 330 and mean of $\\Lambda_{c}=0.61\\pm0.12$$H_{\\rm p}$ for three similar clusters in the LMC. This result does not indicate any strong dependence of $\\Lambda_{c}$ on metallicity. The subsequent study of \\cite{cor02} examined the field population of the Magellanic Clouds and in contrast found evidence for a strong metallicity dependence in the degree of overshoot ($\\Lambda_{c}=0.8$$H_{\\rm p}$ and 0.2$H_{\\rm p}$ for the SMC and LMC respectively. Such studies based on the color-magnitude diagrams of stellar populations do face uncertainty due to contamination by unresolved binary systems (and significant blending of unrelated stars in the case of the Magellanic Clouds), contamination by the surrounding field population, small numbers of stars on the upper main-sequence and systematics introduced by bolometric corrections to place the population on the H-R diagram.  \\cite{bon02} reported the results of light curve modelling of two LMC bump Cepheids and found that they were able to match the luminosities of the Cepheids using masses 15\\% less than predicted by evolutionary models which neglect convective core overshoot and mass-loss. \\cite{kel02} (Paper I) presented results of pulsation modelling for a sample of 20 LMC bump Cepheids. Paper I found a level of core overshoot of $\\Lambda_{c}=0.65\\pm0.03$$H_{\\rm p}$ under the assumption of a LMC abundance of Z=0.008. The models of Paper I incorporated convective energy transfer in the Cepheid envelope through the mixing-length approximation. This approximation is expected to break down at cooler temperatures as the convective region of the Cepheid becomes a substantial fraction of the envelope. A consequence is that these models do not reproduce the red edge of the instability strip. Redward of the red edge the pulsation growth rate remains above zero. Consequently in Paper I we were restricted to bump Cepheids close to the blue edge of the instability strip.  In addition, \\cite{feu00} demonstrates that neglect of the effects of convection can shift the blue edge to temperatures on the order of 400K cooler than predicted by models incorporating the effects of turbulent convection.  In this paper, we present updated models which include convection and turbulent eddy viscosity. The formulation of our models is discussed in Section 2. In Section 3 we analyze for the first time a sample 9 SMC stars and reanalyse 19 LMC Cepheids for comparison. We then discuss the implication of our results for the level of convective core overshoot and its dependence on metallicity in Section 4. ",
        "conclusions": "By matching the observed light and color curves of 19 LMC and 9 SMC bump Cepheids we have been able to place tight constraints on the fundamental stellar parameters of mass, luminosity, effective temperature and metallicity and the secondary parameters of distance and reddening.  The use of light curve fitting to derive metallicity is a new technique developed here.  We have revisited the mass discrepancy that exists between the masses derived from pulsation modelling and the masses predicted from stellar evolution for bump Cepheids. The derived pulsational masses are significantly less than expected from stellar evolution models that do not incorporate extension to the convective core during main-sequence evolution.  Significant overshoot of the convective core during the main-sequence phase is required to bring pulsation and evolutionary masses into agreement.  In addition, we find that the amount of overshoot is a function of metallicity.  The level of convective core overshoot rises from 0.688$\\pm$0.009H$_{\\rm p}$ at LMC metallicity to 0.746$\\pm$0.009H$_{\\rm p}$ at typical SMC metallicity. This trend with metallicity is qualitatively in line with expectation from models that incorporate rotationally-induced mixing provided rotation increases with decreasing metallicity.  Work is currently underway to further refine our technique by the comparison of radial velocity measurements with those predicted by our models. This offers a more stringent test of the details of our pulsation models without limitations imposed by systematic photometric uncertainties (viz.\\ bolometric corrections, photometric transformations and zeropoints)."
    },
    "0601/astro-ph0601013_arXiv.txt": {
        "abstract": "We have studied dust evolution in a quiescent or turbulent protoplanetary disk by numerically solving coagulation equation for settling dust particles, using the minimum mass solar nebular model. As a result, if we assume an ideally quiescent disk, the dust particles  settle toward the disk midplane to form a gravitationally unstable layer within $2\\times 10^3$--$4\\times 10^4$yr at 1--30 AU, which is in good agreement with an analytic calculation by Nakagawa, Sekiya, \\& Hayashi (1986) although they did not take into account the particle size distribution explicitly. In an opposite extreme case of a globally turbulent disk, on the other hand, the dust particles fluctuate owing to turbulent motion of the gas and most particles become large enough to move inward very rapidly within 70--$3\\times 10^4$yr at 1--30 AU, depending on the strength of turbulence. Our result suggests that global turbulent motion should cease for the planetesimal formation in protoplanetary disks. ",
        "introduction": "It is believed that particle settling and growth are important processes leading to the planet formation in protoplanetary disks. Observationally some evidences of dust size growth have been proposed based on dust continuum emission from the protoplanetary disks, such as smaller power-law indices of spectral energy distributions (SEDs) in sub-mm to cm wavelength bands, and fainter trapezium feature of 10$\\micron$ silicate emission, compared with those of the interstellar dust grains (e.g., Beckwith \\& Sargent 1991; Miyake \\& Nakagawa 1993; Beckwith et al. 2000; Kitamura et al. 2002; van Boekel et al. 2003; Przygodda et al. 2003; Wilner et al. 2005). Meanwhile, theoretically the initial stage of dust evolution when micron-sized interstellar dust particles grow into centimeter-sized particles with settling toward the disk midplane have been studied analytically and numerically; for example, Weidenschilling (1980) and Nakagawa et al. (1981, 1986)'s works in a quiescent disk, and Cuzzi et al. (1993) and Weidenschilling (1997, 2004)'s works in local turbulence induced by the shear between the dust layer and the gas near the disk midplane (see also Weidenschilling \\& Cuzzi 1993 and references therein).  In addition to the shear-induced turbulence, it is thought that there exists global turbulent motion in the protoplanetary disks, which is caused by thermal convective and/or magneto-rotational instabilities (e.g., Lin \\& Papaloizou 1980; Balbus \\& Hawley 1991). The turbulent motion of the gas is known to affect the dust evolution processes; turbulence induced relative motion increases the mutual collision rate, that is, the growth rate of the dust particles (e.g., V\\\"{o}lk et al. 1980), turbulent mixing motion lets the particles move diffusively (e.g., Cuzzi et al. 1993), and turbulent eddies trap the dust particles (e.g., Klahr \\& Henning 1997; Cuzzi et al. 2001; Johansen et al. 2004). On the other hand, the disk instabilities, namely, the existence of turbulent regions depend on the spatial and size distributions of the dust particles (e.g., Mizuno et al. 1988; Sano et al. 2000; Nomura 2004). Therefore, self-consistent treatment of the evolution of the dust particles and the turbulent regions (the disk instabilities) is needed in order to understand the very beginning of planet formation process in protoplanetary disks before the dust particles settle toward the disk midplane and form a dusty layer, which could lead to the planetesimal formation.  Moreover, recent observations have been providing a huge amount of data of spectra and SEDs of dust continuum emission as well as molecular line emissions from protoplanetary disks. Theoretically reproducing these observational data have been also developed using detailed disk models (e.g., Kenyon \\& Hartmann 1987; Miyake \\& Nakagawa 1995; Chiang \\& Goldreich 1997; D'Alessio et al. 1998; Dullemond et al. 2001; van Zadelhoff et al. 2001; Aikawa et al. 2002; Nomura \\& Millar 2005). Although rather simple dust models have been used in those previous models, the dust properties affect the physical and chemical structure of the disks, and then the observable properties very much (e.g., D'Alessio et al. 2001; Aikawa \\& Nomura 2005). Thus, we should carefully model the dust evolution in the disks and compare the model predictions with the observational data in order to interpret the observations and understand what is actually going on in protoplanetary disks. The SEDs of young stellar objects have been tried to be modeled by numerically simulating the dust size growth and settling processes in the disks (Suttner \\& Yorke 2001; Tanaka et al. 2005; Dullemond \\& Dominik 2005).  In this paper, we have studied basic behavior of dust evolution in a quiescent or globally turbulent protoplanetary disk, especially focusing on time scales and growing dust size. It is done by numerically solving the coagulation equation, using a simple disk model, as a first step of understanding the initial stage of planet formation and modeling observational properties of protoplanetary disks. In the following section, we present the disk model as well as the basic equations for the vertical and radial motion and the coagulation of dust particles. In \\S 3, we numerically calculate the dust size growth and settling toward the midplane in a quiescent disk, and compare the result with that obtained analytically by Nakagawa et al. (1986; hereafter NSH86). We also discuss the dust evolution in a globally turbulent disk in \\S 4. Finally, the results are summarized in \\S 5. ",
        "conclusions": "We have investigated the dust size growth and settling toward the disk midplane in a quiescent or turbulent protoplanetary disk by numerically solving coagulation equation for settling dust particles.  Our result shows that the dust particles settle toward the disk midplane to form a gravitationally unstable layer at a short timescale ($2\\times 10^3$--$4\\times 10^4$yr at $R=$1--30 AU) if we assume an ideally quiescent disk. The radii of the largest dust particles just before the formation of the unstable layer are 20--1 cm at 1--30 AU. The resulting settling time and evolution of the largest dust radius in our numerical simulation are in good agreement with those obtained by the analytic calculation in NSH86, although they did not take into account the dust size distribution explicitly. This is because most dust mass is included in the dust particles with the largest sizes, and these particles control the dust settling time.  Also, we have discussed the dust evolution in an opposite extreme case of a globally turbulent disk to find that the dust particles are forced to fluctuate by turbulent motion of the gas, and grow to be large enough (32--6 cm at 1--30AU) to move inward very rapidly within a short timescale ($70$--$3\\times 10^4$yr at 1--30 AU). Thus, our result suggests that the disk should be quiescent or the global turbulent motion should cease before most mass of dust particles accrete onto the central star, in order to form planetesimals in protoplanetary disks. Self-consistent treatment of the evolution of the globally turbulent regions and the dust evolution processes is needed in future work. In addition, the dust evolution in a locally turbulent motion induced by the shear between the dust layer and the gas near the disk midplane should be investigated for further understanding of the planetesimal formation process."
    },
    "0601/astro-ph0601539_arXiv.txt": {
        "abstract": "We present observations of dwarf nova oscillations (DNOs), longer period dwarf nova oscillations (lpDNOs), and quasi-periodic oscillations (QPOs) in 13 cataclysmic variable stars (CVs).  In the six systems WW Cet, BP CrA, BR Lup, HP Nor, AG Hya, and V1193 Ori, rapid, quasi-coherent oscillations are detected for the first time.  For the remainder of the systems discussed we have observed more classes of oscillations, in addition to the rapid oscillations they were already known to display, or previously unknown aspects of the behaviour of oscillations.  The period of a QPO in RU Peg is seen to change by 84\\% over ten nights of the decline from outburst---the largest evolution of a QPO period observed to date.  A period-luminosity relation similar to the relation that has long been known to apply to DNOs is found for lpDNOs in X Leo; this is the first clear case of lpDNO frequency scaling with accretion luminosity.  WX Hyi and V893 Sco are added to the small list of dwarf novae that have shown oscillations in quiescence. ",
        "introduction": "Cataclysmic variable stars (CVs) vary in brightness on time scales from seconds to millions of years; the variability ranges in coherence from highly stable orbital and white dwarf spin cycles to aperiodic flickering.  A frequently observed property of CVs that has proved to be one of the most difficult to explain satisfactorily is the quasi-coherent variability seen on time scales of seconds to tens of minutes.  Dwarf nova oscillations (DNOs) and quasi-periodic oscillations (QPOs) were discovered a few decades ago (\\citealt{p38}; \\citealt{p11}).  More recently, Warner, Woudt \\& Pretorius (2003, hereafter Paper III) introduced a third class---the longer period dwarf nova oscillations (lpDNOs).  The phenomenology of DNOs, lpDNOs, and QPOs, together with models to describe them, are reviewed by Warner (2004, hereafter W04).  Amongst the most important characteristics of DNOs are the period-luminosity relation (DNO frequency scales with accretion luminosity) and the time scales (typical DNO periods are $P_{DNO} \\sim 10$~s).  The periods of normal QPOs ($P_{QPO}$) are longer, and the relation $P_{QPO}\\approx 16\\times P_{DNO}$ holds for DNOs and QPOs in a given system.  The period of an lpDNO ($P_{lpDNO}$) is typically midway between $P_{DNO}$ and $P_{QPO}$ in that system, i.e. $P_{lpDNO}\\approx 4\\times P_{DNO} \\approx 1/4\\times P_{QPO}$.  Rapid oscillations are almost exclusively associated with CVs in states of high mass transfer rate ($\\dot{M}$), and are usually not observed in quiescent dwarf novae.  W04 also pointed out that some systems have QPOs at two separate time scales: QPOs with periods $\\sim 1\\,000$~s are seen in addition to QPOs that follow the relation $P_{QPO}\\approx 16\\times P_{DNO}$.   We will refer to the second group of QPOs as kilosecond QPOs.  Since rapid, quasi-coherent oscillations occur in a large fraction of CVs, the physical mechanisms producing DNOs, lpDNOs, and QPOs are likely to be of quite general significance to accretion disc physics.  Judging by the time scale and coherence of these oscillations, DNOs and lpDNOs most likely originate in the inner disc or on the white dwarf itself, while the accretion disc must be the source of the QPOs.  The most viable model of DNOs and lpDNOs ascribes them to the interaction of a weak white dwarf magnetic field with the inner accretion flow (a detailed description of this model for DNOs may be found in \\citealt{dno2}, hereafter Paper II).  The rapid oscillations therefore deserve attention not only because they are a common property of disc-fed accretion in CVs, but also because they probably hold diagnostic potential in the study of the accretion process and the interaction between discs and magnetic fields.  The discovery of a correlation between the ratio of time scales of the oscillations in CVs and those in low-mass X-ray binaries (LMXBs) has further renewed interest in DNOs and QPOs (Woudt \\& Warner 2002, hereafter Paper I; Paper II; \\citealt{p17}; Paper III).  Accretion processes are more readily studied in CVs than in LMXBs---observations of the oscillations in CVs can be obtained with modest-sized ground based telescopes, and may provide insight also into the QPOs found in LMXBs.  This is the fifth in a series of papers on rapid, quasi-coherent oscillations in CVs; we present more observations of these oscillations in the next four sections, and discuss the results in Section \\ref{sec:discuss}.  A more complete account of this work may be found in \\citet{thesis}. ",
        "conclusions": "\\label{sec:discuss}  We have detected rapid, quasi-coherent oscillations in six CVs for the first time, and have observed new aspects of the phenomenology of oscillations in seven more systems.  The QPOs in V1193 Ori show signs of an underlying periodic signal, and are similar to kilosecond QPOs seen in V442 Oph, RX J1643.7+3402, and several more SW Sex stars, which have been interpreted as a low coherence manifestation of the underlying stable rotation of a magnetic white dwarf \\citep{p94}.  However, V1193 Ori shares none of the other characteristics of SW Sex stars.  RU Peg displays QPO periods ranging smoothly over almost an order of magnitude; this may call into question the existence of two clearly separate sets of QPOs discussed by W04.  In HX Peg, however, a QPO and kilosecond QPO have been detected simultaneously, implying that these are distinct phenomena.  Many contradictory and confusing claims about periodic modulations in V426 Oph have been made.  The detection of a DNO in this system casts further doubt on its status as an intermediate polar.  An lpDNO in X Leo was found to follow a period-luminosity relation similar to what is observed for DNOs.  Very little is presently known about lpDNOs; there are probably many features of their phenomenology that have not yet surfaced.  As soon as lpDNOs were identified they were found to be fairly common---although not as common as DNOs---and have up until the present been seen in about 24 CVs (in around one half of these systems the oscillations were at first reported as QPOs or DNOs).  It is perhaps surprising that lpDNOs remained unrecognised for so long; however, there is nowhere in the earlier literature an amplitude spectrum that makes the existence of a third class of oscillations as clearly evident as e.g.\\ the Fourier transform shown as fig. 3 of W04.  Detections of DNOs, QPOs, and lpDNOs in dwarf novae at minimum light are review in W04 (see also \\citealt{brianst}); to these can now be added a DNO in V893 Sco in quiescence, and QPOs in WX Hyi at minimum.  Rapid oscillations are probably more commonly present in quiescent dwarf novae than is currently thought; it is certainly true that oscillations occur much more frequently in outburst than in quiescence, which means that they are not often looked for in quiescent light curves.  For this reason, the observational record is at this point so heavily biased towards outbursting dwarf novae that it is impossible to make a secure estimate of the frequency with which oscillations occur in quiescence.  In summary, we have presented \\begin{enumerate} \\item the first detection of rapid, quasi-coherent oscillations in WW Cet, BP CrA, BR Lup, HP Nor, AG Hya, and V1193 Ori, \\item the detection of additional classes of oscillations in V426 Oph, V1159 Ori, and X Leo, \\item large evolution of a QPO period over 10 nights of the decline from outburst in RU Peg, \\item observations of normal QPOs and kilosecond QPOs simultaneously present in HX Peg, \\item a period-luminosity relation for lpDNOs in X Leo, and \\item the detection of oscillations in quiescence in the dwarf novae WX Hyi and V893 Sco. \\end{enumerate}"
    },
    "0601/astro-ph0601141_arXiv.txt": {
        "abstract": "We present \\xmm\\ observations of a complete sample of five archetypal young radio-loud AGN, also known as Compact Symmetric Objects (CSO) or Gigahertz Peaked Spectrum (GPS) sources. They are among the brightest and best studied GPS/CSO sources in the sky, with radio powers in the range L$_{\\rm{5GHz}}$=10$^{43-44}$ erg s$^{-1}$ and with four sources having measured kinematic ages of 570 to 3000 years. All five sources are detected, and have 2-10 keV luminosities ranging from 0.5 to 4.8$\\times 10^{44}$ erg s$^{-1}$. A detailed analysis was performed, comparing the X-ray luminosities and $\\NH$\\ absorption column densities of the GPS/CSO galaxies with their optical and radio properties, and with those of the general population of radio galaxies. We find that,\\\\ 1) GPS/CSO galaxies show a wide range in absorption column densities with a distribution not different from that of the general population of radio galaxies. We therefore find no evidence that GPS/CSO galaxies could reside in a significantly more dense circumnuclear environment such that they could be ``frustrated'' radio sources $-$ hampered in their development.\\\\ 2) The ratio of radio to X-ray luminosity is significantly higher for GPS/CSO sources than for classical radio sources. This is consistent with an evolution scenario in which young, compact radio sources are more efficient radio emitters than large extended objects, at a constant accretion power.\\\\ 3) Taking the X-ray luminosity of radio sources as a measure of their ionisation power, we find that GPS/CSO sources are significantly underluminous in their [OIII]$_{5007\\AA}$ line luminosity, including a weak trend with age. This is consistent with the fact that the Str\\\"omgren sphere should still be expanding in these young objects. If true, this would mean that here we are witnessing the birth of the narrow line region of radio-loud AGN. ",
        "introduction": "Ever since their discovery, it has been speculated that those compact radio sources that show convex-shaped radio spectra at cm wavelengths, may be young objects \\citep{shklovsky65,blake70}. As a class, they were named Gigahertz Peaked Spectrum (GPS) sources after their characteristic radio spectrum \\citep[see][for a review]{odea98}, most likely caused by synchrotron self absorption \\citep[e.g.]{fanti90,snellen00a}. High resolution Very Long Baseline Interferometry (VLBI) observations have shown that these sources are typically up to a few hundred parsec in size, often exhibiting jet and/or lobe structures on two opposite sides from their central core $-$ the reason why they are also called Compact Symmetric Objects \\citep[CSO; eg.][]{Wilkinson94}. The most compelling evidence that GPS/CSO are indeed young radio sources comes from VLBI monitoring observations, showing that the bright archetypal objects in this class have hot-spot propagation velocities of $\\sim$0.1-0.2c \\citep{owsianik98a,owsianik98b,tschager00}, indicating kinematic ages of $\\sim$10$^{3}$ years. This is in contrast to speculations that GPS/CSO galaxies are small due to confinement by a particularly dense and clumpy  interstellar medium (ISM) that impedes the outward propagation of the jets \\citep{vanbreugel84,odea91}.  Note that a large fraction of GPS sources, in particular those showing a convex radio spectrum peaking at higher frequencies than a few GHz, turn out to be identified with high redshift quasars (z$\\sim$2-3). Their connection with the population of GPS/CSO galaxies at lower redshift is not clear, and they may well be a completely separate class of object which just also happen to exhibit a convex shaped spectrum \\citep{snellen99}. By no means it has been established that the GPS quasars may also represent a young stage of radio source evolution.  Here we present \\xmm\\ X-ray observations of a small but complete sample of all five GPS sources from the \\citet{pearson88} sample with Galactic latitude $|b| > 20$\\degr\\ (Table~\\ref{tab-sample}). These are among the brightest GPS/CSO galaxies in the sky. All these sources appear to be young radio loud AGN with four out of five sources having measured kinematic ages of their hot spots, indicating ages of up to $\\sim3000$~yr. No kinematic age measurement for B1358+624 exists, but based on its size and the age-size measurements of GPS/CSO sources \\citep{polatidis03} we estimate its age to be 2400 yr. Although the sample is limited in size, it allows us to study the correlations between their X-ray, optical and radio properties, and to assess how they compare to those of mature radio galaxies. One expects that the X-ray luminosity is largely a manifestation of the instantaneous accretion power of the central black hole, whereas the radio luminosity is expected to evolve substantially over the life time of the source \\citep[e.g.]{readhead96,snellen00a}. Moreover, X-ray absorption measurements are an excellent means to probe the circum nuclear density of the GPS galaxies.   \\begin{figure*} \\hbox{\\psfig{figure=b0108_388_ldata_delchi.ps,angle=-90,width=0.5\\textwidth} \\psfig{figure=b0710_439_ldata_delchi.ps,angle=-90,width=0.5\\textwidth} } \\hbox{\\psfig{figure=b1031_567_ldata_delchi.ps,angle=-90,width=0.5\\textwidth} \\psfig{figure=b1358_624_ldata_delchi.ps,angle=-90,width=0.5\\textwidth} } \\hbox{\\psfig{figure=b2352_495_ldata_delchi.ps,angle=-90,width=0.5\\textwidth} \\hskip 0.5\\textwidth }  \\caption{Observed \\xmm\\ count rate spectra. The MOS spectra are shown in red and the PN spectra in black. \\label{spectra1}} \\end{figure*}   Several papers on X-ray observations of GPS sources have appeared in the literature, but one should be cautious to interpret these results in terms of X-ray properties of young radio-loud AGN. A first success was obtained by \\citep{odea00} with ASCA. Although they did not detect B2352+495, they obtained a firm detection of GPS galaxy B1345+125 (PKS1345+125). Furthermore, \\citet{guainazzi04} detected B1404+288 (Mkn 668).\\footnote{Prior to the submission of this publication we learned that Guainazzi et al. have detected a number of other GPS/CSO galaxies in X-rays, which do not overlap with our sample \\citep{guainazzi05}} Both are low redshift GPS galaxies exhibiting strong optical line emission and are powerful infrared emitters, and may be not representative to the class of young radio-loud AGN (yet no reliable ages have been measured for these sources). The combined X-ray and radio observations of the GPS {\\em quasar} B0738+313 presented by \\citet{siemiginowska03b} show that, with its prominent kpc-scale X-ray/radio jet, it is certainly not a young radio-loud AGN.  In order to easily compare our results for GPS/CSO galaxies with the X-ray properties of a large sample of AGN by \\citet{sambruna99} we adapt here a cosmology with $H_0 = 75$~km s$^{-1}$ Mpc$^{-1}$ and $q_0 = 0.5$.  \\begin{table} \\centering \\caption{Log of the \\xmm\\ observations. The MOS exposure time and event rates refer to the average value for MOS1 and MOS2.\\label{tab-obs}} {% \\begin{tabular}{@{}lcccc@{}}\\hline Source  & Observation ID & Start Date & Exposures & Event rates \\\\ &         &                   & (MOS/PN) & (MOS/PN)\\\\ &         & d/m/y             & ks & ct s$^{-1}$\\\\ \\hline B0108+388 & 0202520101  & 09/01/2004 & 16.4/12.0 & 1.6/13.6\\\\ B0710+439 & 0202520201  & 22/01/2004 & 14.2/11.2 & 3.2/26.6\\\\ B1031+567 & 0202520301  & 21/10/2004 & 22.2/12.8 & 34.4/93.0\\\\ B1358+624 & 0202520401  & 14/04/2004 & 12.4/12.0 & 23.2/120.6\\\\ B2352+495 & 0202520501  & 25/12/2004 & 15.8/12.8 & 3.7/28.5\\\\ \\hline \\end{tabular} } NOTE -- The event rates refer to the total detector count rates, not the source count rates. \\end{table}  \\begin{table*} \\begin{center} \\begin{minipage}{\\textwidth} \\caption{Observational properties and parameters obtained by modeling the observed X-ray spectra in the range 0.5-10~keV.\\label{tab-res}} \\begin{tabular}{@{}llllll@{}}\\hline\\noalign{\\smallskip} & B0108+388 & B0710+439    & B1031+567   & B1358+624 &B2352+495 \\\\ \\noalign{\\smallskip}\\hline \\noalign{\\smallskip} PN     0.5-10~keV source count rate ($10^{-3}$~cts\\,s$^{-1}$) & $4.7\\pm0.9$   & $89.5\\pm3.0$ &$12.6\\pm2.0$ & $44.3\\pm2.7$  & $8.6\\pm1.2$\\\\ MOS1+2 0.5-10~keV source count rate ($10^{-3}$~cts\\,s$^{-1}$) & $0.71\\pm0.25$ & $33.0\\pm1.1$ & $4.3\\pm0.5$ & $15.8\\pm0.9$   & $3.4\\pm0.4$\\\\ \\noalign{\\smallskip} Galactic \\nh\\ ($10^{20}$cm$^{-2}$)                 & 5.80        & 8.11         & 0.56        & 1.96      &12.4\\\\ \\noalign{\\smallskip} Normalization \\footnote{Statistical errors correspond to $\\Delta C=1$ (68\\% confidence limits).} ($10^{-5}$ph s$^{-1}$keV$^{-1}$ cm$^{-2}$@ 1 keV) & $3.3\\pm1.1$\\footnote{ The 3$\\sigma$ lower limit is $1.1\\times10^{-5}$ph s$^{-1}$keV$^{-1}$ cm$^{-2}$. The source is detected at the $7\\sigma$ level.} &$8.1\\pm0.6$ & $1.2\\pm0.2$ & $6.1\\pm1.6$ & $1.1\\pm0.2$\\\\ Power law slope \\footnote{Brackets indicate that the power law slope was fixed to this value.} ($\\Gamma$)                           & (1.75)    & $1.59\\pm0.06$ & (1.75)    & $1.24\\pm0.17$ & (1.75) \\\\ Intrinsic \\nh\\ ($10^{22}$cm$^{-2} $) & $57\\pm20$\\footnote{ The 3$\\sigma$ lower limit is $1.8\\times10^{23}$~cm$^{-2}$.} & $0.44\\pm 0.08$ & $0.50\\pm0.18$     & $3.0\\pm0.7$ & $0.66\\pm0.27$\\\\ \\noalign{\\smallskip} Flux (2-10 keV) \\footnote{Including absorption.} ($10^{-13}$~erg\\, s$^{-1}$cm$^{-2}$) & 0.50    &  4.0 & 0.51 & 4.8 & 0.41\\\\ Luminosity (2-10 keV)\\footnote{Calculated for rest frame energies, ignoring absorption.} ($10^{44}$~erg\\, s$^{-1}$)         & 1.18     & 2.16           &  0.22        &   1.67    & 0.046\\\\ C-statistic/bins                     &137.5/99  & 497.4/401     & 104.2/97 &    122.9/99  & 126.3/99 \\\\ \\noalign{\\smallskip} \\hline \\end{tabular} \\end{minipage} \\end{center} \\end{table*}  \\begin{figure*} \\centerline{ \\psfig{figure=gps_xray_radio_corr.ps,width=0.5\\textwidth} \\psfig{figure=gps_xray_oiii_corr.ps,width=0.5\\textwidth} } \\caption{ Left: Correlation between X-ray luminosity and radio luminosity of GPS/CSO sources (filled squares), compared to the radio-loud sample of active galactic nuclei by \\citet{sambruna99}. Right: Idem, for the [O III] line emission. In both figures the dotted lines are the average correlations for the radio loud sample obtained by  \\citet{sambruna99}. The galaxies from the current sample are shown in blue, whereas Mkn 668 and PKS1345+125 are shown in green and red, respectively. \\label{correlations}.} \\end{figure*} ",
        "conclusions": "We have presented \\xmm\\ observations of a sample of all the GPS/CSO galaxies with $|b|>20$\\degr\\ from the \\citet{pearson88} catalog, four of which have measured kinematic time scales for the jet expansion. All five of the sources are detected by \\xmm\\ thereby increasing the number of X-ray detected GPS/CSO galaxies from 2 to 7. These detections allow us to compare the X-ray properties of GPS/CSO galaxies with those of other radio loud AGN.  The results presented here support the hypothesis that GPS/CSO galaxies represent the young phases in the evolution of radio-loud AGN. The alternative explanation that the radio sources are compact due to confinement by an exceptionally high density of the interstellar medium in those  galaxies, seems extremely unlikely in view of the low intrinsic X-ray absorption column densities, which ranges from $\\NH = 4\\times10^{21}$~cm$^{-2}$\\ to $6\\times10^{23}$~cm$^{-2}$, and has a distribution similar to other radio loud AGN. {% After submission of our manuscript, a preprint by \\citet{guainazzi05} arrived at apparently different conclusions based on a sample of five different GPS galaxies observed by \\chandra\\ and \\xmm. Only one of the five GPS galaxies in their sample has \\nh $< 10^{22}$~cm$^{-2}$, whereas this is 75\\%$\\pm$26\\% for their control sample. Note, however, that in our sample three out of five have \\nh $< 10^{22}$~cm$^{-2}$. Taking both samples together, this means that four out of ten GPS galaxies, or 40\\%$\\pm$20\\% have \\nh $< 10^{22}$~cm$^{-2}$\\ consistent with the control sample.  The fact that the absorption columns toward GPS galaxies are consistent with those of other radio galaxies suggests that the interstellar medium densities in the cores of GPS/CSO galaxies are similar to those of other radio loud AGN.} A  similar conclusion was reached by \\citet{pihlstroem03} based on HI radio absorption observations, but the X-ray data provide stronger constraints, as the total column density contributes to the absorption, including ionized regions of the interstellar medium. This may contribute to the fact that in all cases the X-ray column densities are higher than the HI column densities.  A difference between the radio column densities and the X-ray column densities is also that the X-ray column is measured toward the central source, whereas the HI column density is toward the jets, which extend outside the central region. If the difference between radio and X-ray column density is dominated by those geometrical effects, this would be additional evidence against the confinement scenario, since the jets have apparently been able to pierce through the dense local regions that contribute most to the X-ray absorption column.  Although the $\\NH$\\ distribution of our sample cannot be distinguished from other radio-loud AGN, the ratio of radio to X-ray luminosity shows that GPS galaxies have a strong tendency to be relatively radio bright. This supports the view that for the same thrust of the jets younger radio sources are relatively bright \\citep{kaiser97,snellen00a}. However, the data is inconclusive concerning whether GPS-galaxies represent the most radio-luminous phases in the lifes of radio-loud AGN, as would be the case if the source develops in a density profile that drops of as power law with distance \\citep{kaiser97,pihlstroem03}, or whether they are still in their brightening phase. This would be the case if the interstellar medium is best described by a King profile with a relative uniform density within $\\sim$1~kpc of the center and then dropping of as a power law \\citep[e.g.][]{snellen00a}. In the latter case the brightest evolutionary phase of  radio-loud AGN would be represented by the compact steep spectrum sources (CSS), which have  more extended radio emission than GPS/CSO galaxies. A similar study to this one concerning CSS-galaxies can clarify this issue.   Finally, we find that GPS/CSO galaxies are relatively weak in \\oiii\\ line emission. This is again in support of the idea that GPS/CSO galaxies represent the very earliest stages of the evolution of radio-loud AGN, since narrow line regions need time to build up to their equilibrium size, and young narrow line regions are therefore not as bright as fully developed ones. The narrow line region is powered by the UV flux of the central source, but also shocks induced by the expanding jets are likely to contribute to their formation. For those very young radio-loud AGN we advocate here that their emission line regions are powered by the UV flux from the central sources. A case in point is that the relatively weakest \\oiii\\ source, {B0108+388}, has also the highest X-ray column density, which suggests that a dusty torus is partially blocking the UV light. The low \\oiii\\ luminosity therefore shows that the birth of the narrow emission line region must coincide more or less with the birth of the radio jet. {% However, we caution that we only have a limited knowledge of the nature of the ionization mechanism for the [O III] line emission. i.e. both shocks from the jets, as the UV radiation from the AGN may contribute to the ionization. Moreover, compact steep spectrum (CSS) sources, probably representing a more advanced evolutionary state of radio galaxies than GPS galaxies, show that the forbidden line emission tends to be aligned with the radio jet \\citep{devries99}. This phenomenon is not well understood, but it should be accounted for if one wants to build  a more detailed model of the evolution of emission line nebulae in radio galaxies.}  In summary, the findings from this X-ray study lends further support to the theory that GPS/CSO galaxies represent the early phases of radio-loud AGN, in which also the narrow line emission nebula is still in the early phases of its evolution. Future X-ray studies may help to further clarify the relation between age or jet-size, the extent or brightness of the emission line region and the power of the X-ray emission. It is important to include also compact steep spectrum (CSS) sources in such a study, as they are likely to represent the next phase in the evolution of radio-loud AGN. Comparing their X-ray to radio luminosity may help clarify whether the radio emission declines already during the GPS-phase, or first increases, then peaks around the CSS-phase and from then on weakens."
    },
    "0601/astro-ph0601188_arXiv.txt": {
        "abstract": "We present a systematic spectral analysis of 350~bright GRBs observed with BATSE, with high spectral and temporal resolution. Our sample was selected from the complete set of 2704 BATSE GRBs, and included 17 short GRBs. To obtain well-constrained spectral parameters, four different photon models were fitted and the spectral parameters that best represent each spectrum were statistically determined. A thorough analysis was performed on 350~time-integrated and 8459~time-resolved burst spectra. Using the results, we compared time-integrated and time-resolved spectral parameters, and also studied correlations among the parameters and their evolution within each burst. The resulting catalog is the most comprehensive study of spectral properties of GRB prompt emission to date, and provides constraints with exceptional statistics on particle acceleration and emission mechanisms in GRBs. ",
        "introduction": "Previous spectral studies of Gamma-Ray Burst (GRB) prompt emission have shown compelling evidence that a simple emission mechanism cannot entirely account for the observed spectra. However, spectral parameters that provide constraints on GRB emission mechanisms often depend on functional forms used to fit the data, as well as integration timescales of the spectra. A consistent, comprehensive spectral study of GRB prompt emission in large quantity is therefore crucial to unveiling the nature of GRBs.  A total of 2704 GRBs were observed with the Large Area Detectors (LADs) of the Burst and Transient Source Experiment (BATSE) on board the {\\it Compton Gamma-Ray Observatory}, in the energy range of $\\sim$ 30 -- 1900 keV. BATSE provided the largest database currently available of GRBs from a single experiment, with good spectral and temporal resolution. Many of the BATSE GRBs were bright enough for high-time resolution spectroscopy within the bursts as well as time-integrated spectroscopy. We present here a systematic spectral analysis of 350 bright BATSE GRBs. The sample size is more than twice that of the previous BATSE spectral catalog \\cite{pre00}, which included only time-resolved spectroscopy. Our analysis included 350 time-integrated spectra and 8459 time-resolved spectra, and four different photon models were used to fit all the spectra. The resulting catalog is the largest and most comprehensive study to date of GRB prompt emission spectra. ",
        "conclusions": "The GRB spectral database obtained in this work is derived from the most sensitive and largest database currently available. Therefore, these results set a standard for spectral properties of GRB prompt emission with exceptional statistics. Our results can provide reliable constraints for existing and future theoretical models of GRB emission and particle acceleration mechanisms."
    },
    "0601/astro-ph0601377_arXiv.txt": {
        "abstract": "We investigate the extent to which correlated distortions of the luminosity distance-redshift relation due to large-scale bulk flows limit the precision with which cosmological parameters can be measured. In particular, peculiar velocities of type 1a supernovae at low redshifts, $z < 0.2$, may prevent a sufficient calibration of the Hubble diagram necessary to measure the dark energy equation of state to better than 10\\%, and diminish the resolution of the equation of state time-derivative projected for planned surveys. We consider similar distortions of the angular-diameter distance, as well as the Hubble constant. We show that the measurement of correlations in the large-scale bulk flow at low redshifts using these distance indicators may be possible with a cumulative signal-to-noise ratio of order 7 in a survey of 300 type 1a supernovae spread over 20,000 square degrees. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601694_arXiv.txt": {
        "abstract": "We show that a Magellanic Cloud origin for the warp of the Milky Way can explain most quantitative features of the outer HI layer recently identified by Levine, Blitz \\& Heiles (2005). We construct a model similar to that of Weinberg (1998) that produces distortions in the dark matter halo, and we calculate the combined effect of these dark-halo distortions and the direct tidal forcing by the Magellanic Clouds on the disk warp in the linear regime. The interaction of the dark matter halo with the disk and resonances between the orbit of the Clouds and the disk account for the large amplitudes observed for the vertical $m=0,1,2$ harmonics.  The observations lead to six constraints on warp forcing mechanisms and our model reasonably approximates all six.  The disk is shown to be very dynamic, constantly changing its shape as the Clouds proceed along their orbit. We discuss the challenges to MOND placed by the observations. ",
        "introduction": "The warp of the outer Milky Way, known since 1957 \\citep{kerr57}, has been quantitatively determined for the first time by \\citet{levine05}. It can be described as a superposition of three and only three of the lowest order vertical harmonics of a disk: a {\\em dish-shaped} $m=0$, a {\\em integral-sign-shaped} $m=1$, and a {\\em saddle-shaped harmonic} $m=2$.  The lines of nodes for each are close to coincident and nearly radial.  The amplitudes of each reaches 7--10\\% of the radius of the disk.  A number of possible warp producing mechanisms have been suggested including long-lived eigenmodes, forcing by halo triaxiality, persistent cold-gas accretion, and tidal excitation.  We show here that the origin of this warp can be well-described as the tidal interaction of the Magellanic Clouds with the disk and dark matter halo of the Milky Way.  The interaction of the dark matter halo with the disk and resonances between the orbit of the Clouds and the disk account for the large amplitudes of the three harmonics and their approximate shape and orientation. ",
        "conclusions": "\\subsection{Comparison with N-body Simulations}  The predictions of W98 were checked by several groups using N-body simulations. \\citet{garcia-ruiz02} used a hybrid N-body particle--ring code and did not find the warp predicted in W98. They offer the failure of the linear theory as the culprit but did not investigate a variety of models. Similarly, \\citet{mastro05} performed a simulation including both the gaseous, stellar and dark components of the Large Cloud and the Milky Way and remark that the effect on the Milky Way is negligible.  Conversely, \\citet{tsuchiya02}, using a hybrid code that includes both a potential expansion and a tree code, obtained amplitudes of $m = 0,1,2$ warps in the larger halo model that show very good agreement with the observed amplitudes in the Milky Way. He also shows that the amplitude of the warp depends strongly on the dark halo mass model.  It is difficult to reconcile these contradictory findings but three possibilities obviously occur: 1) the linear theory does not apply; 2) the simulations do not apply due to numerical difficulties; and 3) the two simulations with null warps have chosen unlucky sets of parameters.  Both \\citet{garcia-ruiz02} and \\citet{tsuchiya02} choose methods that explicitly treat the multiple scales. \\citet{garcia-ruiz02} used tilted rings to represent the disk and a tree code to represent the halo. Although the use of rings limits the investigation to the $m = 1$ response, it should be sufficiently sensitive to the halo excitation without strong particle number issues.  \\citet{tsuchiya02} used an expansion algorithm to represent the halo gravitational field and a tree code to represent the disk to better represent the multiple scales \\citep[see][for additional discussion of the dynamic range problems that motivate this choice]{tsuchiya02}. We feel that the discrepancy between the results of these two groups is most like Item (3): the Garc\\'ia-Ruiz et al. model is not particularly warp producing.  The good qualitative correspondence between W98 and \\citet{tsuchiya02} suggests that the linear theory captures the underlying physics although is likely to differ in detail. For example, the linear theory over-predicts the height in the edge of the disk (Fig. \\ref{amps}), and this motivates our truncation of the predictions in the outer disk.  \\citet{mastro05} primarily emphasize the affect of the Milky Way tides on the Large Cloud and do not tailor their approach to treat the multiple scale problem per se. We suggest that their report of no disk warp results from Items (2) and (3).  \\subsection{Other Explanations}  Most warp theories depend on the bending response of the disk and this response is strongly affected by the existence of a self-gravitating dark matter halo.  The theories may be roughly grouped as follows: 1) bending modes may be primordially excited and persistent \\citep[e.g.][]{sparke88}; 2) the response of the disk to the non-axisymmetric shape of the halo; 3) tidal excitation (as we have discussed here) and 4) the warp may be produced by the response of the disk to cosmic infall \\citep[e.g.][]{jiang99}. The topic was nicely reviewed by \\citet{binney91}. Subsequently, \\citet{nelson95} argued against long-lived modes.  The triaxiality of the Milky Way halo remains uncertain.  \\citet{helmi04} reports a prolate halo (q=1.25) while \\citet{johnston05} finds an oblate halo (q=0.8-0.9) using Sgr dwarf constraints.  Such modest triaxiality seems unlikely to produce the observed vertical $m=2$ feature. In addition, cosmic infall more naturally produces tilted rings, an $m=1$ feature; the observed vertical $m=2$ may require a conspiracy of several inflow directions.  \\subsection{Relevance for Modified Gravity}  Although the success of our model favors the existence of a dark-matter halo in Nature, many find it seductive to modify gravity to produce the observed rotation velocities in the Galaxy without dark matter \\citep[modified Newtonian dynamics or MOND][]{milgrom83a,milgrom83b,bekenstein84}. Might the tidal theory also apply in MOND? It is beyond the scope of this letter to repeat our calculations using MOND, it seems plausible that direct forcing of the Galactic disk by the Clouds in MOND may provide warp amplitudes in excess of the original predictions without a dark halo (HT). Similarly, we would expect the disk modes and frequencies to be qualitively similar, the excess restoring force of the halo being produced by the MOND force.  However, MOND would have to (1) admit bending modes with similar morphology to those in the Newtonian theory and (2) these modes would have to conspire to frequencies that couple them to the direct forcing by the Clouds in such a way that they assumed the same orientation as in the ``clockwork'' described in \\S5. Because this clockwork and the $m=2$ to $m=1$ amplitude ratio depends on the simultaneous halo and disk excitation, agreement with the observations described in \\S2 seems rather unlikely and thereby disfavors MOND.  We encourage detailed predictions.  \\subsection{Relevance for other Galactic Systems}  A large fraction of other warped galaxies also show warps in their HI layers and a significant fraction of these are asymmetric.  Warps in these systems are generally analyzed with a program such at ROTCUR \\citep{begeman89} or one of its derivatives, but invariably are forced to fit the $m = 1$ warp only.  This paper, and the work of \\citet{levine05} show that at least three harmonics ought to be fit to the warps of external galaxies, especially those with asymmetric warps, which can be caused by a superposition of these harmonics.  The degree to which various harmonics are present in a warp can produce important constraints on whether the warp is due to a satellite, a triaxial halo, cold gas inflow, or some primordial excitation.  All warps may not be alike.  \\subsection{Summary}  We have demonstrated a plausible mechanism for the excitation of the warp that explains all of its general features.  A number of warp producing mechanisms have been explored elsewhere in addition to tides.  It is possible that different mechanisms may act in different galaxies and possibly in concert in a single galaxy. Nonetheless, our simple model nicely reproduces the observed features in the Milky Way HI gas layer.  The existence of massive companions, the Magellanic Clouds, and the prediction from linear perturbation theory and at least one corroborating N-body simulation suggests that a tidal explanation is viable.  Our model depends on the gravitational response of the halo and thereby suggests that the dark matter is not an artifact of modifying the laws of gravity. Conversely, given a dark halo, we argue that the tide from the Magellanic Clouds must be affecting the Milky Way disk, and, given the quality of the agreement, it seems to be the dominant mechanism.  This model then promises an additional constraint on the distribution of dark matter. Although we have emphasized the gas layer response beyond the stellar disk, the effect of other satellite encounters such as the recent accretion of Sagittarius dwarf may be detectable in future high-resolution surveys and may help determine the properties of the inner halo.  Warp observations are important because promise to reveal aspects of the dark matter distribution that are otherwise observationally inaccessible.  Further analysis of warps may provide more precise constraints on the profile of dark matter in the Milky Way and nearby external galaxies."
    },
    "0601/astro-ph0601686.txt": {
        "abstract": "We have made an extensive study of the UV spectrum of $\\eta$ Carinae, and find that we do not directly observe the star and its wind in the UV. Because of dust along our line of sight, the UV light that we observe arises from bound-bound scattering at large impact parameters (e.g., 0.033\\arcsec). We obtain a reasonable fit to the UV spectrum by using only the flux that originates outside 0.033\\arcsec. This explains why we can still observe \\etastar\\ in the UV despite the large optical extinction --- it is due to the presence of an intrinsic coronagraph in the $\\eta$ Carinae system, and to the extension of the UV emitting region. It is not due to peculiar dust properties alone. We have computed the spectrum of the purported companion star, and show that it could only be directly detected in the UV spectrum preferentially in the Far Ultraviolet Spectroscopic Explorer (FUSE) spectral region (912-1175\\,\\AA). However, we find no direct evidence for a companion star, with the properties indicated by X-ray studies and studies of the Weigelt blobs, in UV spectra. This might be due to reprocessing of the companion's light by the dense stellar wind of the primary. FUSE observations, at epochs when (and if) the opening of the bow shock is along our sightline, should have a better chance of detecting the companion spectrum. The UV spectrum is dominated by low ionization lines, many exhibiting P~Cygni profiles. Some of the strongest lines that can be readily identified include \\ion{C}{2} $\\lambda$ 1335 (UV1), \\ion{Si}{2}  $\\lambda\\lambda$1304, 1309 (UV3); $\\lambda$ 1264 (UV4), $\\lambda\\lambda$ 1527, 1533 (UV 2); $\\lambda\\lambda$  1808, 1817 (UV1), \\ion{S}{2} $\\lambda\\lambda$ 1250, 1253 (UV1), \\ion{Al}{2} $\\lambda$1671, \\ion{N}{1}\\ $\\lambda\\lambda$ 1493, 1495 (UV4), \\ion{Mg}{2} $\\lambda\\lambda$ 2796, 2803, as well as numerous \\ion{Fe}{2} lines. Higher excitation lines due to \\ion{Al}{3} $\\lambda\\lambda$1855, 1863 and \\ion{Si}{4} $\\lambda\\lambda$1394, 1403 can be identified.  A previous identification of \\ion{C}{4} $\\lambda\\lambda 1548, 1552$ must, because of severe blending, be considered as uncertain. The terminal velocity, as derived from numerous emission lines, is less than 600\\,\\kms, with preferred values around 520\\,\\kms. This value is consistent with that seen in optical spectra. Broad \\ion{Fe}{2} and [\\ion{Fe}{2}] emission lines are detected in spectra taken in the SE lobe, 0.2\\arcsec\\ from the central star. These lines arise in the stellar wind --- thus the Space Telescope Imaging Spectrograph (STIS) and the Hubble Space Telescope (HST) are resolving, at some wavelengths, the stellar wind of Eta Carinae.  The wind spectrum shows some similarities to the spectra of the B \\& D Weigelt blobs, but also shows some marked differences in that high excitation lines, and lines pumped by Ly$\\alpha$, are not seen. The resolution of the stellar wind at optical wavelengths, and the detection of the broad lines, lends support to our interpretation of the UV spectrum,  and to our model for $\\eta$ Carinae. ",
        "introduction": "Eta Carinae is one of the most luminous and spectacular stars in our galaxy, and exhibits variability on a wide range of time scales (Zanella, Wolf, \\& Stahl \\citeyear{ZWS84_shell}, Davidson and Humphreys \\cite{DH97_rev}). A major breakthrough to understanding some of the variability was made by Damineli \\cite{Dam96_cyc} who found that nebular line strengths varied periodically on a time-scale of 5.54 years. A similar time-scale was also present in infrared data (Whitelock \\etal\\ \\citeyear{WFK94_eta_var}). The cycle was confirmed to be periodic when the event predicted for December 1997 came on schedule (Damineli \\etal\\ \\citeyear{DKW00_bin}, Feast \\etal\\ \\citeyear{FWM01_IRvar}). Variability seen in radio data can also be related to a similar time-scale (Duncan \\etal\\ \\citeyear{DWRL99_eta_radio}, Duncan \\& White \\citeyear{DW03_eta_radio}).  On the basis of radial velocity variations, Damineli, Conti \\& Lopes (\\citeyear{DCL97_eta_bin}) postulated that Eta Carinae is a binary system characterized by high eccentricity, a hotter companion, and strong colliding winds. X-ray observations, which reveal an apparent X-ray eclipse, appear to confirm the binary hypothesis (Ishibashi \\etal\\ \\citeyear{ICD99_xray}, Corcoran \\etal\\ \\citeyear{CIS01_xray}). The duration ($\\sim3$ months) of the eclipse, together with the rapid variability of high excitation nebular lines, indicates that the orbit is highly elliptical, with an eccentricity greater than 0.8 (see, e.g., Corcoran \\citeyear{CIS01_xray}).  However, fitting the long duration of the X-ray minimum has been a challenge that was tentatively circumvented by enhanced mass-loss at periastron (Corcoran et al. \\citeyear{CIS01_xray}) or with a tilted angle of the colliding wind shock cone (Pittard \\& Corcoran \\citeyear{PC02_xray}, Ishibashi \\citeyear{Ish01_orbit}). In addition, radiative transfer effects and the complicated and severely blended profiles observed from the ground make it difficult to measure, and to interpret, radial velocity measurements. These difficulties led Davidson \\cite{Dav99_bin} to question the validity of the derived orbital parameters. Indeed, high spatial resolution observations with the Space Telescope Imaging Spectrograph (STIS) did not reveal the expected velocity shifts in emission lines associated with the primary star (Davidson \\etal\\ \\citeyear{DIG00_bin}). An alternative set of orbital parameters, based on analysis of the X-ray light curve, has been given by Ishibashi \\cite{Ish01_orbit}. Thus, while the binary model is generally accepted, the nature of the companion star, its orbit, and the influence of the companion on the major outbursts of the 1840s and 1890s are uncertain. Other interpretations of the variability have been suggested, most notably shell ejections (e.g., Davidson \\etal\\ \\citeyear{DMH05_Ha}, Martin \\etal\\ \\citeyear{MDH05_4686}). A combination of shell ejections and binarity might be needed to explain variability seen in {\\it Hubble Space Telescope (HST)} data (e.g., Smith \\etal\\ \\citeyear{SDG03_eta_lat}).  %Since the companion star has not been directly observed, Information on the nature of the companion comes primarily from indirect arguments. Analysis of X-ray data indicates that the companion should have a mass loss rate of approximately $\\Mdot=1.0 \\times 10^{-5}$\\,\\Msunyr, while estimates of the terminal velocity range from 1700\\,\\kms\\ (Corcoran \\etal\\ \\citeyear{CIS01_xray}) to 3000\\,\\kms\\ (Pittard \\& Corcoran \\citeyear{PC03_xray}), with the higher velocity estimates now preferred. Other constraints on the companion come from its influence on the spectrum of the Weigelt blobs (Verner \\etal\\ \\citeyear{VGB02_bd_blob}; Verner, Bruhweiler \\& Gull \\citeyear{VBG05_bd_blob}). These suggest that it is an O type star with an effective temperature between 34,000\\,K and 38,000\\,K.  With the aim of clarifying the nature of the 5.5 year periodicity, and to gain insights into both the nature of the primary and companion stars, we initiated a major observational multi-wavelength campaign, using the Hubble Space Telescope (HST Eta Carinae Treasury project; PI K. Davidson), FUSE (PI T. Gull), X-ray satellite observatories (PI. M. Corcoran), and numerous ground-based observatories (e.g., VLT-UVES; PI. K. Weis), to observe Eta Carinae through a full variability cycle. Some of the HST data used in this paper are based on data obtained as part of that campaign, while other data were obtained as part of earlier HST programs to understand Eta Carinae. The campaign has confirmed the 5.5 year periodicity (Whitelock \\etal\\ \\citeyear{WFM04_event}, Corcoran \\citeyear{Cor05_xray_mon}, White \\etal\\ \\citeyear{WDC05_radio}) and provided a wealth of new data to help understand the complex system that is Eta Carinae.  An introduction to the extensive literature on Eta Carinae can be obtained from the reviews by Humphreys \\& Davidson \\cite{HD94_LBV_rev}, and Davidson \\& Humphreys \\cite{DH97_rev}, and three relatively recent workshops devoted to Eta Carinae and related objects (Morse, Humphreys, \\& Damineli \\citeyear{Eta_wrk_99}; Gull, Johansson, \\& Davidson \\citeyear{Hven_wrk_shop}; Humphreys \\& Stanek \\citeyear{GT05}).  \\subsection{The Primary Star: \\etastar}  Ground-based spectra of Eta Carinae reveal a complex spectrum of H, \\ion{He}{1}, \\ion{Fe}{2}, and [\\ion{Fe}{2}] emission lines with 2 principal components. There is a narrow nebular-like spectrum $(V_{\\hbox{FWHM}} < 40\\,\\kms)$, and a broad spectrum which indicates gaseous outflows with a velocity of approximately 500\\,\\kms. As other components are also seen, it is difficult to discern the underlying nature of the primary star, \\etastar. Using the reflected spectrum Hillier \\& Allen \\cite{HA92_eta} suggested that the spectrum of \\etastar\\ is similar to the extreme P~Cygni star HDE\\,316285. Later observations with the HST confirm this suggestion (Davidson \\etal\\ \\citeyear{DEW95_FOS}; Hillier \\etal\\ \\citeyear{HDIG01} [hereafter HDIG]). With the STIS on the HST it is possible to obtain the spectrum of the central star, uncontaminated by lines from the adjacent Weigelt blobs and the Homunculus. The Weigelt blobs are spatially unresolved condensations first seen in ground based speckle studies of $\\eta$ Carinae (Weigelt, \\& Ebersberger \\citeyear{WE86_blobs}; Hofmann \\& Weigelt \\citeyear{HW88_blobs}). It is the Weigelt blobs that give rise to the narrow nebular spectrum (Davidson \\etal\\ \\citeyear{DEW95_FOS}, \\citeyear{DEJ97_UV_blobs}). The Little Homunculus also contributes to the narrow nebular spectrum (Ishibashi \\etal\\ \\citeyear{Ish03_lit_hom}; Smith \\citeyear{Smi05_lit-hom}).  HDIG were able to model the March 1998 optical spectrum, taken a few months after the 1997 spectroscopic event (the X-ray minimum began in mid-December of 1997; see, e.g., Corcoran \\citeyear{Cor05_xray_mon}) using CMFGEN -- a non-LTE line blanketed code designed for modeling stars with stellar winds (Hillier \\& Miller \\citeyear{HM98_blank}, \\citeyear{HM99_WC}). The code assumes spherical symmetry, and that the star and its wind is spatially unresolved. They were able to obtain a good fit to the H, \\ion{He}{1}, and \\ion{Fe}{2}\\ emission line spectrum using a luminosity of $5 \\times 10^6$\\,\\Lsun\\ for the primary star, and assuming a mass-loss rate of $1.0\\times 10^{-3}$\\,\\Msunyr\\ and solar mass-fraction of iron. The best fit model confirmed that the central star suffered, in March 1998, severe circumstellar extinction with a total visual extinction of 7 magnitudes. The primary difference between the March 1998 spectrum, and later spectra, is that the P~Cygni profiles tend to be stronger, and more prevalent in the March 1998 data set. In the optical, most emission line strengths (i.e, EWs) are very similar to those obtained on 4-Jul-2002 --- indeed the two data sets are in better agreement with each other than with the CMFGEN model. The major exceptions are the \\ion{He}{1}\\ profiles which show significant profile variations. Surprisingly, there is no strong indication that the terminal velocity of the outflow has changed. Changes in the H$\\alpha$ profile with time have been discussed by Davidson \\etal\\, \\cite{DMH05_Ha}, while other changes will be the subject of future papers.  While the fit to the emission line spectrum was satisfactory (but see HDIG for more details), there were two fundamental discrepancies between the best fit model and the observations. First, the models predicted much stronger absorption components associated with  \\ion{H}{1}, \\ion{He}{1}\\ and \\ion{Fe}{2} P~Cygni emission lines than was observed. Second, the fit to the spectrum shortward of $1600\\,$\\AA\\ was very poor -- emission features did not usually match and the UV spectrum was much less absorbed, by the wind, than predicted by the model. Since the emission lines sample the whole wind, but the wind absorption components sample only one line of sight, two possible causes were suggested. First, the discrepancies could arise because Eta's wind is asymmetric. This would not be surprising since the Homunculus is bipolar, and other LBVs, such as AG~Car, are known to possess asymmetric winds (e.g., Schulte-Ladbeck \\citeyear{SCH94_AG}). For this to work, the wind would have to be more ionized along our sightline. Recent observations of the reflected optical spectrum by Smith \\etal\\ \\cite{SDG03_eta_lat} provide direct evidence that the wind is aspherical and probably axisymmetric (bipolar), with stronger Balmer absorption near the poles.  Interferometric observations by van Boekel \\etal\\ \\cite{Boek03}, confirm the bipolar geometry of the wind. Second, the ionization of the wind in some regions could be influenced by the ionizing radiation field of the companion. This is appealing since it provides a simple explanation for some of the observed spectral changes.  In this paper we reexamine the formation of the UV spectrum. We identify major wind lines seen in UV spectra, discuss the terminal velocity of the wind, and estimate the intervening H column density. Extensive foreground absorptions, caused by both circumstellar and interstellar matter, strongly influence the UV spectrum and have been discussed elsewhere (Gull \\etal\\ \\citeyear{GVB05_abs}; Nielsen, Gull, \\& Vieira Kober, \\citeyear{NGK05}). We investigate the reason for our model's failure to explain the observed UV spectrum. Several alternatives are considered including wind asymmetries, ionization of the outer wind by a companion star, the direct influence of a companion spectrum, the influence of dust, and the spatial extension of the UV emitting region. We show that the original model, when extended and interpreted differently, can explain many of the features seen in the UV spectrum.  The paper is organized as follows: In \\S \\ref{Sec_obs} we discuss the observations and data reduction while in \\S \\ref{Sec_dist_red} we examine the distance, visual magnitude and reddening of \\etastar, the primary star associated with $\\eta$ Carinae. Sections \\ref{Sec_nat_UV}, \\ref{Sec_line_id} and \\ref{Sec_vinf} discuss the nature of the UV spectrum, the identification of UV wind lines, and the terminal velocity of the wind. The spectrum of the companion star, and its possible influence of the observed spectrum is discussed in \\S \\ref{Sec_companion_star}. In \\S \\ref{Sec_fuse} we examine the FUSE spectrum, while the \\ion{H}{1}\\ column density along our sight line is derived in \\S \\ref{Sec_HI_col}. The Model used for the analysis, and improvements made to it for the UV analyses, are described in \\S \\ref{Sec_model}, while the possible importance of wind asymmetries and flow times are discussed in \\S \\ref{Sec_asym}. The creation of the UV spectrum is discussed in \\S \\ref{Sec_UV_mystery}. Finally, in \\S \\ref{Sec_out_wind_spec} we discuss how STIS observations allow us to study the spatial structure of the outer wind at optical wavelengths. ",
        "conclusions": "\\label{Sec_conc}  A fundamental prerequisite to understanding the reason for the ejection of the Homunculus is to determine the fundamental properties of the primary star. Unfortunately the primary star is shrouded in a dense wind. In addition, dust extinction does not allow an uninterrupted view of the primary star, and its wind.  The UV spectrum of Eta Carinae is dominated by low ionization lines, many exhibiting P Cygni profiles. Some of the strongest lines that can be readily identified include \\ion{C}{2} $\\lambda$ 1335 (UV1), \\ion{Si}{2}  $\\lambda\\lambda$1304, 1309 (UV3); $\\lambda$ 1264 (UV4), $\\lambda\\lambda$ 1527, 1533 (UV 2); $\\lambda\\lambda$  1808, 1817 (UV1), \\ion{S}{2} $\\lambda\\lambda$ 1250, 1253 (UV1), \\ion{Al}{2} $\\lambda$1671, \\ion{N}{1}\\ $\\lambda\\lambda$ 1493, 1495 (UV4), \\ion{Mg}{2} $\\lambda\\lambda$ 2796, 2803, as well as numerous \\ion{Fe}{2} lines. Higher excitation lines due to \\ion{Al}{3} $\\lambda\\lambda$1855, 1863, and probably \\ion{Si}{4} $\\lambda\\lambda$1394, 1403 can also be identified. The identification of \\ion{C}{4}\\ $\\lambda\\lambda$ 1548, 1552 is uncertain because of severe line blending.  We have shown that we do not directly observe the star and its wind in the UV --- rather, because the inner regions are occulted by dust, we only observe the spectrum created in the outer regions (i.e., at large impact parameters) of the stellar wind. This helps partially explain the flatness of the UV extinction law --- the flatness does not simply reflect the properties of the dust. Importantly, it indicates that we cannot simply use the observed UV spectrum to determine the properties of the dust causing the circumstellar extinction. Further the dust in the densest knots will suffer greater shielding from UV radiation, and will have different properties compared to the dust located in other regions. This affects our ability to understand the excitation of the Weigelt blobs. The results reinforce the belief that our view of Eta Carinae is biased. The view from other directions would be significantly different. It is still probable the ionizing flux from the companion star is also influencing the observed spectrum. The high hydrogen column density towards the central star ($\\log \\, $N(H)$=22.5$) is confirmed from analysis of the Ly$\\alpha$ profile in HST MAMA data.  We show that the FUSE spectral region is, in principle, the best wavelength region to detect directly the primary star. However, no direct evidence of the companion star, with the properties indicated by X-ray studies (i.e, $\\Mdot \\sim 10^{-5}\\,\\Msunyr$, and $\\Vinf=3000\\,$\\kms) is seen in current FUSE data, or in MAMA data. This may be partially a consequence of reprocessing of the companion light by the dense wind of the primary. Alternatively, it may indicate the parameters of the O star, as inferred from the X-ray and Weigelt blob analyses, are incorrect.  The best fit to the FUSE data is obtained using the spectrum for the model of the primary originating beyond 0.033\\arcsec. The fit, not surprisingly, is far from perfect. As our study has shown modeling of the UV spectra of Eta Car is extremely difficult. It is necessary to allow for the extended nature of the UV emitting region and occulation by dust. Further, the companion will modify the ionization structure of the wind, and while the companion is probably the dominant light source, in the FUSE spectral region, its spectrum will be modified by the dense wind of the primary. Finally, there is extensive evidence for an axi-symmetric wind  which also needs to be allowed for in future modeling.  %and Ly$\\beta$ %profiles profiles in HST MAMA data, and FUSE data.  The terminal velocity of the primary's wind lies between 500 and 600\\,\\kms\\ with values near the lower end preferred. This range is consistent with that determined by HDIG from analysis of the optical spectrum. Surprisingly, we find no convincing evidence for higher velocity components.  With the STIS on the HST we have resolved the stellar wind of Eta Carinae. Broad \\ion{Fe}{2}\\ emission lines are observed directly. These broad wind lines can be seen at a distance of 0.2\\arcsec\\ (and beyond) from the central source, and also indicate a wind terminal velocity of approximately 500\\,\\kms. The wind spectrum shows some similarities to the spectra of B \\& D Weigelt blobs, but also shows some marked differences in that high excitation lines, and lines pumped by Ly$\\alpha$, are not seen.  \\blankline The observations were made with the NASA/ESA Hubble Space Telescope under HST-GO and STIS-GTO programs through the STScI under NAS5-26555 and with the NASA/CNES/CSA Far Ultraviolet Spectroscopic Explorer, which is operated for NASA by The Johns Hopkins University under NASA contract NAS5-32985.  T.R.G., G.S., R.C.I. \\& D.J.H. acknowledge support from the FUSE Guest Investigator program. The HST Treasury project is supported by NASA programs GO-9420 and GO-9973. K. W. acknowledges support by the state of North Rhine-Westphalia (Lise Meitner fellowship). We would also like to thank K. Davidson and R. M. Humphreys for providing useful comments on an earlier draft of the manuscript.  \\def\\SA_BCM#1{in  NATO ASI Seri. C 341, Stellar Atmospheres: Beyond Classical Models, ed. L. Crivellari, I. Hubeny \\& D. G. Hummer (Dordrecht: Kluwer), #1}  \\def\\ASP#1{in ASP Conf. Ser. 22, Nonisotropic and Variable Outflows from Stars, ed. L. Drissen, C. Leitherer, \\& A. Nota (San Francisco: ASP) p.~#1}  %1988 \\def\\prco#1{in Polarized Radiation of Circumstellar Origin, ed. G. V. Coyne, A. M. Magalh\\~aes, A. F. J. Moffat, R. E. Schulte-Ladbeck, \\& S. Tapia (Vatican Observatory),~#1}  %Physics of luminous blue variables. \\def\\plbv#1{in IAU Colloq. 113, Physics of Luminous Blue Variables, ed. K. Davidson, A. F. J. Moffat, \\& H. J. G. L. M. Lamers (Dordrecht: Kluwer),~#1}  \\def\\lbv#1{1997, in ASP Conf. Ser. 120, Luminous Blue Variables: Massive Stars in Transition, ed. A. Nota, \\& H.J.G.L.M. Lamers (San Francisco: ASP),~#1}  \\def\\IAU70#1{1976, in IAU Symp. 70, Be and Shell Stars, ed. A. Slettebak (Dordrecht: Reidel)~#1}  \\def\\etaworkshop#1{in ASP Conf. Ser. 179, Eta Carinae at the Millenium, ed. J.A. Morse, R.M. Humphreys, \\& A. Damineli (San Francisco: ASP),~#1}  \\def\\Hveneta#1{in ASP Conf. Ser. 242, Eta Carinae and Other Mysterious Stars: The Hidden Opportunities of Emission Line Spectroscopy, ed. T.R. Gull, S. Johansson, \\& K. Davidson (San Francisco: ASP),~#1}"
    },
    "0601/nucl-th0601019_arXiv.txt": {
        "abstract": "Nuclear structure of $^7$Be, $^8$B and $^{7,8}$Li is studied within the {\\it ab initio} no-core shell model (NCSM). Starting from high-precision nucleon-nucleon (NN) interactions, wave functions of $^7$Be and $^8$B bound states are obtained in basis spaces up to $10\\hbar\\Omega$ and used to calculate channel cluster form factors (overlap integrals) of the $^8$B ground state with $^7$Be+p. Due to the use of the harmonic oscillator (HO) basis, the overlap integrals have incorrect asymptotic properties. We fix this problem in two alternative ways. First, by a Woods-Saxon (WS) potential solution fit to the interior of the NCSM overlap integrals. Second, by a direct matching with the Whittaker function. The corrected overlap integrals are then used for the $^7$Be(p,$\\gamma$)$^8$B S-factor calculation. We study the convergence of the S-factor with respect to the NCSM HO frequency and the model space size. Our S-factor results are in agreement with recent direct measurement data. We also test the spectroscopic factors and the corrected overlap integrals from the NCSM in describing the momentum distributions in knockout reactions with $^8$B projectiles. A good agreement with the available experimental data is also found, attesting the overall consistency of the calculations. ",
        "introduction": "Introduction} The $^7$Be(p,$\\gamma$)$^8$B capture reaction serves as an important input for understanding the solar neutrino flux \\cite{Adelberger}. Recent experiments determined the neutrino flux emitted from $^8$B with a precision of ~9\\% \\cite{SNO}. On the other hand, theoretical predictions have uncertainties of the order of 20\\% \\cite{CTK03,BP04}. The theoretical neutrino flux depends on the $^7$Be(p,$\\gamma$)$^8$B S-factor that needs to be known with high precision. Many experimental and theoretical investigations were devoted to this reaction. Experiments were performed using direct measurement techniques with proton beam and $^7$Be targets \\cite{Filippone,Baby,Seattle} as well as by indirect methods when a $^8$B beam breaks up into $^7$Be and proton in the Coulomb field of a heavy target \\cite{BBR86,Be7pgamm_exp}. Theoretical calculations needed to extrapolate the measured S-factor to the astrophysically relevant Gamow energy were performed with several methods: the R-matrix parametrization \\cite{Barker95}, the potential model \\cite{Robertson,Typel97,Davids03}, and the microscopic cluster models \\cite{DB94,Csoto95,D04}.  In this work, we present the first calculation of the $^7$Be(p,$\\gamma$)$^8$B S-factor starting from the {\\it ab initio} wave functions of $^8$B and $^7$Be. We apply the {\\it ab initio} no-core shell model (NCSM) \\cite{NCSMC12,NKB00}. In this method, one considers nucleons interacting by high-precision nucleon-nucleon (NN) potentials. There are no adjustable or fitted parameters. Within the NCSM, we study the binding energies and other nuclear structure properties of $^7$Be, $^8$B as well as $^{7,8}$Li, and calculate overlap integrals for the $^8$B and $^7$Be bound states. Due to the use of the harmonic-oscillator (HO) basis, we have to correct the asymptotic behavior of the NCSM overlap integrals. This is done in two alternative ways. First, by fitting Woods-Saxon (WS) potential solutions to the interior part of the NCSM overlap integrals under the constraint that the experimental $^7$Be+p threshold is reproduced. Second, by a direct matching of the NCSM overlap integrals and the Whittaker function. The corrected overlap integrals are then utilized to calculate the $^7$Be(p,$\\gamma$)$^8$B S-factor as well as momentum distributions in stripping reactions. We pay special attention to the convergence of the S-factor with respect to the NCSM model space size and the HO frequency.  In Sect.~\\ref{sec:spectroscopy}, we present the NCSM results for $^7$Be, $^8$B and $^{7,8}$Li energies, ground-state radii and electromagnetic moments and transitions. The calculation of cluster form factors and the correction of their asymptotics is described in Sect.~\\ref{sec:overlap}. A test of the corrected NCSM overlaps and spectroscopic factors for the momentum distributions in the stripping reaction $(^7{\\rm Be}+p)+A\\longrightarrow \\ ^7{\\rm Be}+X$ is discussed in Sect.~\\ref{sec:momdis}. The S-factor calculation with sensitivity and convergence studies is presented in Sect.~\\ref{sec:sfactor}. Conclusions are drawn in Sect.~\\ref{sec:conc}. ",
        "conclusions": ""
    },
    "0601/astro-ph0601627_arXiv.txt": {
        "abstract": "We present the catalog of X-ray sources detected in a shallow \\chandra\\ survey of the inner 2\\degree$\\times$0.8\\degree\\ of the Galaxy, and in two deeper observations of the Radio Arches and Sgr B2. The catalog contains 1352 objects that are highly-absorbed ($N_{\\rm H}$$\\ga$$4\\times10^{22}$ cm$^{-2}$) and are therefore likely to lie near the Galactic center ($D$$\\approx$8 kpc), and 549 less-absorbed sources that lie within $\\la$6 kc of Earth. Based on the inferred luminosities of the X-ray sources and the expected numbers of various classes of objects, we suggest that the sources with $L_{\\rm X} \\la 10^{33}$ \\ergs\\ that comprise $\\approx$90\\% of the catalog are cataclysmic variables, and that the $\\approx$100 brighter objects are accreting neutron stars and black holes, young isolated pulsars, and Wolf-Rayet and O stars in colliding-wind binaries. We find that the spatial distribution of X-ray sources matches that of the old stellar population observed in the infrared, which supports our suggestion that most of the X-ray sources are old cataclysmic variables. However, we find that there is an apparent excess of $\\approx$10 bright sources in the Radio Arches region. That region is already known to be the site of recent star formation, so we suggest that the bright sources in this region are young high-mass X-ray binaries, pulsars, or WR/O star binaries. We briefly discuss some astrophysical questions that this catalog can be used to address. ",
        "introduction": "The exquisite sensitivity of the {\\it Chandra X-ray Observatory} has provided vast improvements in our understanding of faint, hard X-ray sources. \\chandra\\ observations of distant galaxies allow us to study the X-ray population at luminosities similar to those accessible in our own Galaxy with wide-field X-ray instruments like the {\\it Rossi X-ray Timing Explorer} All-Sky Monitor and the {\\it BeppoSAX} Wide-Field Camera ($L_{\\rm X} \\sim 10^{36}$~\\ergs), while observations of our own Galaxy are sensitive to sources a million times fainter than found with previous wide-field surveys. One of the most dramatic products of this improvement in sensitivity are the \\chandra\\ observations of the Galactic center (Wang, Gotthelf, \\& Lang 2002; Baganoff \\etal\\ 2003)\\nocite{wgl02,bag03}. Whereas previous imaging surveys identified dozens of X-ray sources against a background of bright Galactic diffuse emission \\citep{wat81,pav94,pt94,sbm01,sid99,sak02}, % \\chandra\\ observations have revealed thousands of individual X-ray sources \\citep[e.g.,][]{wgl02,mun03a} and discrete, filamentary features \\citep[e.g.,][]{lwl03,yz05a}. This high concentration of X-ray sources is not surprising. The 2\\degree$\\times$0\\fdeg8 \\citep[300$\\times$125 pc for a distance of 8 kpc;][]{mcn00} region centered on the Galactic center contains roughly 1\\% of the Galactic mass (Lauhardt, Zylka, \\& Mezger 2002)\\nocite{lzm02}. Moreover, unlike the Galactic bulge, star formation has occurred continuously in the central region of the Galaxy, as is strikingly illustrated by the $\\ga$60 ultra-compact HII regions in the giant molecular cloud Sgr B2 \\citep{dgg98}, and by three young, dense clusters of massive stars \\citep[the Arches, the Quintuplet, and the Central Parsec;][]{kra95,fig99}.  \\begin{figure*}[th] \\centerline{\\epsfig{file=survey_color.eps,height=0.9\\linewidth,angle=90}} \\caption{Mosaic image of the 0\\fdeg8$\\times$2\\degree\\ field along the Galactic plane, centered on \\sgrastar. The raw image has been adaptively smoothed with the CIAO tool {\\tt csmooth} for display purposes. The prominent features are the Sgr A complex at the center of the image, two bright X-ray binaries and their dust scattering halos at $l$=0.275\\degree and $-$0.88\\degree. Only a fraction of the brightest point sources are visible in this image. Note that the smoothing algorithm introduces significant artifacts, especially at the edges of the deep observations centered on Sgr B2 and the Arches. } \\label{fig:img} \\end{figure*}  A wealth of questions can be addressed with \\chandra\\ observations, and with subsequent comparisons to multi-wavelength catalogs. For instance, a variety of studies of the synthesis of compact, accreting binaries have been designed to explain the large number of X-ray sources \\chandra\\ detects in the Galactic center, and the results constrain, for example, the amount of angular momentum dissipated in the common envelope phase (e.g., Pfahl, Rappaport, \\& Podsiadlowski 2002; Belczynski \\& Taam 2004; Liu \\& Li 2005; Ruiter, Belczynski, \\& Harrison 2005). \\nocite{prp02,bt04,ll05,rbh05} \\chandra\\ also can detect outbursts from transient low-mass X-ray binaries (LMXBs) at much lower flux levels than are accessible with traditional wide-field X-ray surveys, which provides unique insight into the duty cycles and emission mechanisms of compact objects accreting at very low rates \\citep[$\\dot{M} \\la 10^{-11}$ \\msun\\ yr$^{-1}$; e.g.,][]{king00,wmw02,ww02,inz05,sak05}. Combining \\chandra\\ and radio observations can be used to identify young stars with powerful winds, which helps to constrain the rate at which massive stars have formed recently in the Galactic center \\citep{yz02,lyz04,mun05b}.  Several observational studies have discussed the population of X-ray sources in \\chandra\\ surveys of the Galactic center. \\citet{wgl02} presented images from shallow (12~ks) \\chandra\\ exposures of the 2\\degree$\\times$0\\fdeg8 around \\sgrastar, and gave a general overview of the number of X-ray sources and the properties of the diffuse X-ray emission. Takagi, Murakami, \\& Koyama (2002)\\nocite{tak02} studied the properties of X-ray sources associated with HII regions in Sgr B2. \\citet{yz02} and \\citet{lyz04} and studied the X-ray emission from massive stars in several clusters, including the Arches and Quintuplet. Finally, \\citet{mun03a,mun04b,mun05a}, presented a comprehensive study of the population of X-ray sources with $L_{\\rm X} = 10^{31} - 10^{33}$~\\ergs\\ that were discovered in a deep (625~ks) set of \\chandra\\ observations of the inner 25 pc around \\sgrastar. However, aside from the inner 25 pc, no catalog containing all of the X-ray sources found within the inner 150 pc of the Galactic center has been published.  In this paper, we rectify this by reporting the locations and basic properties of the X-ray sources detected in \\chandra\\ observations of the inner 2\\degree$\\times$0\\fdeg8 around the Galactic center. In \\S\\ref{sec:id}, we present the locations of these X-ray sources \\citep[excluding the inner 8\\arcmin\\ covered by the catalog in][]{mun03a} in order to facilitate searches for multi-wavelength counterparts. In \\S\\ref{sec:phot}, we present the fluxes and basic spectral properties in order to constrain the origin of the X-ray emission, and to serve as baseline measurements for future searches for transient sources. In \\S\\ref{sec:dist}, we examine the spatial distribution of the X-ray sources to determine how they are related to the stellar population that is observed in the infrared. In \\S\\ref{sec:ns}, we study the luminosity distribution of the X-ray sources, and report variations in the relative numbers of bright sources that could be related to recent star formation. ",
        "conclusions": "We have presented a catalog of X-ray sources with fluxes between a few $10^{31}$ and $10^{35}$ \\ergs. The majority of the region was covered with tiled pairs of 12~ks observations, in which the 50\\% completeness limit is $1\\times10^{33}$~\\ergs. Two deeper exposures were more sensitive. A 50~ks toward the Radio Arches was complete to $4\\times10^{32}$~\\ergs, and a 100~ks observations toward Sgr B2 was complete to $1\\times10^{32}$~\\ergs. The sensitivity of the Radio Arches observation was a factor of two poorer than would be naively expected based on the exposure time, because the Galactic center produces strong diffuse emission against which point sources are more difficult to detect.  Our survey encompasses a non-trivial fraction of the mass of stars in the Galaxy. Integrating the \\citep{lzm02} models for the stellar distribution over our survey area, we find that it encloses a stellar mass of $\\sim$$10^{9}$ \\msun, or 1\\% of the Galactic value. In Table~\\ref{tab:pop}, we list the various classes of object that could be detected as X-ray sources in our image if they were located near the Galactic center. We also include predictions for the total number of each class of sources encompassed by the survey region, based on the references in the table. The estimates are based on the assumption that the star formation rate is $\\sim$1\\% of the Galactic value, or 0.01 \\msun\\ yr$^{-1}$ \\citep{fig04}. However, if the recent rate of star formation in the central 150 pc of the Galaxy is $\\sim$10\\% of the total Galactic rate, as is suggested by indirect measurements of the Lyman-alpha flux in the region \\citep{cl89,fig99}, then WR/O stars, luminous pulsars, and high-mass X-ray binaries each could be an order-of-magnitude more numerous. Below we describe the considerations that went into developing that table, and some implications that the observed population of sources have for understanding the evolution of the accreting binaries that make up the majority of our sample.  \\subsection{Comparison to the Local Galactic X-ray Population}  We start by comparing the amount of X-ray flux per unit stellar mass from the point sources with $L_{\\rm X}$$>$$5\\times10^{32}$~\\ergs\\ in our survey to that in the local Galaxy as identified by \\citet{saz05}. For a cumulative number-flux distribution of the form $N(>L_{\\rm X}) = N_0 (L_{\\rm X}/L_{\\rm X,0})^{-\\alpha}$, the specific luminosity produced by sources with $L_{\\rm X,min} < L_{\\rm X} < L_{\\rm X,max}$ is \\begin{equation} {\\cal L}_{\\rm X,tot} = \\frac{\\alpha N_0 L_{\\rm X,0}}{\\alpha -1} \\left[ \\left(\\frac{L_{\\rm X,min}}{L_{\\rm X,0}}\\right)^{-\\alpha + 1} - \\left(\\frac{L_{\\rm X,max}}{L_{\\rm X,0}}\\right)^{-\\alpha + 1}\\right], \\end{equation} where $N_0$ is the normalization at a luminosity of $L_{\\rm X,0}$ in units of sources per solar mass. For the Galactic center, we have found that $\\alpha = 1.5$ and that $N_0 = (4\\pm2)\\times10^{-7}$ sources \\msun$^{-1}$ at $L_{\\rm X,0} = 5\\times10^{32}$ \\ergs\\ (Tab.~\\ref{tab:ns}, Fig.~\\ref{fig:dist}). So, the the specific luminosity of X-ray point sources with luminosities between $L_{\\rm X,min} = 5\\times10^{32}$ \\ergs\\ and $L_{\\rm X, max} = 10^{34}$ \\ergs\\ in our survey is ${\\cal L}_{\\rm X,tot} = (5\\pm2)\\times10^{26}$ \\ergs\\ \\msun$^{-1}$. In the local Galaxy, \\citet{saz05} find that $\\alpha\\approx 1.2$ and $N_0 = (6\\pm2)\\times10^{-4}$ sources \\msun$^{-1}$ at $L_{\\rm X,0} = 2\\times10^{30}$ \\ergs (where for $K$ in their Eq. 5, $N_0 = K/\\alpha$, and we have estimated the uncertainty based on that of total specific luminosity for sources with $L_{\\rm X} < 10^{34}$ \\ergs). So, the specific luminosity of sources with luminosities between $L_{\\rm X,min} = 5\\times10^{32}$ \\ergs\\ and $L_{\\rm X, max} = 10^{34}$ \\ergs\\ is locally ${\\cal L}_{\\rm X,tot} = (1.0\\pm0.3)\\times10^{27}$ \\ergs\\ \\msun$^{-1}$. Therefore, the specific luminosities of X-ray sources with $5\\times10^{32} < L_{\\rm X} < 10^{34}$ \\ergs\\ in the local Galaxy and in the Galactic center are consistent within their uncertainties.  To understand the population of X-ray sources at the Galactic center, it is notable that in the local Galactic neighborhood all the X-ray sources with $10^{32} < L_{\\rm X} < 10^{34}$ \\ergs\\ (2--10 keV) are cataclysmic variables (CVs) with magnetic white dwarfs and orbital periods of several hours (intermediate polars; Sazonov \\etal\\ 2005). Therefore, it is conceivable that most of the X-ray sources with $L_{\\rm X} < 10^{34}$~\\ergs\\ in our survey are magnetic CVs \\citep[see also][]{mun04b,lay05}. We would expect CVs to have the same spatial distribution as the old ($\\ga$ Gyr) population of stars, which dominate the infrared light from the Galaxy, and indeed the distribution of X-ray sources is identical to the inferred distribution of stars (Fig.~\\ref{fig:dist}). As described in \\citet{mun04b}, intermediate polars have particularly hard, intrinsicly-absorbed spectra that are consistent with those of the sources with $L_{\\rm X} \\la 10^{33}$ \\ergs\\ in Figure~\\ref{fig:hit}.  Few CVs have $L_{\\rm X} > 10^{33}$, and only one CV has been observed at $L_{\\rm X} \\ga 5\\times10^{33}$~\\ergs\\ \\citep[GK Per in outburst; e.g.,][]{krw79}, so we expect brighter sources to be more luminous objects such as LMXBs, HMXBs, WR/O stars in colliding-wind binaries, or pulsars (Table~\\ref{tab:pop}). The excess of bright X-ray sources observed in the Radio Arches field (Fig.~\\ref{fig:ns}), in which the Arches and Quintuplet clusters are striking evidence recent active star formation \\citep{fig99}, can be explained if some of the bright sources are HMXBs, WR/O stars, or young pulsars. Such sources have short lifetimes, and so should be concentrated near regions of active star formation. Such sources also have softer spectra than intermediate polars in the 2--8 keV band, and so could explain why the sources brighter than $\\sim$$10^{33}$ \\ergs\\ in Figure~\\ref{fig:hit} are systematically softer than the faint ones.  \\begin{deluxetable}{lcccc}% \\tablecolumns{5} \\tablewidth{0pc} \\tablecaption{X-Ray Sources in the Galactic Center\\label{tab:pop}} \\tablehead{ \\colhead{Object} & \\colhead{$\\log(L_{\\rm X})$} & \\colhead{Number} & \\colhead{Number} & \\colhead{References} \\\\ \\colhead{} & \\colhead{$\\log({\\rm erg~s}^{-1})$} & \\colhead{in GC} & \\colhead{Detectable} & \\colhead{} } \\startdata CVs & 29.5--33.5 & $10^{5}$ & $10^{3}$ & [1,2] \\\\ % WR/O Stars & 31--34 & $10^{3}$ & 10 & [3,4,5] \\\\ Pulsars & 29.3--35 & $10^{6}$ & 10 & [6,7] \\\\ LMXBs & 30--39 & $10^{3}$ & 10 & [8,9,10] \\\\ HMXBs & 31--38 & $10^{3}$ & 50 & [11] \\enddata \\tablerefs{[1] \\citet{ver97}; [2] \\citet{saz05}; [3] \\citep{fig04}; [4] \\citet{pol87}; [5] \\citet{ber97}; [6] \\citet{ba02}; [7] \\citet{cl97}; [8] \\citet{wij02}; [9] \\citet{ll05}; [10] \\citet{kon02}; [11] \\citet{prp02}.} \\tablecomments{ We list order-of-magnitude estimates of the total population of various X-ray sources in our field, along with the number that should be detected in our survey.} \\end{deluxetable}   \\subsection{Comparison to Theoretical Models}  Several binary population synthesis calculations have been carried out in order to interpret the population of X-ray sources in the Galactic center described by \\citet{wgl02} and \\citet{mun03a}. These models outline how the numbers of each class of X-ray sources constrain various combinations of parameters in the binary evolution models and assumptions about the physics of systems accreting at low rates.  For instance, \\citet{prp02} suggested that several hundred of the X-ray sources in the \\citet{wgl02} survey could be neutron stars accreting from the winds of $>$3\\msun\\ binary companions \\citep[see also][]{bt04}. The total numbers of wind-accreting neutron stars, and the fractions of systems with companions more and less massive than 8\\msun, varies by a factor of a few depending upon the the magnitudes of the kicks imparted to the neutron stars at birth. These theoretical predictions have motivated searches for infrared counterparts to the X-ray sources in the Galactic center surveys \\citep{lay05,ban05}. However, \\citet{ll05} suggested that accretion could be inhibited by the magnetospheres of the neutron stars \\citep[see also][]{dp81} at the low mass-transfer rates considered by \\citet{prp02}, so that wind-accreting neutron stars may not be luminous enough to be detected in our \\chandra\\ survey.  There also is theoretical disagreement as to whether CVs contribute significantly to the population of X-ray sources in our survey. \\citet{ll05}, who base their calculations on the binary evolution code of Hurley, Tout, \\& Pols (2002)\\nocite{htp02}, suggest that CVs are not luminous enough to be detected in large numbers from the Galactic center. However, this disagrees with similar calculations by \\citet{rbh05}, who use the StarTrack code \\citep{bel05}, and predict that significant numbers of luminous CVs should be detectable from the Galactic center. The main difference between the two codes is in their prescriptions for calculating the rate of mass transfer, which leads \\citet{ll05} to predict mass transfer rates $\\sim$100 times lower than those used by \\citet{rbh05} for identical systems with orbital periods of several hours (A. Ruiter, private communication). Systems with orbits of several hours have the highest accretion rates \\citep[][see also Howell, Nelson, \\& Rappaport 2001]{pat84}\\nocite{hnr01} and therefore are the most luminous in X-rays, which makes them the most important contributors among CVs to our survey. Our comparison with the local Galactic population of X-ray sources suggests that CVs are indeed both numerous and luminous enough to explain the population of X-ray sources in our image \\citep[see also, e.g.,][]{ver97,ei99,srr05}, so our results could be taken as further, indirect evidence in support of the prescription used by \\citet{bel05} and \\citet{rbh05}.  Finally, \\citet{bt04} and \\citet{ll05} predict that there are a few thousands LMXBs consisting of a neutron star accreting from a white dwarf in our survey region. In these models, most of the neutron stars in these systems form through the accretion-induced collapse of an ONe white dwarf. Under the assumptions of \\citet{ll05}, a few percent of the LMXBs are persistently bright enough to be detected in our survey. Moreover, most of these LMXBs should be transient, so if one assumes a standard duty cycle of $\\sim$1\\%, there should be $>$50 LMXBs in outburst with $L_{\\rm X} > 10^{36}$ \\ergs\\ in the field at any given time (and thousands in the Galaxy).  In contrast, there are only two persistent LMXBs this luminous in the field, 1E 1743.1--2843 and 1E 1740.7--2942 \\citep{wgl02}, and several decades of occasional X-ray observations have only revealed about a dozen transients with $L_{\\rm X} \\ga 10^{36}$ \\ergs\\ \\citep[and only $\\sim$150 in the rest of the Galaxy; see, e.g.,][]{liu01,wij05}. The production of a large number of these transient LMXBs appears to be a common feature of models in which neutron stars can form through accretion-induced collapse (e.g., \\S5.6 of Iben, Tutukov, \\& Yungelson 1995)\\nocite{ity95}. Under the models of \\citet{bt04} and \\citet{ll05}, possible ways to accommodate the small number of bright transients in our field are to assume that the efficiency with which the envelope of a star can be ejected during the common envelope phase is low so that many of the binaries merge rather than forming LMXBs, or to assume that the accretion-induced collapse of an ONe white dwarf does not form a neutron star."
    },
    "0601/astro-ph0601557_arXiv.txt": {
        "abstract": "The Swift X-Ray Telescope often observes a rapidly decaying X-ray emission stretching to as long as $ t \\sim  10^3$ seconds after a  conventional prompt phase. This component is  most likely due to a prompt emission  viewed at large  observer angles $\\theta > 1/\\Gamma$, where $\\theta\\sim 0.1$ is a typical viewing angle of the jet and $\\Gamma\\geq 100$ is the  Lorentz factor of the flow during the prompt phase. This can be used to estimate the  prompt  emission radii, $r_{em} \\geq 2  t c /\\theta^2 \\sim 6 \\times 10^{15}$ cm. These radii are   much larger than is assumed within a framework of a fireball model. Such large emission radii can be reconciled with a fast variability, on time scales as short as milliseconds, if the emission is beamed in the bulk outflow frame, \\eg due to a  random relativistic motion of ''fundamental emitters''. This may also offer a possible explanation for  X-ray flares  observed during early afterglows. ",
        "introduction": "Recently launched \\Swift satellite  \\citep{Gehrels} together with  a network of ground based observations have been providing scientific community with crucial information on Gamma Ray Bursts (GRBs). Besides the landmark detection of afterglows from short  GRBs \\citep[\\eg][]{Gehr05},  \\Swift has gathered crucial data on developments of GRBs at early  times. This is especially important since  early observations provide clues to the  properties of the ejecta, like  its composition, lateral distribution of energy etc. At late times the energy is mostly transfered to the forward shock, properties of which can hardly be used to probe the ejecta. A number of surprising results related to early afterglows have emerged \\citep[\\eg][]{Tagliaferri,Nousek,Chincarini,obrien}: (i) early, $t \\leq 10^3$ s,   rapidly-decaying X-ray component, (ii) X-ray flares occurring  at $t \\sim 10^2-10^4$ s, (iii) shallower than expected initial decay (or hump) of the afterglow. These features are common, but the  light curves show a large variety. In this letter we discuss the  first  two mentioned effects, \\ie rapidly-decaying component and X-ray flares, since both can be related to the prompt emission (as opposed to afterglow) and can thus be used to probe the ejecta and the  central engine. ",
        "conclusions": "\\label{Conc}   In this letter we first point out that the interpretation of  the initial fast-decaying part of the X-ray GRB light curves as a prompt emission seen at large angles,  and a generic estimate of jets' opening angle allows a measurement of the radius of prompt emission, which turns out to be relatively large, $> 10^{15}$ cm. On basic grounds, $\\gamma$-ray  emission should be generated before the deceleration radius $r_{dec} \\sim \\left( { E_{iso} \\over 4 \\pi \\rho c^2 \\Gamma_{dec}^2} \\right)^{1/3}\\sim 10^{16}- 10^{17}$ cm, when most energy of the outflow is given to the surrounding medium (here $ E_{iso}$ is isotropic equivalent energy,  $ \\rho$ is density of external medium, $\\Gamma_{dec}$ is  Lorentz factor at $r_{dec}$). \\footnote{Note, that $r_{dec}$ defined above is {\\it independent} of ejecta content, contrary to the claim in \\cite{zk05}, see \\cite{LZK}.} The inferred emission radius is within this limit.  The estimate of the emission radius is very simple, and, in some sense, generic. It can hardly be consistent with the fireball model, unless extreme assumptions are made about the parameters (\\eg very large Lorentz factor). On the other hand, there are  alternative models (\\eg the electromagnetic model  \\citep{l05}, see also \\cite{thompson}) that  place prompt emission radii  at large distances, just before the deceleration radius $r_{dec} $.    Secondly, we show how   models  placing emission at large radii may be able to reproduce a short time scale variability of the prompt emission and explain later  X-ray flares. This can be achieved if  the prompt emission is beamed in the rest frame of the outflow, which may be due to  an internal relativistic motion of ''fundamental emitters''.  What can produce a relativistic motion  in the bulk frame? It can be due, for example, to a  relativistic Burgers-type turbulence (a collection  of randomly directed shock waves). It is not clear how such turbulence may be generated. Alternatively, relativistic internal sub-jets can  result  from   reconnection occurring in highly magnetized  plasma with $\\sigma \\gg 1$, where $\\sigma$  is a plasma magnetization parameter \\citep{KC84}. In this case the  matter outflowing from a reconnection layer  reaches relativistic speeds with $\\gamma_{out} \\sim \\sigma$ \\cite[]{lu03}. Internal synchrotron emission by such jets, or Compton scattering of ambient photons, will  be strongly beamed in the frame of the outflow. Note, that {\\it this model does not require late engine activity } to produce flares.   One of the main observational complications is that at observer times larger than  the conventional prompt phase, the X-ray light curve is a sum of the tail of the prompt emission, coming presumably from internal dissipation in the ejecta, and the forward shock emission. It is not obvious how to separate the two  components. For example, GRBs which do not show a fast initial decay may be dominated by the forward shock emission from early on \\citep{obrien}. This uncertainty also affects estimates of the emission radius since the  end of the steep decay may be related to the emergent afterglow emission and not to the jet opening  angle (or observer's angle, in case of a structured jet), see Fig. \\ref{GRBafter}. Another complication is that at these intermediate times, $10^3 \\leq t  \\leq 10^4$ s,  even the forward shock emission  itself  often does not conform to the standard afterglow models, showing  flatter than expected profiles \\citep[\\eg][]{Nousek}.  A consequence of the model is that {\\it some} short GRBs may be just a single spike directed towards an observer of a long GRBs. In our model the shorter spikes are  highly beamed, less frequent and produce harder emission. This can apply only to {\\it some} short GRBs since as a class they are well established to have different origin than  long GRBs (from non-observation of a supernova signature and coming from a distinctly different host galaxy population)."
    },
    "0601/astro-ph0601282_arXiv.txt": {
        "abstract": "{The possible existence of  strange stars in the universe will help in the understanding of various properties of quantum chromodynamics, like asymptotic freedom and chiral symmetry restoration, which is otherwise very difficult to prove in laboratory experiments. Strange star properties were calculated using large color approximation with built-in chiral symmetry restoration. A relativistic Hartree Fock calculation was performed using the Richardson potential as an interquark interaction. This potential has the asymptotic freedom and a confinement-deconfinement mechanism built into it and the present calculation employs an application of this potential with modified two scale parameters $\\Lambda$ and $\\Lambda^{\\prime}$, to find a new set of equations of state for strange quark matter. The linear confinement string tension from lattice calculations is 350 $MeV$ and the Coulomb -like part has the parameter 100 $MeV$ from deep inelastic scattering experiments. We also consider the effect of temperature, on gluon mass in a simple way, in addition to the usual density dependence. We obtained a set of new equation of states for strange quark matter and also found that the transition temperature from hadronic matter to strange matter is at 80 MeV, close to the 100 MeV estimated in litarature.Formation of strange stars may be the only signal for formation of quark-gluon plasma with asymptotic freedom and chiral symmetry restoration and this may be observable through many processes -such as for example through delayed $\\gamma$ ray afterglow. ",
        "introduction": "Two decades have passed since the existence of strange matter and strange stars was proposed by Witten (1984), and it is still difficult to prove or disprove the existence of a state of strange quark matter in its purest form. Besides the mass radius relation, efforts have been made to find other signatures for the existence of quark stars. The existence of absorption bands at 0.35 $keV$ with harmonics at 0.7, 1.05, 1.4 etc, in 1E 1207$-$5209 was interpreted in \\cite{mpla} as evidence for strange stars. It was then observed that six stars show harmonic band emission at energies corresponding to the same compressional modes at which the absorption bands were invoked (\\cite {MNRAS}). The SPI spectrometer on the INTEGRAL satellite detected a bright 511 keV line in the $\\gamma$ - ray spectrum from the bulge of the galaxy with a spherically symmetric distribution (\\cite{knodjean}). This has stimulated research to the fundamental physics that describes cosmological dark matter. Several authors (\\cite{bhscp04}, Oaknin and Zhitnitsky, 2005) have suggested that the line comes from positron annihilation, possibly in strange stars made of antiquarks. The implication of cosmic separation of phases (Witten, 1984) is that strange stars (SS) can exist since the early universe and can themselves explain dark matter. A crucial condition for the survival of these stars, given by Alcock and Olinto (1989), is a large value of surface tension $\\sim(178~MeV)^3$ which can be obtained from strange star equation of state (EOS) (\\cite{oursurf}). In his recent review, Weber (2005) argued that the constraints on mass radius relations from observational data support that the compact objects SAX J1808.4$-$3658, 4U 1728$-$34, RX J1856.5$-$3754 and Her X$-$1 are strange stars rather than neutron stars and on the basis of its low value of temperature, PSR 0205$+$6449 may be considered as a strange star. He also discussed the possibility of explaining some astrophysical phenomena like X-ray bursts, gamma-ray bursts and soft gamma repeaters with a strange star model. The most controversial is the isolated compact star RX J1856.5$-$3754, which has a featureless thermal spectrum, and its X-ray data fits quite well with a blackbody distribution. With a distance estimate of 120 pc, Drake et al. (2002) gave the radiation radius of the star to be 4.12$\\pm$0.68 km. This indicates that the stellar radius is 4 km to 8 km. Later, with a better estimation of the distance of 175 pc, the estimated radiation radius $R^\\infty \\simeq 8$ km. However, there are controversies regarding the surface temperature of this star. With an estimated surface temperature ($T^\\infty$) of 63.5 $eV$, $R^\\infty \\simeq 4.4$ km, and with $T^{\\infty}  \\simeq 32 eV$, $R^\\infty \\simeq 8$ km (\\cite{bur03}).  There is evidence, as recently discussed by Bombaci, Parenti and Vidana (2004) - for the existence  of compact strange and neutron stars with larger radius and a shift from the latter to the former can account for delayed gamma ray bursts.   PSR 1913+16 is the first  pulsar to be discovered in a double neutron star system is, also known as the Hulse Taylor pulsar (1975). It is estimated to have a mass of 1.441 M$_\\odot$ and the recently discovered pulsar (Lyne et al., 2004) PSR~J07373039A, forming a binary pulsar system, has a mass of 1.337 M$_\\odot$. Thus the mass of pulsars may lie near these values, the observed variation of periodicity of the pulsars indicates the variation in their masses. Although it may be a stringent lower-limit constraint for neutron stars masses, which are bound mainly by gravitational forces, the same does not hold for compact objects, which are self bound, like a system of quark matter. Theoretically there is no physical reason to have a lower limit for the mass of strange stars, where strong interactions play an important role in stellar binding. The same EOS will allow even a very small quark nugget.   In the literature, there are several EOSs for strange quark matter, starting from the Bag model (\\cite{afo86}, \\cite{han86}, \\cite{ket95}) to the recent models like the mean field model with interacting quarks (\\cite{D98}), the strange star EOS model in the perturbative quantum chromodynamics (QCD) approach (\\cite{fps01}), those of the chiral chromodielectric model (\\cite{mft03}) or Dyson-Schwinger model (\\cite{blas99}), etc. In our present work, we will consider only the relativistic mean field model by Dey et al. (1998).   In the relativistic Hartree-Fock calculation of a strange quark matter EOS to form a compact star (\\cite{D98}), a phenomenological inter-quark interaction, namely the Richardson potential, was used. Strange stars are more compact than conventional neutron stars and fit into the Bodmer (1971) -- Witten (1984) hypothesis for the existence of strange quark matter. They were able to produce a very compact stellar structure with a mass of about 1.44 $M_\\odot$ and radius of about 7 km. The original Richardson (1979) potential was designed to obtain the mass spectrum of heavy mesons (Charmonium and Upsilon). It includes two features of the qq force, the asymptotic freedom (AF) and confinement, however, with the same scale, ${\\Lambda}$.   The ${\\Lambda}$ parameter needed was 400 MeV. It was further applied to light meson spectroscopy  and baryon properties with the same value of ${\\Lambda}$ (Crater and van Alstine, 1984, Dey, Dey and Le Tourneux, 1986). Although both the confinement and AF scales are chosen to be the same in the Richardson potential, it is the confinement, not the AF part, that plays the more important role in determining the static properties of hadrons. However, when the AF part is important (as in strange star property calculations), ${\\Lambda}$ would be much smaller, around 100 MeV, as it is the scale for AF, obtained from perturbative QCD. Indeed, for self bound high density strange quark matter, Dey et al. (1998) found that ${\\Lambda}$ needed to be $\\sim 100 $ MeV. In the present paper, proper scales are chosen for both the parameters, as determined from a calculation by Bagchi et al. (2004) by fitting the masses and recently determined magnetic moments of $\\Omega^-$ and $\\Delta^{++}$.  In section~\\ref{sec:model} we describe the basic formalism of the model as given in Dey et al. (1998) and the calculations incorporating finite temperature effects. In section~\\ref{sec:modify} we present a fine tuning of the Dey et al. (1998) model, with a modified version of the Richardson potential, and also incorporate a temperature dependence of gluon mass. In section~\\ref{sec:result} we show the results and the plots, and draw conclusions and summarize in section~\\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc} A set of new EOSs for strange matter is presented - both for zero and very high $T$, using the Richardson potential with the value of ${\\Lambda}^{\\prime}~ \\simeq~300~{\\rm to}~350~MeV$ and the value of ${\\Lambda}~= ~100~MeV$, the two scales for confinement and AF respectively. This explains both the properties of the deconfined quark matter (our present work) and the properties of confined 3u and 3s baryons (Bagchi et al., 2004). Although the confinement part ${\\Lambda}^{\\prime}~$ is stronger, leading to a sharp surface, it is softened by medium effects, developing a screening length. Inside the star, the AF part is more important. We have found that for a wide range of parametric variation, the strange matter EOS gives a minimum of energy per baryon (E/A) that is lower than ${(E/A)}_{Fe^{56}}=~930.6~MeV$ which ensures that the system is stable, and unlike neutron star - like structures, our strange stars are not solely bound by gravitational forces but also by strong forces. Moreover, considering temperature dependent screening of the potential, we have found that strange stars can sustain stable configurations up to a temperature of 80 $MeV$.  The indirect proof of the existence of strange stars is becoming more prominent because there are many observed features from compact objects that cannot be explained by the standard neutron star model, like spectral features and the controversies built around the star RX J1856.5$-$3754. Another feature, superbursts, faces difficulties in using the standard carbon burning scenario in the neutron star crust. The carbon, nitrogen and oxygen mass fraction ($Z_{CNO}$) must be larger than $Z_{CNO,\\odot}$, where the latter refers to the standard value found in the sun (Cooper et al. 2005). Sinha et al. (2002) predicted that an alternative to the carbon burning scenario in a neutron star to explain superbursts, can be formation of diquark pairs on the strange star surface and their subsequent breaking, thus giving a continuous and prolonged emission of energy that can be comparable to the superbursts. Recently, Page and Cumming (2005) tried to explain superbursts as carbon burning of matter accumulated on the electron cloud of a strange star. There are many other features that can be fitted better with a strange star than a neutron star model. So, the study of strange matter and strange stars is important and the formulation of a realistic strange star model is needed."
    },
    "0601/astro-ph0601474_arXiv.txt": {
        "abstract": "{We present results from archival Very Large Array (VLA) data and new VLA observations to investigate the long term behavior of the circular polarization of M81*, the nuclear radio source in the nearby galaxy M81. We also used the Berkeley-Illinois-Maryland Association (BIMA) array to observe M81* at 86 and 230 GHz.  M81* is unpolarized in the linear sense at a frequency as high as 86 GHz and  shows variable circular polarization at a frequency as high as 15 GHz.  The spectrum of the fractional circular polarization is inverted in most of our observations.  The sign of circular polarization is constant over frequency and time. The absence of linear polarization sets a lower limit to the accretion rate of $10^{-7} M_\\odot y^{-1}$. The polarization properties are strikingly similar to the properties of Sgr A*, the central radio source in the Milky Way. This supports the hypothesis that M81* is a scaled up version of Sgr A*. On the other hand, the broad band total intensity spectrum declines towards milimeter wavelengths which differs from previous observations of M81* and also from Sgr A*. ",
        "introduction": "The nearby spiral galaxy M81 (NGC\\,3031) shares many properties with the Milky Way. It is similar in type, size and mass and it also contains a nuclear radio source, M81*, that is most likely associated with a supermassive black hole. M81* has been studied extensively using Very Long Baseline Interferometry (VLBI) in the past. \\citeN{BietenholzBartelRupen2000} resolved M81* into a stationary core with a one sided jet. Multi-wavelength (\\citeNP{HoFilippenkoSargent1996}) and sub-millimeter observations (\\citeNP{ReuterLesch1996}) showed many similarities between M81* and Sgr A*, the central radio source in our Milky Way (\\citeNP{MeliaFalcke2001}). A jet model of Sgr A* has been applied to M81*, where it can reproduce the radio flux density and the size of the radio core by changing the accretion rate (\\citeNP{Falcke1996}). The sizes of both radio sources show a $\\sim 1/\\nu$ dependency on the frequency (e.g. \\citeNP{BietenholzBartelRupen2004} for M81*, and \\citeNP{BowerFalckeHerrnstein2004} and \\citeNP{ShenLoLiang2005} for Sgr A*).  M81* is an apparent transitional object between Sgr A* and high luminosity AGN.  As the brightest of the nearby low luminosity AGN (LLAGN), it is 5 orders of magnitude brighter than Sgr A* at radio wavelengths.  M81* is substantially underluminous at X-ray wavelengths ($L\\sim 10^{-5} L_{edd}$), yet not as much as Sgr A* ($L\\sim 10^{-10} L_{edd}$). Still, it is the faintest LLAGN we can study.  Furthermore, the polarization properties of M81* and Sgr A* are very similar. Sgr A* shows circular polarization in absence of linear polarization (\\citeNP{BowerFalckeBacker1999}, \\citeNP{BowerBackerZhao1999}, \\citeyearNP{BowerWrightBacker1999}) and we detected the same behaviour in M81* (\\citeNP{BrunthalerBowerFalcke2001}.) The polarization properties of Sgr A* and M81* are in contrast to the properties of most radio jets in active galactic nuclei where linear polarization often exceeds circular polarization by a large factor (e.g.~\\citeNP{WardleHomanOjha1998}; \\citeNP{RaynerNorrisSault2000}). The absence of linear polarized emission in Sgr A* can be explained as a consequence of the accretion flow.   \\citeN{BowerFalckeSault2002} investigated the long term behavior of the circular polarization in Sgr A* from archival VLA data and showed that {\\it i)} the circular polarization is variable on timescales of days to months, {\\it ii)} the sign of the circular polarization stayed constant over the entire time range of almost 20 years, and {\\it iii)} the average spectrum of circular polarization is inverted.  After the discovery of circular polarization in M81* we used the VLA to investigate the variability of the circular polarization on short timescales. We used additional archival VLA data to investigate the long term behavior of the circular polarization of M81*. ",
        "conclusions": "We have presented VLA observations of M81* from 1994 until 2002 that show that circular polarization is present at 4.8, 8.4, and 15 GHz in absence of linear polarization. The fractional circular polarization is variable on timescales of days and months and not correlated with the total flux density of the source. The sign of the circular polarization was, if detected, at all frequencies and times always positive. The polarization properties are strikingly similar to the properties of Sgr A*, the central radio source in the Milky Way. This supports the hypothesis that M81* is a scaled up version of Sgr A*. \\citeN{AitkenGreavesChrysostomou2000} and \\citeN{BowerWrightFalcke2003} detected linear polarization at 230 GHz and higher frequencies that also shows variability (\\citeNP{BowerFalckeWright2005}; \\citeNP{MarroneMoranZhao2006}). Given the similarity between M81* and Sgr A* we expect to see also linear polarization in M81* at higher frequencies."
    },
    "0601/astro-ph0601642_arXiv.txt": {
        "abstract": "We have been constructed a brand-new radiation hydrodynamics solver based upon Smoothed Particle Hydrodynamics (SPH), which works on parallel computer system. The code is designed to investigate the formation and evolution of the first generation objects at $z \\gtrsim 10$, where the radiative feedback from various sources play important roles. The code can compute the fraction of chemical species e, H$^+$, H, H$^-$, H$_2$,  and H$_2^+$ by fully implicit time integration. It also can deal with multiple sources of ionizing radiation, as well as the radiation at Lyman-Werner band. We compare the results for a few test calculations with the results of one dimensional simulations, in which we find good agreements with each other. We also evaluate the speedup by parallelization, that is found to be almost ideal, as far as the number of sources is comparable to the number of processors. ",
        "introduction": "\\label{intro} One of the important objectives of cosmology today is to understand the way how the first generation stars or galaxies are formed and how they affected the later structure formation, and how they reionized the surrounding media. These problems have been studied intensively in the last decade, hence we already have some knowledge on these issues. However, studies which properly address the radiation transfer effects are still at the beginning, in spite of their great importance.  In order to investigate the radiation transfer effects on such issues in realistic cosmological density field, we need the code which can deal with hydrodynamics, gravity, chemical reactions, and naturally radiation transfer in three dimension. So far, various kinds of numerical approaches have been taken by several authors at different levels. There are several codes which focusing on the fragmentation of primordial gas \\citep{Abel00,Bromm99,Bromm02,NU99,NU01,Yoshida03}, without radiation transfer. There  also are the radiation transfer codes in which hydrodynamics are not coupled \\citep{Abel99, NUS01, Raz02, MFC03}. Some of these codes have the merit that they can include diffuse radiation field \\citep{NUS01, Raz02, MFC03}. Utilizing this strong point, these trials can give answers to the cosmic reionization problem to some extent, however, they do not fit to investigate the problem of structure formation because of the dynamical nature of the phenomena and radiative feedback effects on star/galaxy formation. Moreover, cosmic reionization problem itself is also  coupled with structure formation, since the formed stars, galaxies and black holes are the sources of reionization. Therefore we need some assumptions on the nature and number of the sources when we try to investigate the reionization of the Universe with this type of simulations.  Another approach to these issues is suggested by \\citet{GA01}, in which paper they propose a new type of approximation, called as Optically Thin Variable Eddington Tensor method (OTVET). They solve the moment equations of transfer equation on the assumption that the closure relation is given by the Eddington tensor which is assessed under the optically thin approximation. The advantage of the approximation is that radiation transfer can be coupled with hydrodynamics, since this approximation reduces the computational cost dramatically. Therefore they could trace the cosmic reionization history in a self-consistent fashion.  Recently, \\citet{Mellema05} and \\citet{saru} have constructed fully coupled Radiation Hydrodynamics (RHD) codes applying on the cosmic reionization problem. They couple the adoptive mesh refinement hydrodynamics codes with adoptive ray-tracing codes. They successfully trace the radiation hydrodynamic expansion of HII region, as well as the formation of shadows behind dense clumps. However, their codes still do not include H$_2$ chemistry and transfer of Lyman-Werner band radiation, although they are crucial for the calculation of first structure formation.  Another branch of radiation hydrodynamic approach was proposed by \\citet{KB00}, in which they show a scheme of ray tracing in Smoothed Particle Hydrodynamics (SPH). SPH is often applied to the simulations on galaxy formation and star formation, because of its Lagrange nature and simplicity. They create grid points on  light rays from sources to SPH particles. The code has an advantage that the ray tracing is naturally incorporated into Lagrangian scheme, since they utilize the neighbor lists of particles to find the grid points. We have parallelized their scheme in case light rays from the source can be approximated as parallel \\citep{SU04a}. However, the parallelized code cannot be applied to the problems with multiple sources, that are crucially required for the first structure formation.  In this paper, we describe the newly developed RHD code by ourselves, in which fully coupled hydrodynamics, gravity, chemical reactions with H$_2$ are included, as well as radiation transfer of ionizing photons and the Lyman-Werner band photons. One of the characteristic feature of this code is that we employ SPH. Radiation transfer part of the code is a natural extension of our previous code described in \\citet{SU04a}. Thanks to the newly developed radiation transfer and parallelization scheme, we can include multiple sources in our new code. Since it is a particle code, it also has an advantage that it has the good affinity to GRAPE system \\citep{FMK05} which can be dedicated to the gravitational force calculations and neighbor search. We also show the outcomes of some calculations on the test problems, that are compared with the results of one dimensional simulations.  This paper is organized as follows. We describe the detail of the code in section \\ref{code}. Then we show the results of test calculations in section \\ref{results}, and summarize in the final section \\ref{summary}. ",
        "conclusions": "\\label{summary} We have constructed the new code for radiation hydrodynamics designed to be applied to the feedback effects of first generation stars/galaxies. The code can deal with multiple sources, and it can solve the transfer of ionizing photons and photodissociating photons.  The code is already parallelized, and almost ready to be installed in the newly developed full size huge PC cluster ``FIRST'' in University of Tsukuba. We compare the numerical results with one dimensional calculations, and find the good agreements.   \\bigskip We thank the anonymous referee for important comments. We also thank  M. Umemura and T. Nakamoto, K. Ohsuga for careful reading of the manuscript.  We thank N. Shibazaki for continuous encouragement. We thank H. Maki for providing the 1D-PPM code. We also thank T. Kitayama for providing the 1-dimensional results, and thank all the collaborators in TSU3 project ( http://www.mpa-garching.mpg.de/tsu3 ) for fruitful discussions during the workshop in CITA. The analysis has been made with computational facilities at Center for Computational Science in University of Tsukuba and Rikkyo University. This work was supported in part by Ministry of Education, Culture, Sports, Science, and Technology (MEXT), Grants-in-Aid, Specially Promoted Research 16002003 and Young Scientists (B) 17740110.  \\appendix"
    },
    "0601/astro-ph0601368_arXiv.txt": {
        "abstract": "A new single-pulse behaviour has been identified in two pulsars, B0919+06 and B1859+07.  Normally both stars emit bright subpulses in a region near the trailing edge of their profile.  However, occasionally these stars exhibit ``events'' wherein the emission longitude gradually decreases, by about their profile width, remains in this position for typically several tens of pulses, and then gradually returns over a few pulses to the usual longitude.  The effect bears some resemblance to a profile ``mode change'', but here the effect is gradual and episodic.  When the separate profiles of the normal and ``event'' emission are studied, they reveal a broad and complex profile structure in each pulsar---but one which can probably be understood correctly in terms of the geometry of a single conical beam. Possibly the effect entails an extreme example of ``absorption''-induced profile asymmetry, as suspected in other pulsars. Alternatively, shifting sources of illumination within the pulsar's beam may be responsible. ",
        "introduction": "Pulsar emission phenomena have proven to be a rich area of study, and these specific effects have often provided needed insights for theory building.  The \"classical\" six such effects, drifting subpulses, nulling, profile mode changing, microstructure, polarization modes and \"absorption\" are well known, and several further well defined behaviours have been identified in recent years.  We describe below a further such effect which has now been found in two poorly studied pulsars in the course of analysing sensitive new Arecibo polarimetric pulse-sequence (hereafter PS) observations.  The two stars, B0919+06 and B1859+07, both exhibit highly asymmetric average profiles with a long leading region of weak emission culminating in a bright \"component\" with a sharp trailing edge.  What attracted our attention, however, was that their bright subpulses, usually found in the trailing region, occasionally moved progressively over near the leading profile edge for several score pulses and then progressively back again.  This effect -- which we denote as an \"emission shift\" -- is not a drift, nor is it strictly a profile mode change.  Nor is it an effect easily visualized in terms of a rotating subbeam ``carousel'' sytem.  It is possible that the effect is related to the increasingly fascinating and potentially important property of \"absorption\", or even some more exotic ``Christmas lights'' phenomenon whereby different parts of the profile illumine at different times. Therefore, we thought it important to present a preliminary study of the phenomenon so that other pulsar investigators may become aware of it. \\S 2 describes our observations, \\S 3 \\& 4 examine the character of the phenomenon from the perspectives of individual pulses and average profile, respectively.  \\S 5  \\& 7 assess the implied emission geometries of the two stars, \\S 7 entertains possible explanations and \\S 8 summarizes the results.  \\begin{table} \\caption{Available Observations} \\begin{tabular}{ccccc} \\hline Frequency  &  MJD & Resolution  & pulses & events \\\\ (MHz)     &            & ($\\deg$)     &               &   \\\\ \\hline B0919+06  &&& \\\\ 1404 & 44857 &  0.42 & 3582 &      2    \\\\ 1404 & 44859 &  0.84 & 1886 & none  \\\\ 1400 & 52854 &  0.43 & 1115 &      2    \\\\ 327 & 52916 &  0.43 & 4180 &      1    \\\\ B1859+07  &&& \\\\ 1400 & 52739 & 0.64 &  1021 &      6    \\\\ 1400 & 53372 & 0.64 &  2096 &   8-10 \\\\ \\hline \\end{tabular} \\end{table}    \\begin{figure*} \\includegraphics[width=80mm]{fig1acolour2.ps} \\includegraphics[width=80mm]{fig1bcolour2.ps} \\caption{Pulse-sequence polarization displays showing ''emission shift'' effect. Two such events are seen in the B0919+06 sequence (left) and some seven in the B1859+07 (right) display.  In both cases the onset and return usually occurs over a duration of a few pulses with bright subpulses appearing to shift from the usual emission pattern to an earlier phase.  Both panels display a 1000-pulse sequence.  The total power $I$, fractional linear $L/I$, PA $\\chi$, and fractional circular polarization $V/I$ are colour-coded in each of the respective four columns according to their respective scales at the left of the diagram.  The latter three columns are plotted only when a given sample falls above a 2-sigma noise threshold.} \\label{Fig1} \\end{figure*}  \\section[]{Observations} The observations used in our analyses were made using the 305-m Arecibo Telescope in Puerto Rico.  The primary L and P-band polarized PSs were acquired using the upgraded instrument together with the Wideband Arecibo Pulsar Processor (WAPP\\footnote{http://www.naic.edu/$\\sim$wapp}) between 2003 July and early 2005  The ACFs and CCFs of the channel voltages produced by receivers connected to orthogonal linearly polarized feeds\\footnote{A quadrature hybrid was inserted between the feed and receivers of the P-band system on 11 October 2004, making it thereafter an orthogonal circular system.} were 3-level sampled.  Upon Fourier transforming, up to 256 channels were synthesized across a 100- (L) or 25-MHz (P) bandpass with a sampling time of roughly a milliperiod, providing overall resolutions well less than 1$\\deg$ longitude.  The Stokes parameters have been corrected for dispersion, interstellar Faraday rotation, and various instrumental polarization effects.  The L-band observations usually recorded four 100-MHz channels centred at 1275, 1425, 1525 and 1625 MHz, and the lower three of these were added together appropriately for greater sensitivity.  Older Arecibo observations at 1400 MHz (Stinebring \\etal\\ 1984) were used for comparison as detailed in Table 1.    \\section[]{The Pulse-Sequence Phenomenon} Figure~\\ref{Fig1} displays two 1000-pulse sequences (hereafter PSs), one of B0919+06 (left) and the other of B1859+07 (right).  Both exhibit episodes in which the emission moves sharply earlier leaving the usual region of bright subpulses almost empty.  Two such ``events'' are seen in B0919+06 and some 7 in B1859+07.  In both stars the emission moves earlier by an amount which is of the scale of the profile width.  Each event a) commences and exhibits an orderly shift, over a few pulses, of bright emission toward earlier longitude, b) remains in this state for up to some 20 or so pulses, and c) then gradually returns over a few pulses to the usual configuration and phase of emission.  In B0919+06, theses events are rather rare---one typically in several thousand pulses---which may account for why they have not been reported earlier.  In B1859+07, however, they typically occur every 1--2 hundred pulses, and in Fig.~\\ref{Fig1} they even appear quasi-periodic though much less so in the MJD 53372 observations (not shown).  In total power, the events almost resemble an effect which could be produced by an instrumental timing fault, and perhaps this is another reason that they have not been noted earlier.  In polarization they show up more clearly, in part because the linear and circular is plotted only when it exceeds a noise threshold.  This 2-sigma threshold is shown as a narrow white bar on the $I$-$L/I$ color bar, but it can only be seen on the B1859+07 bar at about the 5\\% level because in B0919+06 the signal-to-noise ratio (hereafter S/N) is so high that it disappears into the white region just above zero intensity.  \\begin{figure} \\includegraphics[width=80mm]{fig2colour.ps} \\caption{Pulse-sequence polarization display showing a ''shift'' effect in B0919+06 at 327 MHz as in Fig.~\\ref{Fig1}.  Here only a 200-pulse sequence is depicted.  The event begins about pulse 3750 and extends to pulse 3710.} \\label{Fig2} \\end{figure}  Close inspection shows that the events have different and characteristic linear polarization angles (hereafter PAs):  In B0919+06 the PA rotates smoothly from green ($+60\\deg$) on the trailing edge through cyan ($+30\\deg$) to almost blue ($+15\\deg$) on the leading edge of the usual profile, but during an event the PA reaches full purple ($-15\\deg$).  Similarly, in B1859+07 note that there is an region of smooth PA rotation that begins with near orange ($-75\\deg$) and rotates through green ($+60\\deg$) to cyan ($+30\\deg$) and even blue ($0\\deg$) during an event.  Also, there is a PA jump on the trailing edge to blue ($0\\deg$).  In B0919+06's normal emission pattern note that there is a region of enhanced linear polarization between about --3 and --5\\degr longitude and further that during the two events bright subpulses extend well in advance of --6\\degr.  It is also worth noticing that the region of reduced linear polarization on the leading edge of the normal pattern---which we have taken as the longitude origin---also tends to occur on the trailing side of the event pattern.  And both the trailing part of the normal pattern and leading side of the events show some weak negative circular polarization.  B1859+07's circular polarization is more interesting: we see consistent weak positive at about --2\\degr and a suggestion of weak negative circular around +2\\degr longitude.  Figure~\\ref{Fig2} shows an expanded view of the one event encountered in our only 327-MHz observation of B0919+06 on MJD 52916.  At this lower frequency the overall profile width is larger so that the event is less dramatic. Interestingly, it is hardly discernible in in total power, but in the linear and even circular polarization we see that the band of highly linearly polarized subpulses moves sharply to earlier longitudes following pulse number 3750 and remains in this configuration until or just after pulse 3800.  The shift is somewhat more apparent in the three polarization columns because the noise threshold defines the leading and trailing edges of the plotted area.  Note, though, that the PA values during the event represent a smooth rotation toward the yellow-orange angle ($-80\\deg$) from the blue-magenta-red (0 to $-60\\deg$) PA range that is characteristic of the usual emission pattern.   \\begin{figure*} \\includegraphics[width=80mm]{fig3a.ps} \\includegraphics[width=80mm]{fig3b.ps} \\caption{Total (bold) and partial average (lighter curves) profiles of B0919+06 at 327 MHz on MJD 52916 (left) and 1400 MHz on MJD 52854 (right) in the upper panels.  Solid curves are the total intensity (Stokes $I$), dashed the total linear (Stokes $L$), and dotted the circular polarization (Stokes $V$). The lower panels give PA histograms corresponding to the total profile of those samples having PA errors smaller than 3\\degr as well as the average PA curves for both the total and partial profiles.  The total profiles are very asymmetric with gradual leading and sharp trailing edges.  The partial averages include only those pulses associated with the ``events'' in the respective PSs. Note the markedly different profile forms but highly similar PA behavior.} \\label{Fig3} \\end{figure*}  \\begin{figure*} \\includegraphics[width=80mm]{fig4a.ps} \\includegraphics[width=80mm]{fig4b.ps} \\caption{Total (bold) and partial average (lighter curves) profiles of B1859+07 at 1400 MHz on MJD 52739 (left) and 53372 (right) as in Fig.~\\ref{Fig3}.  Here the partial averages are scaled up by a factor of 1.5 for easier comparison with the total profiles.  The lower panels give PA histograms of those samples having PA errors smaller than 3\\degr as well as the average PA curves for both the total and partial profiles.  The total profiles are again asymmetric with gradual leading and sharp trailing edges, whereas the partial averages including only the event pulses are nearly symmetric.  Again note the markedly different profile forms but highly similar PA behavior.} \\label{Fig4} \\end{figure*}    \\section[]{Average Profile \\& Modulation Properties} Let us now try to understand the ``shift'' events further by constructing sets of total (bold) and partial average (lighter curves) profiles.  These are shown in Figures~\\ref{Fig3} for B0919+06 and \\ref{Fig4} for B1859+07.  The top panels give the total intensity (solid curves), total linear polarization (dashed lines) and circular polarization (dotted lines).  The lower panels show the total PA distribution of samples having angle errors less than 3\\degr as well as average PA curves of the total (bold) and partial (lighter) averages.  The asymmetric total profiles of B0919+06 are well known from Stinebring \\etal\\ (1984) and Blaskiewicz \\etal\\ (1991), and even in the latter it was understood that the single bright peak at  21 cms aligns with the trailing component at 430 MHz, probably owing to the work of Phillips \\& Wolszczan (1991).  Fig.~\\ref{Fig3} shows us that the star's profiles (bold curves) have three main components at both frequencies.  At 327 MHz (left diagram) the center and trailing components are almost equally bright with the leading one much less so.  During an event, however (lighter curves), the leading component brightens up, the central one remains almost unchanged, and the trailing feature's intensity drops to half or less.  Note that the strong linear polarization under the trailing feature simply disappears during the event.  Then we see a polarized feature associated with the central component, a leading-edge extension of the region of smooth PA rotation, and indeed a poorly defined linear PA in the region under the trailing component.  A similar, but more dramatic behavior is seen at 1400 MHz in the right-hand display (Note the finer longitude scale).  Here the total profile asymmetry is so extreme that we see little apart from the trailing component.  The central feature is visible in neither the total power nor the linear polarization, and the leading component shows only as a small bump in both.  During the events, however, (lighter curves) the profile position shifts markedly earlier, and the three main components are in clear evidence.  The leading component is again much brighter, the central feature now as bright as the usual trailing feature, and the latter's intensity now again reduced to half or so.  We also see that the PA traverses of the total and event profiles track each other closely.  Finally, recall that the events in this pulsar are rather rare as shown in Table 1, so that the partial averages at 327 and 1400 MHz are comprised of only some 36 and 77 pulses, respectively.   Figure~\\ref{Fig4} gives a further set of total and ``event'' partial average profiles for the two observations of B1859+07 at 1400 MHz on MJDs 52739 and 53372. This star's two total average profiles (bold curves) are also very asymmetric with long gradual leading regions and sharp trailing edges.  Both observations also show several weak features which may be evidence of underlying component structure.  Then, during the events---which are considerably more frequent in this pulsar---the profiles (lighter curves) become more nearly symmetric.  While the total profiles do not exhibit any recognizable profile form, the partial averages are much more interesting.  Notice first the tripartite form of the total linear polarization $L$: we see a trailing feature aligned with the usual bright narrow total-power trailing component at about +9\\degr longitude, then a broad central ``hump'' between about $-8$ and +7\\degr, and finally a leading feature around $-10\\deg$.  The latter's linear power is small, but note in the lower panel that two different populations of bright polarized subpulse samples contribute to the power in this leading-edge region---one at a PA of roughly $-10\\deg$ and the other at some +80\\degr---contributing to the nearly complete depolarization on the leading edge of the partial profile.  We can also see that this polarization-modal behavior on the leading edge mirrors that on the trailing edge, which there has primary polarization-mode power at about +90\\degr immediately followed by secondary polarization-mode power at about 0\\degr.  Such adjacent configurations of modal power---producing 90\\degr PA ``flips'' and linear polarization ``nulls''---are most often seen on the outside edges of conal beams (Rankin \\& Ramachandran 2003).  Thus, they seem to mark the sightline crossings of the outer conal beam edge and indicate that even this extremely asymmetric profile is produced in substantial part by a full traverse through a conal emission beam.  Again, it is worth noting that the PA distribution reflects the properties of the total profile, but of course emission at longitudes earlier that about $-5\\deg$ only occurs during the ``events'' as can readily be seen in the right-hand panel of Fig.~\\ref{Fig1}.  Moreover, we see weak antisymmetric circular polarization in the total profiles about the longitude of the origin, and while it is only a few percent positive (near $-2\\deg$) and negative (near +2\\degr) it is produced by a population of highly circularly polarized subpulses as can again be seen in the righthand panel of Fig.~\\ref{Fig1}.  This in turn suggests some core activity in the center part of the profile and argues that overall we should regard the star's profile as reflecting emission from both a core and a cone.  Finally, Weltevrede \\etal\\ (2005) include both stars in their survey of pulse modulation characteristics.  In the case of B0919+06, not surprizingly they find evidence of a weak modulation with a $P_3$ of 50 $P_1$ or more.  In the case of B1859+07, their fluctuation spectra are almost ``white'', perhaps because the active leading-edge region was too weak to be detected.    \\section[]{B0919+06 Profile Geometry} Pulsar B0919+06 is bright over a very large spectral range and has thus been well observed both at very high and very low frequencies.  At 21 cms and above, its profile exhibits only a single component---the one we have identified as its trailing component---as we have seen in Fig.~\\ref{Fig3} (right).  Then, in a range including 300--500 MHz, it shows two components (\\eg, Fig.3, left), spaced by about 3\\degr.  At all lower frequencies the resolution has been such that these two features apparently merge into one broad feature.  The entire sweep of this star's profile evolution can be seen in Figure~\\ref{Fig5}, where a set of high quality total power profiles have been time-aligned over a frequency range from 50 to 5000 MHz (from Hankins \\& Rankin 2006).  \\begin{figure} \\includegraphics[width=80mm]{fig5.ps} \\caption{Time-aligned average profiles of B0919+06 between 49 and 4870 MHz. All were recorded using the Arecibo Observatory but those at 61 and 103 MHz which acquired at the Pushchino Radio Astronomy Observatory.  Bars indicating the instrumental resolutions are given at the left of the diagram.  From Hankins \\& Rankin (2006).} \\label{Fig5} \\end{figure}  Of course, these profiles should not have been aligned as they are:  The star's profiles are basically three-componented, so the middle (probably core) feature would provide a better alignment point.  This would shift the longitude origin some 4\\degr earlier---near the leading peak of the 430-MHz profile---so that the second peak as well as the higher frequency profile peaks align as a trailing conal component.  Then, the time-alignment of the 100-MHz and lower profiles seems to indicate that they too are dominated by this trailing component, and in no other profile do we clearly see the central component.  Quality polarimetry is available for B0919+06 at many frequencies from a number of different instruments.  In addition to Blaskiewicz \\etal\\ and Hankins \\& Rankin above, Gould \\& Lyne (1998), von Hoensbroech \\& Xilouris (1997), Rankin \\etal\\ (1989), Stinebring \\etal\\ (1984), Suleymanova \\etal\\ (1988) and Weisberg \\etal\\ (1999) have published profiles between 60 MHz and 5 GHz. Unfortunately, most of this work gives little insight into the leading ``ramp'' region of most interest to us here.  Another hint may come from the size of the triangular ``ramp'' on the profile leading edges.  We first see evidence of this in the asymmetry of the 1408-MHz profile, which we know from the above discussion is associated with the ``events'' and therefore with activity in the earlier components.  Here, its overall scale (see also Fig.~\\ref{Fig3}) is hardly 10\\degr, but from 430 MHz down it appears to have a nearly constant extent of about twice this.  The nearly symmetrical triangular ``ramp'' on this star's trailing edge at low frequencies is no less strange.  We see no hint of it in our 327-MHz profile, but at 111 MHz, it is a significant feature.  Nor can this be attributed to scattering; Phillips \\& Wolszczan (1992) have shown that scattering dominates the star's profile at 25 MHz, but appears to distort it little at 50 MHz.  Moreover, the available polarimetry at low frequencies (Suleymanova \\etal\\ 1988; Hankins \\& Rankin 2006) suggest that the triangular ``ramps'' are highly linearly polarized at an angle rotating about +9\\degr/\\degr over the entire profile.  These triangular low frequency features---not really components---are quite unusual and deserve further study.  Finally, returning to Fig.~\\ref{Fig3} the partial profiles provide further indications of the profile's structure, but not enough to significantly improve the geometrical models in Paper VIb.  Slightly revising the values there (see Table 5), a putative core width of some 4.7\\degr,  1-GHz profile width of about 10\\degr, and a PA rate of +9\\degr/\\degr indicates an inner cone geometry with magnetic latitude and sightline impact angles of 53 and 5.1\\degr, respectively, and a sightline cut some 25\\% inward of the outside conal 3-db edge.  This may be compatible with the apparent lack of conal spreading at very low frequency.  \\section[]{B1859+07 Profile Geometry} A more limited record of published observations is available for B1859+07, in part because its relatively short period (0.644 s) combined with a large dispersion measure (253 pc/cm$^3$) make it difficult to observe at lower frequencies.  Polarimetric observations have been published only by Gould \\& Lyne (1998) and Weisberg \\etal\\ (2004).\\footnote{Both observations show anomalies at lower frequencies: The former's 300-MHz profile is poorly resolved and defined polarimetrically; whereas the latter authors were unable to detect significant linear polarization.} Both show single profiles with a slow rise and steep trailing edge.  The average profile width increases with decreasing frequency; its half-power dimension is hardly 10\\degr\\ around 21 cm and may be as much as 20\\degr\\ in the 400-MHz region.  At no frequency does the star exhibit much fractional linear linear polarization, 20\\% at most.  Gould \\& Lyne's 1408-MHz profile resembles those in Fig.~\\ref{Fig4} in both form and fractional linear; even the PA seems to show both the leading ramp and abrupt $-90\\deg$ ``jump''.   The PA rate is about +6\\degr/\\degr, and as we had observed above, the clear ``patches'' of modal polarization on the edges of the event partial profile, together with the evidence for low frequency spreading, have the pattern of an outer cone.  Moreover, the emission near the center of the profile suggests core activity---and altogether argues that the underlying emission pattern is that of a triple ({\\bf T}) or possibly a five-component ({\\bf M}) geometry.  We can then compute the basic emission geometry as above using the procedures outlined in Rankin (1993b) for Table 5.  Estimating the outside (conal) half-power 1-GHz width of the profile as 20\\degr\\ from the ``event'' profiles in Fig.~\\ref{Fig4}, we find that the star has an outer cone geometry with $\\alpha$ and $\\beta$ some 31\\degr\\ and 4.8\\degr, respectively.  Interestingly, this geometry also implies a core width of 6.0\\degr, and while we are not able to measure this width accurately from the profiles, we can get some indication of its value from the circularly polarized power in individual pulses referred to above.  If we measure the typical separation between the respective leading positively and weaker trailing negatively circularly polarized subpulses associated with the core emission in Fig.~\\ref{Fig1} (\\eg, near pulse 150 or 650), the interval is clearly a few degrees.  In total power the central feature in the ``event'' partial profiles of Fig.~\\ref{Fig4} appears quite broad, but could also be comprised of several merged components (if intrinsically a double cone {\\bf M} profile).    \\section[]{How are ``Events'' to be Understood?} Since first noting the ``emission shift'' phenomenon a few months ago, we have tried to understand its character and causes.  We considered, for instance, whether the ``events'' might be produced by a displacement of the emitting region to a higher altitude within the polar flux tube.  This idea at first seemed appealing given the gradual nature of the ``event'' onsets and returns. However, on exploring the idea quantitatively, the displacements seemed too large in terms of vertical height in the magnetosphere. In both pulsars the magnitude of the swing is about 15 degrees, and interpreted in terms of light-travel-time would correspond to a vertical displacement about one quarter of the light-cylinder radius (5,500km in the case of B0919+06 and 8,000km for B1859+07). This figure greatly exceeds the height above the neutron star's surface within which the radio emission is widely believed to originate, and must be considered implausible. Furthermore, we have demonstrated in both pulsars a consistency with the picture of a single emission cone for both normal and shifted emission, based on the continuous PA change across the profile. A sudden increase in emission altitude would surely lead to an abrupt change to some new PA pattern and a dramatic widening of the total emission cone, and neither of these features are observed.  Rather, on closer inspection, the usual profile seems to reflect an incomplete illumination of the full geometric traverse of the sightline through the emission cone.  During the ``events'', the situation is very different:  here the partial profiles have recognizable forms, dimensions, component structures and even sensible classifications (\\eg, Rankin 1983a).  In both cases, the usual total profile is unusually asymmetric, but the partial profiles of the ``events'' have the usual and familar level of overall conal symmetry.  Similarly with the PA behaviour, especially of B1859+07:  the PA histograms suggest a complete profile with modal power on both edges, but this power is only observed on the leading edge during ``events''.  Given these circumstances, we began to ask whether the ``events'' could be constructively viewed as a kind of profile ``mode change''.  Surely we do see a good deal of evidence for more or less asymmetric profile illumination in the different modes of a given star.  What makes the ``emission shifts'' stand out, though, is their gradual character:  usually we see that the stars take a few single pulses to go from the ``normal'' emission pattern to that of an ``event''---and in general mode changes seem to occur on a time scale of a single pulse.  We can think of no good case of a gradual mode-change onset.\\footnote{We do have some evidence for a gradual as well as sharp character to mode changes in pulsar B0943+10 (Suleymanova \\etal\\ 1998; see also Rankin \\& Suleymanova 2005)} A recognizable reconfiguration of the emission pattern within about about a single stellar rotation is thus usually implied in a ``mode change''.  If the ``events'' are not a form of ``mode change'' (or even if they are), we have seen above that in both pulsars they entail a change in the emission pattern, from one which is concentrated ``late'' on the trailing edge of the profile---that is well past the probable longitude of the magnetic axis---to one which strongly favors a region earlier than the zero longitude.  Such a configuration and effect is very suggestive of the problematic ``absorption'' phenomenon first identified in pulsar B0809+74 (Bartel \\etal, Bartel 1981), but now identified in a number of situations (\\eg, Rankin 1983b).  Indeed, in the current context ``absorption'' provides an interesting model:  if one looks carefully again at Fig.~\\ref{Fig1}, it {\\em does} appear that the leading part of the emission pattern is usually obscured, but during an ``event'' it is rather the trailing part which is obscured.  A good deal of evidence in other pulsar contexts does indeed suggest that a partial or complete blockage of the radiation from either the leading of trailing side of the magnetic axis longitude does occur.  The most interesting recent evidence about the phenomenon comes from studies of the drifting-subpulse patterns of pulsars B0943+10 (Deshpande \\& Rankin 1999, 2001; Rankin \\etal\\ 2003; Rankin \\& Suleymanova 2005) and B0809+74 (Rankin \\etal\\ 2005a,b,c), where the rotating-subbeam ``carousel'' is usually visible only at longitudes later and earlier than the central longitude, respectively.  Fascinatingly, however, this is not always the case: in B0809+74 the ``absorption'' characteristics are frequency dependent and possibly are in some cases specific to one polarization mode.  In B0943+10 recent studies show that the ``absorption'' properties can change very slowly; the profile changes entailed in transitions from the star's Q to B emission modes take of order an hour to completely change from favouring emission prior to the central longitude to nearly precluding it.  However, here we have deliberately and cautiously maintained inverted commas around the word ``absorption'', since the presence of an absorbing agent between ourselves and the source may not be the only interpretation for the emission shifts in B0919+06 and B1850+07. We might equally suppose that different regions of the open emission cone defined by the profile structure illumine at different times, in the manner of ``Christmas lights'', so that the emission flips from one side of the magnetic axis to the other. However, both explanations demand dynamic, irregular changes in the pulsars' magnetospheres--whether in the absorbing or emitting material-- for which no model has yet appeared, and there would seem to be no observational way to distinguish between the two hypotheses.  An interesting point is that both B0919+06 and B1859+07 usually emit late and the ``events'' shift the emission earlier.  Other prominent instances of ``absorption'' show no such asymmetry: as noted above B0943+10's effect usually curtails emission after the central longitude, whereas just the opposite is true for B0809+74.  Thus, we cannot be sure whether there might be other pulsars which could usually emit early and then throw their emission later during ``events''.  Indeed, an in some ways very similar effect is seen in B2034+19, but here the changes are quick, bounded by nulls and are easily classed under the ``mode-change'' rubric (Redman 2006).  It is also possible that in some stars the ``events'' occur much more frequently so that they are seen merely as chaotic subpulse modulation;  possibly pulsar B1604--00 provides such an instance (\\eg, Rankin 1988).  Taking the pulsars B0919+06, B1859+07 and B1604-00 as the most convincing exemplars of the emission shift phenomenon, we attempted to find some common features among their basic physical parameters. Certainly, their spin-down ages give no clue: they range from the ``middle-aged'' B0919+06 at 0.5Myr to B1604-00 at 20Myr. Much the same is true of their inferred surface magnetic fields and their corresponding light-cylinder values. The only possible hint is that their rotation periods are fairly close to one another (0.43s, 0.64s and 0.43s respectively) and somewhat shorter than the average for the general population of ``slow'' pulsars. Physically, this may suggest that the observed $\\it {gradual}$ emission shift requires an intermediate-size light-cylinder for its operation. ",
        "conclusions": "We identify what seems to be a new aspect of pulsar behavior wherein the emission usually comes from a trailing region of the profile, but then occasionally shifts earlier over a few pulses to illuminate the leading part of the profile for some 20--50 pulses, and then shifts back again over a few pulses.  Both B0919+06 and B1859+07 exhibit the effect, the former perhaps once in 2000 pulses and the latter about 10 times more often.  Both stars exhibit asymmetric single profiles which defy ready classification, though in the case of B0919+06 enough structure has been discerned to suggest that it has three main components (Rankin 1993b).  However, when partial profiles are constructed only of the pulses framing the ``events'', more recognizable profile forms are found---and in both cases we have been able to identify some of the underlying component structure and work out their probable basic emission geometry.  It is interesting that B0919+06 is among the stars thought by Lyne \\& Manchester (1988) to represent a class of ``partial cones''---and indeed we certainly find that this classification is appropriate given that only during its ``events'' is the leading part of its cone illuminated.  It will be interesting to see if other pulsars with asymmetric single profiles thus classified by Lyne \\& Manchester also exhibit similar ``events''---and whether other stars exhibiting the effect will be found to throw their emission later rather than earlier.  Our primary interest in these stars and their behaviour, however, lies in the possibility that the ``events'' represent an example of the profile ``absorption'' effect.  In both pulsars it is possible to interpret the emission shift as if emission in the leading part of the profile is normally obscured and then, for brief intervals, the obscuration shifts to the trailing part of the star's profiles. ``Absorption'' appears to exhibit this sort of asymmetry about the longitude of the magnetic axis in other pulsars, and it also exhibits a gradual onset in some other instances. However, we may see the effect in terms of shifting illumination within the defined emission cone, so that the emitting zone appears to shift laterally as different regions switch on and off. Clearly, we do not yet understand physically what is the origin of such observed phenomena, but the starting point must be define its observational characteristics."
    },
    "0601/gr-qc0601096_arXiv.txt": {
        "abstract": "This paper reports on the methods and results of a theoretical analysis to design an insulator which must provide a thermally quiet environment to test on ground delicate temperature sensors and associated electronics. These will fly on board \\esa's \\lisa Pathfinder (\\lpf) mission as part of the thermal diagnostics subsystem of the \\lisa Test-flight Package (\\ltp). We evaluate the heat transfer function (in frequency domain) of a central body of good thermal conductivity surrounded by a layer of a very poorly conducting substrate. This is applied to assess the materials and dimensions necessary to meet temperature stability requirements in the metal core, where sensors will be implanted for test. The analysis is extended to evaluate the losses caused by heat leakage through connecting wires, linking the sensors with the electronics in a box outside the insulator. The results indicate that, in spite of the very demanding stability conditions, a sphere of outer diameter of the order one metre is sufficient. ",
        "introduction": "}  \\lisa Pathfinder (\\lpf) is an \\esa mission, whose main objective is to put to test critical parts of \\lisa (Laser Interferometer Space Antenna), the first space borne gravitational wave (GW) observatory \\cite{bender}. The science module on board \\lpf is the \\lisa Test-flight Package (\\ltp)~\\cite{lpfall}, which basically consists in two test masses in nominally perfect geodesic motion (free fall), and a laser metrology system, which reports on \\emph{residual deviations} of the test masses' actual motion from the ideal free fall, to a given level of accuracy \\cite{gerhar}.  In order to ensure that the test masses are not deviated from their geodesic trajectories by external (non-gravitational) agents, a so called Gravitational Reference System (GRS) is used~\\cite{rita}. This consists in position sensors for the masses which send signals to a set of micro-thrusters; the latter take care of correcting as necessary the spacecraft trajectory, so that at least one of the test masses remains centred relative to the spacecraft at all times. The combination of the GRS plus the actuators is known as \\emph{drag-free} subsystem\\footnote{ The term \\emph{drag-free} dates back to the early days of space navigation, when it was used to name a trajectory correction system designed to compensate for the effect of atmospheric drag on satellites in low altitude orbits.}.  The \\emph{drag-free} is of course a central component of \\lisa, and needs to be operated at extremely demanding levels of accuracy. The laser metrology system should then be sufficiently precise to measure relative test mass deviations. The overall level of noise acceptable for \\lisa is defined in terms of rms acceleration spectral density, and has been set to \\begin{equation} S_{a,{\\rm LISA}}^{1/2}(\\omega)\\leq 3\\!\\times\\!10^{-15}\\,\\left[ 1 + \\left(\\frac{\\omega/2\\pi}{3\\ {\\rm mHz}}\\right)^{\\!\\!2}\\right]\\, {\\rm m}\\,{\\rm s}^{-2}/\\sqrt{\\rm Hz} \\label{eq.1} \\end{equation} in the frequency range $\\quad 10^{-4}\\,{\\rm Hz}\\leq\\omega/2\\pi\\leq 10^{-3}\\,{\\rm Hz}$. This is equivalent to $S_h^{1/2}$\\,$\\sim$\\,4$\\times$10$^{-21}$ Hz$^{-1/2}$, with the same frequency dependence.  Because \\lpf is a \\emph{technological mission}, aimed to assess the feasibility of \\lisa, its ultimate goal has been relaxed to~\\cite{toplev} \\begin{equation} S_{a,{\\rm LPF}}^{1/2}(\\omega)\\leq 3\\!\\times\\!10^{-14}\\,\\left[ 1 + \\left(\\frac{\\omega/2\\pi}{3\\ {\\rm mHz}}\\right)^{\\!\\!2}\\right]\\, {\\rm m}\\,{\\rm s}^{-2}/\\sqrt{\\rm Hz} \\label{eq.2} \\end{equation} in the frequency range $1\\,{\\rm mHz}\\leq\\omega/2\\pi\\leq 30\\,{\\rm mHz}$, i.e., one order of magnitude less demanding, both in noise amplitude and in frequency band.  Equation (\\ref{eq.2}) gives the \\emph{global} noise budget. This is naturally made up of contributions from different perturbative agents, such as temperature and magnetic field fluctuations, GRS and interferometer noise, etc. As a general rule, a requirement on the magnitude of each of the various perturbing factors is set at a 10\\,\\% fraction of the total. In the case of temperature fluctuations, this is equivalent to \\begin{equation} S_{T}^{1/2}(\\omega)\\leq 10^{-4}\\,{\\rm K}/\\sqrt{\\rm Hz}\\ , \\quad 1\\,{\\rm mHz}\\leq \\omega/2\\pi \\leq 30\\,{\\rm mHz} \\label{eq.3} \\end{equation}  Because temperature stability is important, a decision has been taken to place high precision thermometers in several strategic spots across the \\ltp ---as part of what is called \\emph{Diagnostics Subsystem}~\\cite{lobo} \\footnote{ The Diagnostics Subsystem of the \\ltp also includes magnetometric measurements and a charged particle flux detector.}. Such high precision temperature measurements will be useful to identify the fraction of the total system noise which is due to thermal fluctuations only, and this will in turn provide important debugging information to assess the performance of the \\ltp.   \\subsection{Temperature measurements \\label{sec.1-1}}  If the temperature gauges are to be sensitive to fluctuations at the level given by (\\ref{eq.3}) then clearly the entire measuring device should be less noisy, typically by a factor of~10. This means that such device, which includes both the sensors \\emph{and} the associated electronics, can generate a maximum level of noise of \\begin{equation} S_{T, {\\rm sensor}}^{1/2}(\\omega)\\leq 10^{-5}\\,{\\rm K}/\\sqrt{\\rm Hz}\\ , \\quad 1\\,{\\rm mHz}\\leq \\omega/2\\pi \\leq 30\\,{\\rm mHz} \\label{eq.4} \\end{equation}  Research work is currently being conducted at \\ieec (Barcelona, Spain) to identify the appropriate sensors and design the better suited front end electronics. But the prototype system needs of course to be tested for compliance with equation~(\\ref{eq.4}). Thus, in order to do a meaningful test, the system must be sufficiently thermally isolated that the observed fluctuations in the readout data can be attributed \\emph{solely} to sensor noise, rather than to a combination of it with real ambient temperature fluctuations. This means temperature fluctiations in the thermomters' placements should again be at least one order of magnitude below the target sensitivity, equation~(\\ref{eq.4}), or \\begin{equation} S_{T, {\\rm testbed}}^{1/2}(\\omega)\\leq 10^{-6}\\,{\\rm K}/\\sqrt{\\rm Hz}\\ , \\quad 1\\,{\\rm mHz}\\leq \\omega/2\\pi \\leq 30\\,{\\rm mHz} \\label{eq.5} \\end{equation}  It turns out that 10$^{-6}\\,{\\rm K}/\\sqrt{\\rm Hz}$ is a truly demanding temperature stability, orders of magnitude beyond the capabilities of normal thermally regulated rooms. We thus need to design a specific thermal insulator to shield the sensors from ambient temperature fluctuations during the test process.  In the ensuing pages we describe in detail the insulator design. It is extremely important to stress at this point that the performance of the insulator, i.e., its ability to screen out ambient temperature fluctuations, \\emph{cannot be checked} experimentally, at least under working thermal conditions in the laboratory. This is because, by definition, the insulator is the tool to check the sensing instruments, \\emph{not viceversa}: we need to rely on the results of theoretical argumentation to make a decission on which is the appropriate thermal insulator for our purposes. An experimental verification of the model is only thinkable under much more extreme conditions, where external temperature fluctuations are orders of magnitude higher than the ones which will be met during the test. ",
        "conclusions": "Temperature fluctuation measurement is very demanding in the \\ltp, and subsequently \\lisa, as reflected by equation~\\eref{eq.4}. Accordingly, very delicate sensor and associated electronics must be designed, and of course tested in ground before boarding.  However, even the best laboratory conditions are orders of magnitude worse than the above requirement, so meaningful tests of the temperature sensing system cannot be tested without suitably screening the sensors from ambient temperature fluctuations. We have addressed how this can be accomplished by means of an insulating system consisting of a central metallic core surrounded by a thick layer of a very poorly conducting material. The latter provides good thermal insulation, while the central core, having a large thermal inertia, ensures stability of the sensors' environemnt. The choice of materials is flexible, so aluminium and polyurethane, which are easily available in the market, has been adopted. Thereafter, the dimensions need to be fixed.  The appropriate sensors for the needs are temperature sensitive resistors, more specifically thermistors ---also known as NTCs. It appears that, because these sensors need to be wired to external electronics, heat leakage through such wires is an effect which needs to be quantitatively assessed to prevent losses. We have analysed this problem, and concluded that it strongly depends on the central metallic core size, and imposes that it be somewhat large.  Laboratory ambient temperature fluctuations, determined by dedicated \\emph{in situ} measurements, are of the order of 10$^{-1}$\\,K/$\\sqrt{\\rm Hz}$ at 1~mHz, and dropping at higher frequencies within the MBW. The required stability conditions at the sensors, attached at the core's surface, thus need an attenuation factor of 10$^{-5}$, or better. Our analysis determines that a central aluminium core of 13~cm of radius, surrounded by a concentric layer of polyurethane 15--20~cm thick, comfortably provides the needed thermal screening which guarantees a meaningful test of the sensors' performance.  The results of this paper are based on modelling. Because our aim is to produce a very stable thermal environment for the temperature sensors, we cannot check \\emph{experimentally} the correctness of our conclusions. We must instead rely on the validity of the hypotheses made ---essentially that heat only flows by thermal conduction--- and on the underlying physical laws which govern heat conduction. Even though there is good reason to believe that both are sufficiently accurate, unexpected behaviour e.g. at the interface between the metal core and the insulator, may partly distort the results. Direct measurements with very large temperature gradients applied across the insulating device are envisaged, and will be reported elsewhere as an auxiliary independent test of the model.  \\ack  We want to thank Albert Tom\\`as, from {\\sl NTE}, for discussions on the subject of this paper. Support for this work came from Project ESP2004-01647 of Plan Nacional del Espacio of the Spanish Ministry of Education and Science (MEC). MN acknowledges a grant from Generalitat de Catalunya, and JS a grant from MEC.     \\appendix"
    },
    "0601/astro-ph0601532_arXiv.txt": {
        "abstract": "We mapped 12 massive protostellar candidates in the \\coj\\ line, which in combination with Zhang et al. (2005) completes an unbiased survey of outflows for all 48 sources with $l$$>$50\\degr\\ in a sample of 101 massive protostellar candidates. We detected outflows in 10 sources, implying 88\\% occurrence frequency of outflows for the 48 sources. This supports the conclusion of previous studies that bipolar outflows are an integral component in the formation process of massive stars. The vast majority of the observed outflows are much more massive ($>$10~\\msol) and energetic ($>$100~\\msol~\\kms) than outflows from low-mass protostars. They also have large mass outflow rates ($>$2$\\times$10$^{-4}$~\\msol~yr$^{-1}$), suggesting large ($\\sim$1$\\times$10$^{-4}$~\\msol~yr$^{-1}$) accretion rates sufficient to overcome radiation pressure of the central massive protostars. We compared the frequency distribution of collimation factors of 40 massive outflows including those of this study with that of 36 low-mass outflows from the literature, and found {\\it no} significant difference between the two. All these results are consistent with the suggestion that massive stars form through accretion as do low-mass stars but with much higher accretion rates. ",
        "introduction": "It has been well established that low-mass stars form by gravitational collapse of dense molecular cores and that accretion disks and bipolar outflows are basic building blocks of the process \\citep{shu87}. Compared to low-mass star formation, very little is known about the formation process of massive ($>$8~M$_\\odot$) stars, such that it is still much debated whether  massive stars form in a qualitatively similar manner to low-mass stars. This debate is mainly because the radiation pressure of massive stars is strong enough to reverse materials infalling at the typical mass accretion rates observed in low-mass young stellar objects (YSOs) \\citep{wc87}. There are currently two major theoretical models in competition: accretion via disks \\citep[e.g.,][]{ys02, mt03} and coalescence of low-mass (proto)stars \\citep[e.g.,][]{bon98}. The most convincing way to differentiate between the two models may be to investigate whether well-defined accretion disks exist around massive YSOs. But it is difficult to detect the accretion disks, if any, because they would be small (at most several hundred AU), short-lived, and easily confused by envelopes. In this study we attempt to distinguish between the two models by observing bipolar molecular outflows that are much easier to detect.  Bipolar outflows from low-mass YSOs have been extensively studied since the early 1980's \\citep[see][and references therein]{bac96}, while systematic studies of outflows associated with massive star formation began much later. These studies are roughly divided into two groups. The first group are single-point CO line surveys towards samples of massive YSOs in search of high-velocity molecular gas \\citep[e.g,][]{sc96a,sri02}. They showed that high-velocity gas is a common feature in massive YSOs. \\citet{sc96a} made a CO J=1$-$0 line survey towards 122 massive YSOs, the vast majority of which are ultracompact (UC) HII regions, and detected high-velocity gas in 90\\% of the 94 sources suitable for line wing analysis. \\citet{sri02} obtained a similar detection rate (84\\%) from a CO J=2$-$1 line survey of 69 massive protostellar candidates. The second group are CO line mappings of subsets of the sources that exhibit high-velocity wings \\citep[e.g.,][]{sc96b,beu02a}. These studies investigated whether the high-velocity wings are caused by bipolar outflows and derived the physical and dynamical parameters of the outflows. \\citet{sc96b} mapped 10 sources with high-velocity gas and identified outflows in 5 sources. If we assume a similar detection rate for the remainder of their sample, the occurrence frequency of outflows would be about 45\\% for their sample. \\citet{beu02a} found outflows in 21 of 26 sources showing high-velocity wings in the sample of \\citet{sri02}, indicating a significantly higher occurrence rate ($\\sim$70\\%). These studies also show that bipolar outflows from massive YSOs are much more massive and energetic than those observed in low-mass YSOs \\citep[see also][]{rm01}. There is some controversy as to whether massive outflows are much less collimated than low-mass outflows (see \\S~3.1).  The present paper presents \\coj\\ line maps of 12 massive star-forming regions at 27$''$ angular resolution (Table 1). Combining the results of this study with those of \\citet[][2005]{zha01}, we complete a census of molecular outflows for all sources with $l$$\\ge$50$^\\circ$ in a flux-limited sample of massive protostellar candidates (see \\S~2.1), and estimate the occurrence frequency of molecular outflows for the sample. We also examine the morphological, physical, and dynamical properties of bipolar outflows observed, and compare them with those of outflows from low-mass YSOs. The selection criteria of the sources and the details of the observations are described in \\S~2. The results are presented in \\S~3. The implications of our results for massive star formation are discussed in \\S~4. We summarize the main results in \\S~5. ",
        "conclusions": "We mapped 12 massive protostellar candidates in the \\coj\\ line, which in combination with \\citet{zha05} completed a census of bipolar molecular outflows for all 48 sources with $l$$>$50$^\\circ$ in the Molinari catalog of 101 massive protostellar candidates. We detected outflows in 10 of the 12 sources, indicating 88\\% occurrence frequency of outflows in the 48 sources. The observed outflows are much more massive ($>$10~\\msol) and energetic ($>$100~\\msol~\\kms) than those observed in low-mass YSOs. These results imply that almost all sources of the Molinari catalog have very massive and powerful molecular outflows. Thus bipolar outflows appear to be a ubiquitous phenomenon in massive star formation as in low-mass star formation, as suggested by some previous studies.  Nearly all outflows observed by this study have mass outflow rates $>$2$\\times$10$^{-4}$~\\msol~yr$^{-1}$, suggesting large accretion rates of $\\sim$1$\\times$10$^{-4}$~$M_\\odot$~yr$^{-1}$. These values are much greater than the typical accretion rates expected and measured in low-mass YSOs and are large enough to overcome the radiation pressure of the central massive protostars. We compared the collimation factors of 40 massive outflows with those of 36 low-mass outflows, and found {\\it no} significant difference between them. This can be naturally understood in the accretion model, but cannot be easily explained by the coalescence model. Therefore, it is likely that accretion plays a dominant role in the formation process of massive stars, although we cannot exclude the possibility that coalescence plays a role in some special circumstances.   We acknowledge support from the Laboratory for Astronomical Imaging at the University of Illinois and NSF AST 0228953."
    },
    "0601/astro-ph0601704_arXiv.txt": {
        "abstract": "We study the emission from an old supernova remnant (SNR) with an  age of around $10^5$~yrs and that from a giant molecular cloud (GMC) encountered by the SNR. When the SNR age is around $10^5$~yrs, proton acceleration is efficient enough to emit TeV $\\gamma$-rays both at  the shock of the SNR and that in the GMC. The maximum energy of primarily accelerated electrons is so small that TeV $\\gamma$-rays and  X-rays are dominated by hadronic processes, $\\pi^0$-decay and synchrotron radiation from secondary electrons, respectively. However, if the SNR is older than several $10^5$~yrs, there are few high-energy particles emitting TeV $\\gamma$-rays because of the energy loss effect and/or the wave damping effect occurring at low-velocity isothermal shocks. For old SNRs or SNR-GMC interacting systems capable of generating TeV $\\gamma$-ray emitting particles, we calculated the ratio of TeV $\\gamma$-ray (1--10~TeV)  to X-ray (2--10~keV) energy flux and found that it can be more than $\\sim10^2$. Such a source showing large flux ratio may be a possible origin of recently discovered unidentified TeV sources. ",
        "introduction": "\\label{sec:intro}  The most probable cosmic-ray accelerator in our Galaxy is the young supernova remnant (SNR). The detection of synchrotron X-rays from shells of young SNRs provides us the strong evidence for electron acceleration up to more than $\\sim10$~TeV \\citep[e.g.,][]{koyama1995,koyama1997, bamba2003a,bamba2003b,bamba2005a,bamba2005b}. So far, the evidence for hadron acceleration, however, has not yet been obtained. High energy $\\gamma$-ray observations may give us important information on the accelerated protons \\citep{naito1994,drury1994,aharonian1994,aharonian1996}. For example, TeV $\\gamma$-rays are detected from the SNRs RX~J1713.7$-$3946 \\citep{enomoto2002,aharonian2004} and RX~J0852.0$-$4622 \\citep{katagiri2005,aharonian2005c}, which can be originated in either the decay of neutral pion, arising from the collision of high energy protons and interstellar matter, or CMB photons up-scattered by accelerated electrons \\citep[e.g,][]{pannuti2003,lazendic2004, ellison2001,bamba2005b,uchiyama2005}. At present, the leptonic process is not yet ruled out.   Recently, a survey of the inner part of our Galaxy has revealed several new TeV $\\gamma$-ray sources \\citep{aharonian2002,aharonian2005,aharonian2005b,aharonian2005d}. For some of them, no counterpart has been found in any other wave lengths yet. Their spectra are rather hard, and the flux above 1~TeV is around $1\\times10^{-11}$~cm$^{-2}$s$^{-1}$. They are extended with the angular size of around $0.1^\\circ$. They should be galactic origin because all are located along the galactic plane. In Table~\\ref{table}, we list the properties of the Galactic TeV $\\gamma$-ray sources that are also observed in the X-ray bands (together with a typical young SNR SN~1006). One of the newly discovered sources, HESS~J1303$-$631 \\citep{aharonian2005b}, was observed by {\\it Chandra}, and no obvious counterpart was revealed \\citep{mukherjee2005}. Then, the flux ratio, defined by \\begin{equation} R_{\\rm TeV/X}=\\f{F_\\gamma(1-10~{\\rm TeV})}{F_X(2-10~{\\rm keV})}~~, \\end{equation} is more than $\\sim2$. On the other hand,  young SNRs, Cas~A, RX~J1713.7$-$3946 and RX~J0852.0$-$4622, have $R_{\\rm TeV/X}$ less than $\\sim2$. Although at present the lower limit on $R_{\\rm TeV/X}$ for HESS~J1303$-$631 is comparable to that of young SNRs, it may become much larger with forthcoming deeper X-ray observations. Furthermore, HESS~J1640$-$465 has an X-ray counterpart that is identified as SNR \\citep{aharonian2005d}. This object has $R_{\\rm TeV/X}\\sim19$. Unlike young SNRs, these TeV sources with large $R_{\\rm TeV/X}$ may show  evidence for hadron acceleration because inverse-Compton scenario requires unusually small magnetic field strength ($\\ll1~\\mu$G). \\begin{table*} \\centering \\begin{minipage}{140mm} \\caption{ Observed properties of Galactic TeV $\\gamma$-ray sources and a typical young SNR, SN~1006.} \\label{table} \\begin{tabular}{cccccc} \\hline Name & $F_\\gamma$(1--10~TeV)% \\footnote{Derived for best fitted values (in units of $10^{-12}$erg~s$^{-1}$cm$^{-2}$).} & $F_X$(2--10~keV)% \\footnote{Derived for best fitted values (in units of $10^{-12}$erg~s$^{-1}$cm$^{-2}$).} & $R_{\\rm TeV/X}$% \\footnote{$R_{\\rm TeV/X}=F_\\gamma$(1--10~TeV$)/F_X$(2--10~keV)~.} & Classification% \\footnote{SNR: Supernova Remnant, PWN: Pulsar Wind Nebula, unID: unidentified TeV $\\gamma$-ray source.} & References% \\footnote{ (1)~\\citet{aharonian2004b}; (2)~\\citet{willingale2001}; (3)~\\citet{aharonian2005f}; (4)~\\citet{delaney2006} (5)~\\citet{aharonian2001}; (6)~\\citet{allen1997}; (7)~\\citet{aharonian2005e}; (8)~\\citet{ozaki1998}; (9)~\\citet{aharonian2004}; (10)~\\citet{pannuti2003}; (11)~\\citet{aharonian2005c}; (12)~\\citet{slane2001}; (13)~\\citet{aharonian2005d}; (14)~\\citet{brogan2005}; (15)~\\citet{sugizaki2001}; (16)~\\citet{aharonian2002}; (17)~\\citet{aharonian2005b}; (18)~\\citet{mukherjee2005}. } \\\\ \\hline Crab              & $56$   & $2.1\\times10^{4}$  & $2.6\\times10^{-3}$ & PWN  & (1), (2)  \\\\ RCW~89            & $16$   & $38$             & $0.41$  & PWN  & (3), (4)  \\\\ Cas~A             & $1.9$  & $1.3\\times10^2$  & $0.014$ & SNR  & (5), (6)  \\\\ SN~1006 (NE)      & $<2.1$ & $19$             & $<0.11$ & SNR  & (7), (8)  \\\\ RX~J1713.7$-$3946 & $35$   & $5\\times10^2$    & $0.07$  & SNR  & (9), (10)  \\\\ RX~J0852.0$-$4622 & $69$   & $>32$            & $<2.1$  & SNR  & (11), (12)  \\\\ HESS~J1813$-$178  & $8.9$  & $7.0$            & $1.3$   & SNR  & (13), (14) \\\\ HESS~J1640$-$465  & $7.1$  & $0.37$           & $19$    & SNR  & (13), (15) \\\\ TeV~J2032$+$4130  & $1.9$  & $<2$             & $>1$    & unID & (16)       \\\\ HESS~J1303$-$631  & $10$   & $<6$             & $>2$    & unID & (17), (18) \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table*}   A large value of $R_{\\rm TeV/X}$ is expected for old SNRs. As the SNR ages,  the shock velocity decreases. In general, primary electron acceleration is limited by synchrotron cooling. Then, the roll-off energy of electron synchrotron radiation is much smaller than that of typical young SNRs, so that  small synchrotron X-ray flux is expected \\citep[e.g.,][]{sturner1997}. It is also important to consider the association with a giant molecular cloud (GMC). Because of large volume, old SNRs may encounter the GMC. In this paper, we study how large $R_{\\rm TeV/X}$ can become for single SNRs and the SNR-GMC interacting systems. ",
        "conclusions": "\\label{sec:dis}  For an old SNR in the radiative phase, the maximum energy of primary electrons is so small that emissions via leptonic processes are vanishingly small both in the X-ray and TeV $\\gamma$-ray bands. On the other hand, there still exists  accelerated protons with energies of more than $\\sim10$~TeV. These particles emit the TeV $\\gamma$-rays via $\\pi^0$-decay and synchrotron X-rays from secondary electrons generated by charged pions. We find that  the ratio \\[ R_{\\rm TeV/X}=\\f{F_\\gamma(1-10~{\\rm TeV})}{F_X(2-10~{\\rm keV})} \\] could be more than $\\sim10^2$. Such  sources may be an origin of recently discovered unidentified TeV sources, and give us the evidence for hadron acceleration. We might have to consider the interaction between the GMC and the SNR with an age of $\\sim10^5$~yrs. For $t_{\\rm age}\\ll10^5$~yrs, SNR radius is so small that there is only a few SNRs interacting with a GMC. On the other hand, if the SNR with  $t_{\\rm age}\\gg10^5$~yrs encounters the GMC, there are few high-energy particles emitting TeV $\\gamma$-rays around shocks of the SNR and the GMC  because of the energy loss effect and/or the wave damping effect occurring at low-velocity isothermal shocks.  Actually detected TeV sources have the energy flux $F_\\gamma(1-10~{\\rm TeV})\\sim10^{-12}$--$10^{-11}$erg~cm$^{-2}$s$^{-1}$ \\citep{aharonian2005,aharonian2005d}. Hence if $R_{\\rm TeV/X}\\la10^{2-3}$, then $F_X(2-10~{\\rm keV})\\ga10^{-14}$~erg~cm$^{-2}$s$^{-1}$. Such diffuse, extended source can be detected with current X-ray telescopes ({\\it Suzaku}, {\\it Chandra}, and {\\it XMM-Newton}). Especially, X-ray Imaging Spectrometer (XIS) onboard {\\it Suzaku} is capable of observing dim, diffuse X-ray sources with low background in the hard X-ray band. Less energetic protons emit GeV $\\gamma$-rays that might have been detected by EGRET or may be detected by GLAST in the future. However, for isothermal shocks, the spectrum of accelerated particles may be hard ($p<2$). Indeed, some of the newly discovered TeV sources  show hard spectrum \\citep{aharonian2005,aharonian2005d}, which may imply $p\\lesssim2$. In such a case, GeV emission becomes dim as shown in Table~\\ref{table2}, where we calculate $R_{\\rm TeV/GeV}$ for the case of Figures~1--3, 5 and 6. Radio emission is also expected. However, as can be seen in Table~\\ref{table2}, the sources with large $R_{\\rm TeV/X}$ have large $R_{\\rm TeV/radio}$, so that the radio emission may be dim. This fact might account for a large scatter of $R_{\\rm TeV/radio}$ found by \\citet{helfand2005}. Furthermore, the source is confined in a galactic plane region and radio emission may be obscured by  diffuse components.   \\begin{table} \\centering \\begin{minipage}{140mm} \\caption{ Calculated flux ratios for various bands.} \\label{table2} \\begin{tabular}{cccc} \\hline & $R_{\\rm TeV/X}$% \\footnote{$R_{\\rm TeV/X}=F_\\gamma$(1--10~TeV$)/F_X$(2--10~keV)~.} & $R_{\\rm TeV/GeV}$% \\footnote{$R_{\\rm TeV/GeV}=F_\\gamma$(1--10~TeV$)/F_{\\rm GeV}$(1--10~GeV)~.} & $R_{\\rm TeV/radio}$% \\footnote{$R_{\\rm TeV/radio}=F_\\gamma$(1--10~TeV$) /F_{\\rm radio}$($10^7$--$10^{11}$~Hz)~.} \\\\ \\hline Fig.~1 & $7.6\\times10^{-2}$ & $0.36$  & $9.8$            \\\\ Fig.~2 & $1.6$              & $0.25$  & $6.6$            \\\\ Fig.~3 & $82$               & $6.6$   & $47$             \\\\ Fig.~5 & $6.9$              & $74$    & $9.7\\times10^3$  \\\\ Fig.~6 & $5.9\\times10^2$    & $5.8$   & $3.5\\times10^2$  \\\\ \\hline \\end{tabular} \\end{minipage} \\end{table}  In this paper, we consider TeV $\\gamma$-ray emission from (1) a single old SNR with $t_{\\rm age}\\sim3\\times10^5$~yrs (\\S~\\ref{subsec:emissionSNRold}), (2) a shocked GMC collided by the old SNR with $t_{\\rm age}\\sim5\\times10^4$~yrs (\\S~\\ref{subsec:emissionGMC}), and (3) an unshocked GMC illuminated by high-energy particles arising at the shock of the old SNR with $t_{\\rm age}\\sim3\\times10^5$~yrs (\\S~\\ref{subsec:emissionGMC2}), respectively. As in the followings, we can roughly estimate, for each case, the expected number of observed TeV sources, $N_{\\rm TeV}$, which have observed flux larger than 3\\% of the Crab flux ($6\\times10^{-13}$~cm$^{-2}$s$^{-1}$ above 1~TeV) and lie in the inner part of the galactic plane with $-30^\\circ<l<30^\\circ$ \\citep[cf.][]{aharonian2005,aharonian2005d}. The number of  detected TeV sources  gives us important information on the energetics of old SNR emission.  Let us first consider the case~(1). Let $E_{p}=E_{p,50}10^{50}$~ergs be the energy of accelerated protons stored in the SNR. Then, the observed flux of TeV $\\gamma$-rays is calculated as $F(>{\\rm TeV})\\sim1\\times10^{-10} n_{{\\rm sh,} 0.86}E_{p,50} d_{\\rm kpc}^{-2}$~cm$^{-2}$s$^{-1}$, where $n_{\\rm sh}=10^{0.86}n_{{\\rm sh,} 0.86}$~cm$^{-3}$ and $d=d_{\\rm kpc}$~kpc are the density of the SNR shell and the distance to the source, respectively. The maximum distance to the source is $d_{\\rm max}\\sim13n_{{\\rm sh,} 0.86}^{1/2}E_{p,50}^{1/2}$~kpc. Then the volume fraction of the survey region is \\[ \\beta_1\\sim\\frac{(\\pi/6)d_{\\rm max}{}^2}{\\pi(10~{\\rm kpc})^2} \\sim0.30~n_{{\\rm sh,} 0.86}E_{p,50}~~. \\] Hence, for assumed SN explosion rate of $1\\times10^{-2}r_{-2}$~yr$^{-1}$, we obtain \\begin{eqnarray} N_{\\rm TeV}&\\sim&3\\times10^3~r_{-2}\\beta_1\\nonumber\\\\ &\\sim&9\\times10^2~r_{-2}n_{{\\rm sh,} 0.86}E_{p,50}~~.\\nonumber \\end{eqnarray}    Next, we calculate $N_{\\rm TeV}$ for case~(3). In order to estimate the  number of GMCs in our Galaxy, we adopt the mass function of GMCs, $f(M_c)\\propto M_c^{-1.5}$ with the maximum and the minimum mass, $M_{\\rm max}=2\\times10^7M_{\\sun}$ and $M_{\\rm min}=30M_{\\sun}$, respectively \\citep{blitz2004}. If the total mass of GMCs in our Galaxy is $2\\times10^9M_{\\sun}$, then the number of GMCs with mass of around $M_c=10^5M_5M_{\\sun}$ is about $7\\times10^2M_5^{-1/2}$. The SNR with an age of $3\\times10^5$~yrs has a radius $R_s\\sim90$~pc, so that the geometrical factor, representing the fraction of accelerated protons colliding the GMC and emitting $\\gamma$-rays, is estimated as $\\alpha=(R_c/R_s)^2/4\\sim1\\times10^{-2}$. The observed flux of TeV $\\gamma$-rays is calculated as $F(>{\\rm TeV})\\sim3\\times10^{-11}\\alpha_{-2} n_{{\\rm c,}2.2}E_{p,50} d_{\\rm kpc}^{-2}$~cm$^{-2}$s$^{-1}$, where $\\alpha=\\alpha_{-2}10^{-2}$, and $n_{\\rm c}=10^{2.2}n_{{\\rm c,}2.2}$~cm$^{-3}$ is the density of the GMC. The maximum distance to the source is $d_{\\rm max}\\sim6\\alpha_{-2}^{1/2}n_{{\\rm c,}2.2}^{1/2}E_{p,50}^{1/2}$~kpc. Then the volume fraction of the survey region is \\[ \\beta_1\\sim\\frac{(\\pi/6)d_{\\rm max}{}^2}{\\pi(10~{\\rm kpc})^2} \\sim7\\times10^{-2}\\alpha_{-2}n_{{\\rm c,}2.2}E_{p,50}~~. \\] Assuming the volume of the galactic disk of $4\\times10^{10}$~pc$^3$, the mean separation of GMCs is $\\ell\\sim4\\times10^2M_5^{1/6}$~pc. Then, the probability that a SNR collides with a GMC is \\[ \\beta_2\\sim\\f{4\\pi}{3}\\left(\\f{R_s}{\\ell}\\right)^3 \\sim5\\times10^{-2} \\ell_{400}{}^{-3} \\left(\\f{R_s}{90~{\\rm pc}}\\right)^3~~, \\] where $\\ell=400~\\ell_{400}$~pc. Therefore, we derive \\begin{eqnarray} N_{\\rm TeV}&\\sim&3\\times10^3~r_{-2}\\beta_1\\beta_2\\nonumber\\\\ &\\sim&10~r_{-2}\\alpha_{-2}n_{{\\rm c,}2.2}E_{p,50} \\ell_{400}{}^{-3} \\left(\\f{R_s}{90~{\\rm pc}}\\right)^3~~. \\nonumber \\end{eqnarray} Hence $E_{p,50}\\gtrsim0.1$ is required so that $N_{\\rm TeV}\\gtrsim1$.   Finally, we consider the case~(2). We assume that  the energy of accelerated protons is $1\\times10^{48}E'_{p,48}$~ergs, which is much smaller than the thermal energy stored in the shocked GMC. Then, the observed  TeV $\\gamma$-ray flux is calculated as $F(>{\\rm TeV})\\sim3\\times10^{-9}n_{{\\rm c,d,}4.1}E'_{p,48} d_{\\rm kpc}^{-2}$~cm$^{-2}$s$^{-1}$, where $n_{\\rm c,d,}=10^{4.1}n_{{\\rm c,d,}4.1}$~cm$^{-3}$ is the density of shocked GMC. The maximum distance to the source is calculated as $d_{\\rm max}\\sim66n_{{\\rm c,d,}4.1}^{1/2}E'_{p,48}{}^{1/2}$~kpc, which is larger than the size of the Galactic disk. Therefore, the volume fraction is  about $\\beta_1\\sim0.3$. Since the SNR with an age of $5\\times10^4$~yrs has a radius $R_s\\sim40$~pc, the collision probability is \\[ \\beta_2\\sim\\f{4\\pi}{3}\\left(\\f{R_s}{\\ell}\\right)^3 \\sim4\\times10^{-3} \\ell_{400}{}^{-3} \\left(\\f{R_s}{40~{\\rm pc}}\\right)^3~~. \\] Hence we obtain \\begin{eqnarray} N_{\\rm TeV}&\\sim&5\\times10^2~r_{-2}\\beta_1\\beta_2\\nonumber\\\\ &\\sim&0.6~r_{-2}(\\beta_1/0.3)\\ell_{400}{}^{-3} \\left(\\f{R_s}{40~{\\rm pc}}\\right)^3~~.\\nonumber \\end{eqnarray} Taking into account for the uncertainties for $R_s$ and/or $\\ell_{400}$, we can expect $N_{\\rm TeV}\\gtrsim1$ for reasonable values of parameters.  At present, the number of unidentified TeV sources showing large $R_{\\rm TeV/X}$ is only a few. If the number of actually detected TeV sources does not increase (i.e., the H.E.S.S. galactic plane survey with the flux larger than 0.03~Crab above 1~TeV is nearly complete at present), then  $E_{p,50}\\lesssim0.01$ is implied, otherwise $N_{\\rm TeV}$ for the case (1) is much larger than a few. This might suggest the escape of high-energy particles from the SNR starts by  $t_{\\rm age}\\sim10^5$~yrs, or injection efficiency is low for old SNRs. Then, the number of TeV sources, $N_{\\rm TeV}$, for the case (3) should be small. On the other hand, the case (2) may be still likely because the number of sources for this case can be comparable to that for the case (1). Hence, some TeV sources are associated with the GMC, but others not. In order to clarify whether the TeV source is associated with a GMC, the CO observation may be important."
    },
    "0601/astro-ph0601548_arXiv.txt": {
        "abstract": "The reliability of modeling the far-IR continuum to $\\cOone$ spectral line ratios applied to the Orion clouds \\citep{W05} is tested by applying the models to simulated data.  The two-component models are found to give the dust-gas temperature difference, $\\DT$, to within 1 or 2$\\,$K.  However, other parameters like the column density per velocity interval and the gas density can be wrong by an order of magnitude or more.  In particular, the density can be systematically underestimated by an order of magnitude or more.  The overall mass of the clouds is estimated correctly to within a few percent.  The one-component models estimate the column density per velocity interval and density within factors of 2 or 3, but their estimates of $\\DT$ can be wrong by 20$\\,$K.  They also underestimate the mass of the clouds by 40-50\\%.  These results may permit us to reliably constrain estimates of the Orion clouds' physical parameters, based on the real observations of the far-IR continuum and $\\cOone$ spectral line.  Nevertheless, other systematics must be treated first.  These include the effects of background/foreground subtraction, effects of the HI component of the ISM, and others. These will be discussed in a future paper \\citep{W05a}. ",
        "introduction": "}  Paper~1 \\citep{W05} examined the ability of the FIR-continuum to $\\cOone$ line intensity ratio to diagnose dust and molecular gas physical conditions. Specifically, the {\\it COBE/DIRBE\\/} 140$\\um$ and 240$\\um$ continuum data \\citep[see][]{dirbex} were compared with the Nagoya 4-m $\\cOone$ spectral line data for the Orion~A \\citep{Nagahama98} and B molecular clouds.  The $\\Ib/\\Ic$ ratio, or $\\rd$, was plotted against the 140$\\um$/240$\\um$ dust color temperature, or $\\Tdc$, for the high signal-to-noise positions ($\\ge 5-\\sigma$ for 140$\\um$, 240$\\um$, and $\\cOone$) in the Orion clouds.  This plot was modeled with LTE and LVG, one-component models and LVG, two-component models; the two-component models fit the data better than the one-component models at the 99.9\\% confidence level.  Tables~1, 2, and 3 of Paper~I list the resultant parameter values of the two-component model fits.  The most noteworthy result is that the two-component models demand the dust-gas temperature difference, $\\DT$, to be zero within $\\pm 1$ or 2$\\,$K. (Note that in the case of the two-component, two-subsample models, the $\\Tdc\\ge 20\\,$K subsample still yields $\\DT = 0\\pm 1K$ {\\it if\\/} a two-component model is fitted to that subsample.  The listed results in Table~2 of Paper~I are those of the one-component model fitted to the $\\Tdc\\ge 20\\,$K subsample.)  This result has important consequences that were briefly mentioned in Paper~I and will be discussed in detail in Paper~III \\citep{W05a}. Consequently, the reliability of the derived $\\DT$ must be tested.  In all of the modeling mentioned in Paper~I, the systematic uncertainties of the derived parameter values were evaluated by applying scale factors to the data. These systematic uncertainties are related to uncertainties in the calibration and in certain assumptions, such as the dust optical depth to gas column density ratio. The combined effect of these uncertainties was estimated to be $\\pm 40\\%$.  Accordingly, scale factors that varied from 0.6 to 1.4 were applied to the data to see how strongly the resultant parameter values would change.  Also, the starting search grid for the two-component models was slightly shifted and re-run.  The magnitudes of the changes in the results provided another test of the systematic uncertainties in the parameter values.  These two tests gave similar estimates of the systematic uncertainties. These systematic uncertainties are demonstrated in Figure~21 of Paper~I, which shows that the column densities per velocity interval and densities of both components are uncertain by factors of a few or by more than an order of magnitude.  (These uncertainties are orders of magnitude larger than the formal uncertainties obtained from the model fits. Accordingly, the latter uncertainties can be ignored.)  While the abovementioned tests provide rough estimates of the reliability of the results, they do {\\it not\\/} measure any biases inherent in the method.  In other words, the range of possible parameter values that result from the modeling and from the tests {\\it may not even include the ``true'' or correct value.\\/}  And we cannot know that these ranges are indicative of the correct values,  because we cannot know the correct values in the first place.  This is in stark contrast to using simulated data.  With simulated data, the true, or input, values can be compared with the resultant values from the model fits.  The tests that were applied to modeling the actual observed data can be repeated on the modeling of the simulated data.  Biases or shortcomings in the modeling technique are then clearly seen.  In the following section and its subsections, the creation of the simulated data and the results of modeling these data is described.  Other systematics are not discussed in the current paper, but are left to Paper~III.  These are the systematic effects that result when the models do not properly characterize the contributions of other phases of the ISM, such as from HI and its dust or from some large-scale foreground/background emission, or when they adopt an improper value of some more basic physical parameter, such as the far-IR spectral emissivity index, $\\beta$. ",
        "conclusions": "}  The reliability of recovering physical conditions in the dust and gas of molecular clouds using the far-IR continuum and the $\\cOone$ line was tested by using simulated data.  These data were created using input beam-average column density and dust temperature maps that crudely represented the inferred physical conditions in the Orion~A and B giant molecular clouds \\citep[see Paper~I][]{W05}.  Input physical parameters, with values similar to those recovered from modeling the actual observed data (see Paper~I), in combination with the column density and dust temperature maps gave us the simulated intensity maps in the 140$\\um$ continuum, 240$\\um$ continuum and $\\cOone$ spectral line.  The simulated maps assumed two subsamples of positions within the clouds and two components.  The two components were component 0, with constant physical conditions within each subsample, and component 1, with constant physical conditions within each subsample, except for spatially varying dust and gas temperatures.  The two subsamples were defined by the component-1 dust temperature, $\\Tdo$: those positions with $\\Tdo<20\\,$K represent one subsample and the positions with $\\Tdo\\ge 20\\,$K represent the other subsample.  The point of the current paper was to apply the models used in Paper~I to the simulated maps to see how well those models recover the input values of the physical parameters.  Given that the simulated maps are based on the two-component, two-subsample models, fitting such models to the simulated data in the noise-free case {\\it might\\/} be expected to recover the inputs perfectly.  However, even in the noise-free case some input parameters could {\\it not\\/} be recovered.  The component-0 and component-1 densities, for example, were an order of magnitude or more different from the inputs. The simulated maps with noise show us that we can obtain the dust-gas temperature difference, $\\DT$, to within 1 or 2$\\,$K {\\it regardless of the specific value of $\\DT$.\\/}  The component zero dust temperature is apparently recovered to within a fraction of a Kelvin, but see Paper~III for further discussion of this.  Recovery of the component-1 dust temperatures and the gas column densities is accurate to within a few percent for 93\\% of the points.  The other 7\\% of the points have column densities too high or too low by a factor of about 5.  This results in overestimate of only 6-7\\% in the total mass.  The simple two-component models applied to the simulations has shown what biases can exist in the model results: \\begin{enumerate} \\item There are noticeable systematic offsets in the derived component-1 dust temperatures and in the derived column densities from their inputs.  These offsets come from forcing a single model curve to fit through the two different subsamples. \\item About 7\\% of the column densities are wrong by factors of 5, as is the case for the two-subsample, two-component models.  Inspite of these systematic errors, the simple two-component models overestimate the total mass of the clouds by only 3 to 6\\%. \\item Despite the varying the scale factors, the inferred component-1 densities are {\\it all\\/} systematically too low by an order of magnitude or more from the input density. \\end{enumerate} Keeping these shortcomings in mind gives us reasonable estimates of the parameter value ranges for all the two-component models as applied to the real observations.  These ranges are listed in Table~\\ref{tbl-8}.  The range for the component-1 density is the kind of range roughly expected for LTE emission of the $\\cOone$ line.  The range for the component-1 column density per velocity interval is given, as stated earlier, by the large-scale cloud properties at the low end and by the necessity of optically thin $\\cOone$ emission at the high end.  For component 0, the lower limit is nearly that of the master search grid. (Note that the $c_0\\nvtcz$ product lower limit is indeed {\\it equal\\/} to that of the master search grid, but the $\\nvtcz$ {\\it itself\\/} is still slightly larger than that.)  Fitting the one-component models to the simulated data shows that these models can provide reasonable estimates of the column density per velocity interval and volume density (i.e., within factors of 2 or 3) provided that these models are applied to the two different subsamples (i.e. with $\\Td$ below and above 20$\\,$K). The estimates of $\\DT$, however, can be wrong by about 20$\\,$K.   The one-component models result in mass estimates that are too low by about 40-50\\%; the continuum-derived mass estimates being worse on average than the $\\cO$-derived mass estimates due to the higher temperature sensitivity of the continuum observations.   These simulations have provided important insights into the reliability of the model results. Yet other questions need to be addressed: \\begin{itemize} \\item What is the effect of the background subtractions used? \\item How will dust associated with HI affect the results? \\item Does changing the spectral emissivity index $\\beta$ appreciably affect the results? \\item Are there alternative kinds of models that would also explain the data? \\item How representative are the results of the clouds as a whole, given that the modeled cloud positions only represent 26\\% of the area of the Orion clouds ? \\end{itemize} Paper~III examines these questions and discusses the scientific implications of the results."
    },
    "0601/astro-ph0601062_arXiv.txt": {
        "abstract": "We discuss the preliminary results of a survey of young ($<$300 Myr), close ($<$50 pc) stars with the Simultaneous Differential Extrasolar Planet Imager (SDI) implemented at the VLT and the MMT.  SDI uses a double Wollaston prism and a quad filter to take 4 identical images simultaneously at 3 wavelengths surrounding the 1.62 $\\mu$m methane bandhead found in the spectrum of cool brown dwarfs and gas giants.  By performing a difference of images in these filters, speckle noise from the primary can be significantly attenuated, resulting in photon noise limited data.  In our survey data, we achieved H band contrasts $>$25000 (5$\\sigma$ $\\Delta$F1(1.575$\\mu$m)$>$10 mag, $\\Delta$H$>$11.5 mag for a T6 spectral type) at a separation of 0.5\" from the primary star.  With this degree of attenuation, we should be able to image (5$\\sigma$ detection) a 2-4 Jupiter mass planet at 5 AU around a 30 Myr star at 10 pc. We are currently completing our survey of young, nearby stars.  We have obtained complete datasets for 35 stars in the southern sky (VLT) and 7 stars in the northern sky (MMT). We believe that our SDI images are the highest contrast astronomical images ever made from ground or space for methane rich companions. ",
        "introduction": "Direct detection of extrasolar giant planets is extremely difficult. Giant gas planets seen in reflected light are $>$20 magnitudes fainter than their primary stars and likely lie within $\\sim$1'' of their primary stars. The problem is slightly easier with younger, hotter planets -- 100 Myr old extra-solar planets are 10$^{4-7}$ times more self-luminous than old (5 Gyr) extra-solar planets, whereas their primary stars are only slightly (2-5 times) brighter at early ages.  In theory, adaptive optics (AO) systems that are ``photon noise limited'' can detect an object up to 10$^5$ times fainter than its primary at separations $>$1''.  However, numerous surveys for extrasolar planets using large telescopes with AO systems have yielded useful limits but few confirmed giant planet candidates (Kaisler et al. 2003, Masciadri et al. 2005, Chauvin et al. 2005, Neuh\\\"auser et al. 2005).  AO surveys for young extrasolar planets only address half of the difficulty of direct detection -- the contrast limit problem.  Beyond the contrast limit problem, all AO systems suffer from a limiting ``speckle noise'' floor (Racine et al. 1999).  Within 1'' of the primary star, the field is filled with speckles left over from instrumental features and residual atmospheric turbulence after adaptive optics correction.  These speckles vary as a function of time and color. For photon noise limited data, the signal to noise S/N increases as t$^{0.5}$, where t is the exposure time.  For speckle-noise limited data, the S/N does not increase with time past a specific speckle-noise floor (limiting contrasts to $\\sim$10$^3$ at 0.5'').  This speckle-noise floor is considerably above the photon noise limit and makes planet detection very difficult.  Interestingly, space telescopes such as HST also suffer from a similar limiting speckle-noise floor due to imperfect optics and ``breathing'' (Schneider et al. 2003).  Direct detection of extrasolar giant planets requires special new instrumentation to suppress this speckle noise floor and produce photon noise limited images.  The VLT, Keck, Subaru, and Gemini are all currently developing dedicated planet-finding cameras which exploit these new instrumentational approaches for speckle suppression.  The Simultaneous Differential Imager (SDI), which our team built and installed at the VLT and MMT (see Biller et al., this conference), is one of the first dedicated planet-finding optical devices to come online on a large telescope.  Simultaneous Differential Imaging is an instrumental method which can be used to calibrate and remove the ``speckle noise'' in AO images, while also isolating the planetary light from the starlight.  This method was pioneered by Racine et al. (1999), Marois et al. (2000), Marois et al. (2002), and Marois et al. (2005). It exploits the fact that all cool (T$_{eff}$ $<$1200 K) extra-solar giant planets have strong CH$_4$ (methane) absorption redwards of 1.62 $\\mu$m in the H band infrared atmospheric window (Burrows et al. 2001, Burrows et al. 2003).  Our SDI device obtains four images of a star simultaneously through three slightly different narrowband filters (sampling both inside and outside of the CH$_4$ features).  These images are then differenced.  This subtracts out the halo and speckles from the bright star to reveal any massive extrasolar planets orbiting that star. Since a massive planetary companion will be brightest in one filter and absorbed in the rest, while the star is bright in all three, a difference can be chosen which subtracts out the star's light and reveals the light from the companion.  Thus, SDI also helps eliminate the large contrast difference between the star and substellar companions (Close et al. 2005; Lenzen et al. 2004; Lenzen et al. 2005) The SDI device has already produced a number of important scientific results: the discovery of AB Dor C (Close et al. 2005) which is the tightest (0.16'') low mass companion detected by direct imaging, detailed surface maps of Titan (Hartung et al. 2004), the discovery of $\\epsilon$ Indi Ba-Bb, the nearest binary brown dwarf (McCaughrean et al. 2005), and evidence of orbital motion for Gl 86B, the first known white dwarf companion to an exoplanet host star (Mugrauer and Neuh\\\"auser 2005). ",
        "conclusions": "\\label{sec:concl}  The novel SDI device at the VLT and MMT has been fully commissioned and is currently achieving attenuations of $>$25000 ($\\Delta$H $>$ 11.5 for a T6 spectral type object, $\\Delta$F1(1.575$\\mu$m)$>$10) at 0.5''.  With these contrasts, we can detect a wide range of substellar objects.  For instance, for AB Dor A (a 70 Myr K1V star 15 pc away) we can detect ($>$5 $\\sigma$) a 5 M$_{Jup}$ planet 12 AU from the star.  For a younger closer star (30 Myr age at 10 pc), we can detect a 2-4 M$_{Jup}$ planet at 5 AU.  We have currently observed 42 of the youngest ($<$300 Myr), nearest ($<$50 pc) stars as part of a survey of young, nearby stars. We have received time at the VLT for followup observations of 8 tentative candidates found as part of the survey. With a total sample size of $\\sim$50 stars, we will be able to place strong constraints on the frequency and semimajor axis distribution of massive extrasolar planets $>$5 AU from their primaries.  From scaling laws derived from the distribution of known radial velocity planets (Marcy et al. 2003, Lineweaver and Grether 2003, Burrows et al. 2003), we expect to detect $\\sim$4 planets for our total sample (see Nielsen et al., this conference). Whether or not we detect planets, our survey will begin to measure the true distribution young massive extrasolar planets $>$5 AU from their primaries and will provide valuable contraints for theories of planet formation and migration."
    },
    "0601/astro-ph0601254_arXiv.txt": {
        "abstract": "We show how to place constraints on cluster physics by stacking the weak lensing signals from multiple clusters found through the Sunyaev-Zeldovich (SZ) effect.  For a survey that covers about $200$ sq. deg.  both in SZ and weak lensing observations, the slope and amplitude of the mass vs. SZ luminosity relation can be measured with few percent error for clusters at $z \\sim 0.5$. This can be used to constrain cluster physics, such as the nature of feedback. For example, we can distinguish a pre-heated model from a model with a decreased accretion rate at more than 5$\\sigma$. The power to discriminate among different non-gravitational processes in the ICM becomes even stronger if we use the central Compton parameter $y_0$, which could allow one to distinguish between models with pre-heating, SN feedback and AGN feedback, for example, at more than 5$\\sigma$. Measurement of these scaling relations as a function of redshift makes it possible to directly observe e.g., the evolution of the hot gas in clusters.  With this approach the mass-L$_{\\rm SZ}$ relation can be calibrated and its uncertainties can be quantified, leading to a more robust determination of cosmological parameters from clusters surveys.  The mass-L$_{\\rm SZ}$ relation calibrated in this way from a small area of the sky can be used to determine masses of SZ clusters from very large SZ-only surveys and is nicely complementary to other techniques proposed in the literature. ",
        "introduction": "Clusters of galaxies are powerful probes of cosmology. In particular, the cluster number density as a function of mass and redshift can provide unique constraints on the growth of cosmological structure and e.g. on dark energy.  The next generation of Sunyaev-Zel'dovich (SZ) surveys, e.g., the Atacama Cosmology Telescope\\footnote{\\tt{http://www.hep.upenn.edu/act/}} (ACT) and the South Pole Telescope\\footnote{\\tt{http://spt.uchicago.edu/}} (SPT), will provide a catalog of thousands of clusters.  The flux of the SZ is proportional to the integral of the density times the temperature of the hot gas in the cluster \\citep{SZ80} and the flux limit of the catalogs will approximately translate into a mass limit. However, the actual value of the cluster's mass for a given SZ flux depends on the details of the SZ-mass relation, which in turn is governed by cluster physics.  Recent work has shown how this relation can be obtained by using numerical simulations (e.g. \\citet{Oh03, daSilva03,Motl,Nagai}) or analytical approximations (e.g. \\citet{DosSantos01,ReidSpergel, Roy05, Ostriker05}). Numerical models find that the SZ-mass relation is expected to be very tight, implying that clusters masses can be directly read out from the SZ observations once the SZ-mass relation has been calibrated. In addition, both numerical work and analytical approximations have shown that the integrated SZ luminosity of a cluster is relatively insensitive to input physics used in the models, such as pre-heating and energy injection by supernovae or Active Galactic Nuclei (e.g. \\citet{Motl,Nagai,Lapi05}). The same models, however, show a larger scatter with mass of the central SZ decrement parameter ($y_0$) and a much stronger dependence of $y_0$ on the cluster physics.  Ideally, we would like an independent and robust measurement of the cluster mass  from observations. This would open the possibility to  not just  calibrate the SZ-mass relation directly from observations, but also to test intra-cluster medium (ICM) models, and therefore to directly probe cluster physics. Arguably, weak gravitational lensing provides a direct way to measure the mass of clusters. The distortions of the background galaxies due to the cluster's  gravitational field are sensitive to the mass along the line of sight, not just the intra-cluster gas. These distortions can be used to reconstruct the cluster's projected mass density (e.g., \\citet{KaiserSquires93}).  The weak lensing signal induced by galaxy clusters has been measured for a large number of systems (e.g., \\citet{Clowe05, Sheldon01} and references therein). The amplitude of the lensing signal depends on the redshift distribution of the faint background galaxies, which can be obtained from photometric observations. Observations from the ground are limited by the atmosphere: both the maximum surface density of background galaxies and the minimum error on the shear are set by the seeing. While observations from space could circumvent this difficulty, there is no wide-field imager available in the near future to do this (ACT will cover a region of a few hundred sq. deg. and SPT 4000 sq. deg.). As both these factors set the minimum cluster mass that can be directly detected from weak lensing, a direct mass determination is possible only for fairly massive clusters, $ \\approx 10^{15} \\Msun$ (e.g., \\citet{Marian}).  Here we circumvent this limitation by taking advantage of the wide multi-wavelength coverage provided by ACT and SPT with their planned extensive optical follow up (Southern Cosmology Survey and Dark Energy Survey).  By combining (stacking) the weak lensing signals from multiple clusters with roughly the same SZ luminosity, otherwise undetectable shear signal can be amplified, and thus an average mass determination can be achieved.  The goal of this paper is to quantify the error in the calibration of the mass-SZ relation achievable by stacking weak lensing observations of SZ clusters. This also will allow one to investigate cluster physics since the scatter and amplitude in the SZ-mass relation can then be directly compared with predictions from numerical simulations. In addition, a well calibrated mass$-$SZ relation can lead to a more robust determination of cosmological parameters from cluster surveys.  Here, we will consider what can be learned by stacking clusters from a survey with the specifications of the two-year ACT SZ survey, assuming that the entire clean region (200 sq. deg.) is covered by optical weak lensing observations with a seeing of $\\sim 0.6$ to $0.7$ arcsec and $20$ source galaxies per square arcminute. The outline of the paper is as follows: in section 2 we compute the expected error in the mass estimator for the stacked clusters. In section 3 we explore how cluster physics can be constrained with the mass-SZ relation calibrated from the finding of section 2. We conclude in section 4 where we also outline the possible consequences of this approach for constraining the evolution of the hot gas in clusters and for obtaining a more robust determination of cosmological parameters from clusters surveys. ",
        "conclusions": "We have shown how, stacking the weak lensing signals from multiple clusters with roughly the same SZ luminosity, the SZ luminosity-mass ($L_{sz}-M$) and SZ central decrement-mass ($y_o-M$) relations can be measured from forthcoming SZ surveys with extensive optical follow up. For example, from a survey of 200 square degrees such as ACT, we find we should be able to determine the scaling slope to $\\pm 0.05-0.2$ and amplitude to 5-20\\% depending on redshift bin. Observations indicate that the entropy of the gas in clusters cannot be explained by gravitational collapse alone, but the cause of the increased entropy remains to be determined.  Different non-gravitational processes affect both the slope and amplitude of these relations, and in particular the $y_o-M$ relation is most sensitive to cluster physics.  Used in conjunction with results from simulations, the expected errors on the measurements of the slope and amplitude of scaling relations imply that we can start to discriminate among different non-gravitational processes affecting the ICM.  We have shown that, for a survey of $\\sim$ 200 sq. deg., the available S/N enables one to measure these scaling relations as a function of redshift: we can distinguish slopes and amplitudes of the $M-L_{sz}$ relation and $M-y_0$ relation for different redshift bins with width $\\Delta z = \\pm 0.1$, thus opening up the possibility to constrain how the mass-SZ scaling evolves with redshift. A mass measurement from weak lensing and a SZ measurement in different redshift bins,  can constrain the evolution of hot gas as a function of redshift, which in turn would enable one to constrain feedback evolution.  Inspection of Fig. \\ref{fig:3} and \\ref{fig:4} provides a quantitative measurement of how physics in clusters, e.g. feedback, can be constrained as a function of redshift. The errors reported here will scale with the survey parameters as $1/\\sqrt{f_{\\rm sky}}$ and $1/\\sqrt{n_g}$.  Thus, using this method, wider surveys can improve constraints on cluster physics just as well as deeper surveys. For a space telescope with $n_g=100$ galaxies per square arcminute and 1000 square degrees, our results would improve by a factor of 5.  A calibration of the SZ-Mass relation with known uncertainties also has consequences for the determination of cosmological parameters.  The number density of clusters within a given mass range or above a given mass is extremely sensitive to cosmology. This sensitivity, in principle, can lead to tight constraints on cosmological parameters such as $\\Omega_m$, dark energy properties and neutrino mass. However, even relatively small systematic errors in clusters mass estimates can lead to systematic errors on the recovered cosmological parameters that are larger than the statistical ones \\citep{HolderHaimanMohr, Battye03,Francis05}. This potentially compromises measurements of parameters like e.g. the dark energy equation of state. The technique outlined in this paper to calibrate the mass-SZ relation may help reduce the amplitude of these systematic errors. The weak lensing mass determination used here is sensitive to the mass along the line of sight, not just the gas. It is true that there can be biases introduced by the large-scale structure local to the clusters, but we argue that this should not be considered a fundamental limitation as this effect in principle can be quantified by means of numerical N-body simulations. If this can be achieved, our estimates above indicate that the residual uncertainty in the calibration of the amplitude of the mass-SZ relation is $~ 5$\\% for a survey with $n_g=20$ galaxies per square arcminute and sky coverage of 200 square degrees. This uncertainty would propagate into a systematic error in the determination of cosmological parameters, that is now reduced to be not larger than the statistical uncertainty achievable from a survey of the same size (see \\citet{Francis05}).  Alternatively, The M-L$_{\\rm SZ}$ relation calibrated in this way from a small area of the sky can be used to determine masses of SZ clusters from very large SZ-only surveys, or the $M-y_0$ relation could be used to constrain cluster physics and this information can then be used to model the $M-L_{sz}$ relation. This approach is nicely complementary to other techniques proposed to calibrate the mass-observable relations such as the self-calibration using the cluster power spectrum of \\citet{MajumdarMohr03}. While only positions and redshifts of clusters are needed to measure the power spectrum and thus to self-calibrate the SZ-Mass relation, it requires observations of large areas of the sky as square and contiguous as possible. The calibration of the M-SZ relation with weak lensing has more stringent follow up requirements (imaging with good seeing in addition to clusters redshifts), but has less stringent geometry requirements."
    },
    "0601/astro-ph0601124_arXiv.txt": {
        "abstract": "We present a non-parametric technique to infer the projected-mass distribution of a gravitational lens system with multiple strong-lensed images. The technique involves a dynamic grid in the lens plane on which the mass distribution of the lens is approximated by a sum of basis functions, one per grid cell. We used the projected mass densities of Plummer spheres as basis functions. A genetic algorithm then determines the mass distribution of the lens by forcing images of a single source, projected back onto the source plane, to coincide as well as possible. Averaging several tens of solutions removes the random fluctuations that are introduced by the reproduction process of genomes in the genetic algorithm and highlights those features common to all solutions. Given the positions of the images and the redshifts of the sources and the lens, we show that the mass of a gravitational lens can be retrieved with an accuracy of a few percent and that, if the sources sufficiently cover the caustics, the mass distribution of the gravitational lens can also be reliably retrieved. A major advantage of the algorithm is that it makes full use of the information contained in the radial images, unlike methods that minimise the residuals of the lens equation, and is thus able to accurately reconstruct also the inner parts of the lens. ",
        "introduction": "The deflection of light caused by a gravitational lens and the amplifying and distorting effect thereof on the images of background sources, provide us with a means to measure the total mass of the lensing object. If only, of course, that it is possible to ``invert'' such a lensing system, i.e. to infer the mass distribution of the lens given the positions and shapes of a set of lensed images of background sources and the redshifts of the lens and the sources. The inversion of gravitational lensing systems is interesting in its own right, since it puts constraints on the spatial dark matter distribution of the lensing objects and thus helps constrain dark matter physics \\citep{n04,d04}, but it may also contribute to cosmology. Good reconstructions of lensing systems may help to constrain the density parameter and the redshift evolution of the dark energy \\citep{y01,s04,m05}.  The use of the gravitational lensing effect to measure masses of point-mass lenses (or ``stars'') was envisaged by \\citet{r64} and \\citet{l64}. The idea of using simple parametric models for extended lenses, such as galaxies or clusters of galaxies, can be traced back to \\citet{dr80}, who use a King model to estimate the mass distribution of the galaxy that produces the two images of QSO~0957+561A,~B. Since then a host of parametric inversion methods has been developed and applied to observed strong lensing systems, such as the ring cycle method of \\citet{k89} or the maximum entropy method of \\citet{w94}. \\citet{kn93} fitted a bimodal lensing potential, constructed from two elliptical pseudo-isothermal potentials, images of the cluster A370, obtained from the ground under excellent seeing conditions. More elaborate parametric models for lensing clusters associate a simple mass distribution, e.g. a power-law of radius, to each galaxy and to the cluster as a whole. The many parameters that define the lens model are then determined by a $\\chi^2$ fit to the observed images \\citep{tkd98}. \\citet{b05} suggest using a gravitational lens as a cosmic magnifying glass to recover structural details about distant sources. These authors assume a parametric form for the lens mass distribution and use a genetic algorithm to non-parametrically reconstruct the surface brightness distribution of the source. If only a few sources are being lensed or if the sources do not sufficiently cover the caustics, parametric methods are clearly the preferred approach.  However, the multitude of arcs and distorted images visible in massive clusters (e.g. \\citealt{s87,lh88}), which are observed routinely now at high spatial resolution with the Hubble Space Telescope (e.g. \\citealt{br05}), contain a wealth of information and call for more flexible and model independent inversion methods. Using the pixelation method \\citep{sw97,a98}, the lens plane is divided into a static grid. The mass in each grid cell and the source positions are estimated so as to construct a solution that best reproduces the observed image positions, subject to regularizing constraints that ensure a smooth mass distribution for the lens that stays close to the luminosity distribution. \\citet{t00} perform a multipole-Taylor expansion of the two-dimensional lensing potential, the coefficients of which are determined by a $\\chi^2$-fit to the observed images. In \\citet{d05}, a non-parametric inversion technique, called {\\sc slap}, is presented and applied to an HST image of the cluster Abell 1689 \\citep{d05b,k02} that, like the pixelation method, makes use of a grid division of the mass distribution of the lens. This time, however, the grid is dynamic:~it is refined iteratively where the mass density is large. Recently, \\citet{d05c} extended {\\sc slap} to {\\sc wslap} in order to also take into account information in the weak lensing regime. Thus, {\\sc slap} and {\\sc wslap} are fast and versatile tools for inverting observed lens systems. However, both methods assume the background galaxies to be point sources, which may lead to an over-estimation of the central mass density of the lens in order to focus the images into very compact sources and to physically implausible regions with negative mass density in the lens plane. Adding weak-lensing information alleviates the dependence of the solution on this minimization threshold. {\\sc wslap} can make use of quadratic programming to avoid unphysical negative mass densities for the lens. However, this limits the analysis to observables that are linear functions of the lens mass density, such as image positions, and does not allow to incorporate e.g. surface brightness information, which depends non-linearly on the lens mass density. \\citet{bra05} also used weak and strong lensing data to invert the X-ray cluster RX~J1347.5$-$1145. Their method evaluates the gravitational potential on a non-dynamic grid. The best fitting gravitational potential is then constructed non-parametrically by minimising a $\\chi^2$ function, starting from a  parametric priorsolution.  The ideal non-parametric lens inversion algorithm {\\em (i)} should be free of any assumptions regarding the mass distribution of the lens or the luminosity distributions of the sources, {\\em (ii)} should not depend on any prior on the lens mass distribution or any regularisation scheme that could bias the solution, {\\em (iii)} should not produce unphysical, i.e. negative, mass densities {\\em (iv)} should be free of any uncontrollable parameters, {\\em (v)} should be easily extendible to any kind of data, both in the strong and weak lensing regimes, without having to change the inner workings of the algorithm or having to worry about features like continuity or differentiability of the objective function that is extremised. These conceptual issues are the main motivation for this paper, rather than computational speed.  Genetic algorithms do an excellent job at fulfilling all these constraints. In this paper, we describe and test a new non-parametric lens inversion technique. The technique makes use of a dynamic grid on which the mass distribution of the lens is approximated by a weighted sum of basis functions and of a genetic algorithm to determine the unknown weights. In the strong lensing regime, where multiple images of each source are available, the following data are offered to the algorithm:~the redshifts of the sources and the lens, and the observed positions of the images. The genetic algorithm generates solutions that satisfy only one minimal constraint:~the back-projected images of a single source should overlap as well as possible in the source plane.  We briefly describe the relevant background to gravitational lens systems and genetic algorithms in Section \\ref{sec:back}. The details of the inversion method are given in Sect. \\ref{sec:invert}. We discuss a number of tests to which we subjected the method in Sect. \\ref{sec:sim}. Finally, our conclusions are summarized in Sect. \\ref{sec:conc}. ",
        "conclusions": "\\label{sec:conc}  The procedure described and illustrated above is a non-parametric method for inverting gravitational lenses, making no a priori assumptions regarding the shape of the lens. We only impose the condition that an acceptable solution must be able to map images of the same source onto overlapping regions in the source plane. The procedure only requires that one can identify which images correspond to the same source. In particular, no information about the sizes of the sources needs to be provided. The size of the grid on which the algorithm will determine the mass distribution of the lens needs to be specified. However, because a single solution can be obtained relatively fast, it is an easy task to try a variety of sizes until one is obtained which generates a lens with a good fitness value. A multi-resolution grid with a few hundred cells is usually sufficient to represent any plausible lens mass distribution. The implementation used for the simulation in this paper employs a population of $250$ genomes. Based on a suite of simulations, this value proved to be sufficiently large to to yield good solutions within an acceptable amount of time. The calculations were done in a distibuted manner, using sixteen Intel \\textregistered{} Xeon \\texttrademark{} 2.4 GHz processors of a computer cluster. Depending on the number of sources, creating a single solution may require several hours. To give a specific example, the $25$ solutions used in the simulation were created in four days.  The simulation discussed in this article, together with the many others we performed, indicate that our inversion technique successfully solves the lens inversion problem. The reconstructed sources lie close to their true positions and their shapes are retrieved quite accurately as well. The procedure determines the mass of the lens very accurately. The averaging procedure guarantees the removal of random fluctuations and yields a smooth mass density very close to the true solution. Note that because the masses of the individual Plummer distributions are always represented by positive numbers in the genome, no negative mass densities will be produced.  Of course, the quality of the reconstruction depends on the quality of the information at hand. The algorithm depends on the availability of multiply imaged sources, which identifies the relevant area in our procedure as the area within the outermost caustic. When this area is sampled well by the sources, we can expect a good reconstruction of the mass density, as indicated by the simulation described previously. A major advantage of the algorithm is that it makes full use of the information contained in the radial images, unlike methods that minimise the residuals of the lens equation, and is thus able to accurately reconstruct also the inner parts of the lens.  Another advantage of a genetic algorithm is the ease with which the fitness criterion can be specified. One simply has to devise a way to associate a fitness value with a specific genome, without worrying about features like continuity or differentiability. E.g., to test whether information about the amplifying effect of the gravitational lens can improve the lens reconstruction, we added a contribution to the fitness value:~the maximum brightness values in the back-projected images of a single source should lie as close as possible to each other. Several simulations indicated that augmenting the procedure in this fashion does not improve the final result since the $\\Vec{\\beta}(\\Vec{\\theta})$ mapping was already very well approximated anyway. Not requiring surface brightness information not only reduces the computational cost, but also makes the method less sensitive to noise in the images. Incorporating information about the shear field at larger radii outside the gravitational lens is straightforward. For each genome the shear components $\\gamma_1(\\Vec{\\theta})$ and $\\gamma_2(\\Vec{\\theta})$ can be calculated and compared with the observed values. The expression for the fitness can be changed so as to penalise genomes with large differences between the measured and the model shear. Further improvements of the genetic algorithm, such as making the mutation amplitude depend automatically on the convergence of the fitness, are easily implemented and are left as future research. An application to real data will be presented in a subsequent paper."
    },
    "0601/astro-ph0601638_arXiv.txt": {
        "abstract": "We use the optical and near-infrared galaxy samples from the Munich Near-Infrared Cluster Survey (MUNICS), the FORS Deep Field (FDF) and GOODS-S to probe the stellar mass assembly history of field galaxies out to $z \\sim 5$. Combining information on the galaxies' stellar mass with their star-formation rate and the age of the stellar population, we can draw important conclusions on the assembly of the most massive galaxies in the universe: These objects contain the oldest stellar populations at all redshifts probed. Furthermore, we show that with increasing redshift the contribution of star-formation to the mass assembly for massive galaxies increases dramatically, reaching the era of their formation at $z \\sim 2$ and beyond. These findings can be interpreted as evidence for an early epoch of star formation in the most massive galaxies in the universe. ",
        "introduction": "In recent years, there has been considerable interest in the relation of the stellar mass in galaxies and their star-formation rate (SFR), since this allows to quantify the contribution of the recent star formation to the build up of stellar mass for different galaxy masses. \\citet{feulner:Cowie1996} were the first to investigate this connection for a $K$-selected sample of $\\sim 400$ galaxies at $z<1.5$ and noted an emerging population of massive, heavily star forming galaxies at higher redshifts, a phenomenon they termed ``down-sizing''.  Later on, the specific star-formation rate (SSFR), defined as the SFR per unit stellar mass, was used to study this relation. ",
        "conclusions": ""
    },
    "0601/astro-ph0601312_arXiv.txt": {
        "abstract": "Since its launch in October 2002, ESA's INTEGRAL observatory has enabled significant advances to be made in the study of Galactic nucleosynthesis. In particular, the imaging Ge spectrometer SPI combines for the first time the diagnostic powers of high resolution gamma-ray line spectroscopy and moderate spatial resolution. This review summarizes the major nucleosynthesis results obtained with INTEGRAL so far. Positron annihilation in our Galaxy is being studied in unprecented detail. SPI observations yield the first sky maps in both the 511~keV annihilation line and the positronium continuum emission, and the most accurate spectrum at 511~keV to date, thereby imposing new constraints on the source(s) of Galactic positrons which still remain(s) unidentified. For the first time, the imprint of Galactic rotation on the centroid and shape of the 1809~keV gamma-ray line due to the decay of $^{26}$Al has been seen, confirming the Galactic origin of this emission. SPI also provided the most accurate determination of the gamma-ray line flux due to the decay of $^{60}$Fe. The combined results for $^{26}$Al and $^{60}$Fe have important implications for nucleosynthesis in massive stars, in particular Wolf-Rayet stars. Both IBIS and SPI are searching the Galactic plane for young supernova remnants emitting the gamma-ray lines associated with radioactive $^{44}$Ti. None have been found so far, which raises important questions concerning the production of $^{44}$Ti in supernovae, the Galactic supernova rate, and the Galaxy's chemical evolution. ",
        "introduction": "Gamma-ray line astronomy has opened a new and unique window for studying nucleosynthesis in our Galaxy. The singular advantage of gamma-ray spectroscopy over other observations is that it offers the opportunity to detect directly and identify uniquely individual isotopes at their birthplaces. The rate at which radioactive decays proceed is in general unaffected by the physical conditions in their environment, such as temperature or density. The interstellar medium is not dense enough to attenuate gamma rays, so that radioactive decays can be observed throughout our Galaxy. Recent reviews on implications of gamma-ray observations for nucleosynthesis in our Galaxy can be found in \\cite{D-P-vB05} and \\cite{Prantzos05}. ",
        "conclusions": ""
    },
    "0601/astro-ph0601682.txt": {
        "abstract": "{ We study the effects of a metallicity variation on the thermal balance and [CII] fine-structure line strengths in interstellar photon dominated regions (PDRs). We find that a reduction in the dust-to-gas ratio and the abundance of heavy elements in the gas phase changes the heat balance of the gas in PDRs. The surface temperature of PDRs decreases as the metallicity decreases except for high density ($n>10^6$ cm$^{-3}$) clouds exposed to weak ($\\chi< 100$) FUV fields where vibrational H$_2$-deexcitation heating dominates over photoelectric heating of the gas. We incorporate the metallicity dependence in our KOSMA-$\\tau$ PDR model to study the metallicity dependence of [CII]/CO line ratios in low metallicity galaxies.  We find that the main trend in the variation of the observed CII/CO ratio with metallicity is well reproduced by a single spherical clump, and does not necessarily require an ensemble of clumps as in the semi-analytical model presented by Bolatto et al. (1999). ",
        "introduction": "The bright line emission from photon-dominated regions (PDRs) is one of the key tracers of the star formation activity throughout the evolution of galaxies in the course of the cosmological evolution. Hence, proper modeling of PDR emission is of central importance for the interpretation of the observations, in order to derive the physical parameters and the chemical state of the ISM in external galaxies. The extensive literature on PDR emission, both observationally and from the modeling side, has largely concentrated on bright Galactic sources and starburst galaxies. The effect of different metallicity for the resulting PDR emission has, up to now, drawn little attention. It is, however, very important in order to cover the full course of galactic evolution, starting with low metallicity material of cosmological origin. Many nearby galaxies, such as dwarf galaxies, irregular galaxies and the Magellanic Clouds have a low metallicity \\citep{lisenfeld98, kunth00, pustilnik02, lee03}. Within the Galaxy, as well as in other spiral galaxies, there is a radial decrease in the metallicity of molecular clouds and associated H{\\small II} regions \\citep{zaritsky94, arimoto96, giveon, bresolin04}. These systems provide the opportunity to study star formation and photon-dominated regions (PDRs) for a variety of metallicities.  In PDRs the molecular gas is heated by the far-ultraviolet (FUV) radiation field, either the strong FUV radiation in the vicinity of hot young stars, or weak average FUV fields in the Galaxy. The gas cools through the spectral line radiation of atomic and molecular species (Hollenbach \\& Tielens 1999, Sternberg 2004). The gas-phase chemistry together with a depth dependent FUV intensity lead to the formation of atomic and molecular species at different depths through the cloud. This typical stratification of PDRs is for example reflected by the the H/H$_2$ and C$^+$/C/CO transitions \\citep{SD95, boger05}. At low visual extinctions the gas is cooled by emission of atomic fine-structure lines, mainly [CII] 158$\\mu$m and [OI] 63$\\mu$m. At larger depths, millimeter, sub-millimeter and far-infrared molecular rotational-line cooling (CO, OH, H$_2$O) becomes important together with the interaction of dust and gas. Physical conditions such as temperature and density can be derived, by  comparing the observed line emissions with model predictions \\citep{lebourlot93, wolfire95, stoerzer96, warin96, kaufman99, zielinsky00, stoerzer00, gorti02}.  [CII] emission is a widely used diagnostic indicator of star formation \\citep{stacey91, pierini99, malhotra00, boselli02b, pierini03, kramer04}. Observations suggest that low metallicity systems have higher [CII] to CO rotational line ratios compared to the Galactic value. In particular, the intensity ratio $I_{\\mathrm{[CII]}}/I_{\\mathrm{CO}}$ may vary from $\\sim 1000$ in the inner Milky Way, up to $\\sim 10^5$ in extremely low metallicity systems \\citep[eg.][]{madden97, mochi98, bolatto99, madden00, hunter01}. Several studies have suggested that a lower abundance of heavy elements affects the chemical structure of PDRs and the cooling line emission, and that estimates of molecular gas masses from the observed CO(J=1-0) line intensities using the standard conversion factor may underestimate the true masses in such objects \\citep{wilson95, israel97, israel03, rubio04}.   \\citet{bolatto99} modelled the metallicity variation of the line ratio [CII]/CO(1-0), for an ensemble of spherical ``clumps'', assuming an inverse relation between the size of the C$^+$ region and the metallicity. However the sizes of the C$^+$, C and CO regions also depend on the chemistry in PDRs and the chemical network is modified at low metallicities \\citep{lequex94}. Additionally it has been suggested that the size of the C$^+$, C and CO regions may also  significantly depend on the overall cloud morphology, e.g. density variations \\citep{hegmann03} and velocity fluctuations \\citep{roellig02}. Moreover the temperature of the molecular gas might depend on the metallicity which affects the observable line intensities \\citep{wolfire95}.   We study the effects of metallicity changes on the temperature and chemical structure of PDRs. In \\S~\\ref{metalsection} we consider the dependence of the PDR gas temperature on the metallicity using a simplified semi-analytic model and compare it with numerical results from full PDR model calculations. Our computations were carried out using an updated version of our spherical KOSMA-$\\tau$ model \\citep{stoerzer96} which was originally adapted from the plane-parallel model presented by Sternberg \\& Dalgarno (1995).  In \\S~\\ref{ciiw} we examine the predicted size of the C$^+$ zones as a function of metallicity. We then model the strength of the [CII] emission and investigate the the dependence of the [CII]/CO(J=1-0) line ratio on the metallicity. Finally we compare the results with observational data in \\S~\\ref{clumpysection}. ",
        "conclusions": "We study the effects of  metallicity variations on the gas temperature and [CII] emission line properties of spherical PDRs. We find that the surface temperature of PDRs at high UV fields varies linearly with metallicity. For low UV fields and high densities this metallicity behavior of the surface temperature is converse, showing an inverse dependence with metallicity due to the dominant H$_2$ heating. We introduce a new two level FUV H$_2$ heating and cooling function that properly accounts for energy losses via vibrational collisional excitations.  We examine the dependence of the C$^+$ envelope on metallicity and find that its geometrical depth scales inversely with $Z$. This produces a higher [CII]/CO(J=1-0) line ratio at lower metallicities.  We used the numerical results from the KOSMA-$\\tau$ model to study the dependence of PDR emission lines with metallicity.  The observed variation of [CII\u00b4]/CO(J=1-0) with metallicity can be explained well by a single-clump model and it is not necessary to refere to an average over a clump ensemble. We conclude that the [CII]/CO(J=1-0) line ratios for sources with differing metallicities do not provide a strong constraint on the clumpy morphology of a molecular clouds."
    },
    "0601/astro-ph0601306_arXiv.txt": {
        "abstract": "An efficient algorithm for adaptive kernel smoothing (AKS) of two-dimensional imaging data has been developed and implemented using the Interactive Data Language (IDL). The functional form of the kernel can be varied (top-hat, Gaussian etc.) to allow different weighting of the event counts registered within the smoothing region. For each individual pixel the algorithm increases the smoothing scale until the signal-to-noise ratio (s.n.r.) within the kernel reaches a preset value.  Thus, noise is suppressed very efficiently, while at the same time real structure, i.e.\\ signal that is locally significant at the selected s.n.r.\\ level, is preserved on all scales. In particular, extended features in noise-dominated regions are visually enhanced. The {\\sc asmooth} algorithm differs from other AKS routines in that it allows a quantitative assessment of the goodness of the local signal estimation by producing adaptively smoothed images in which all pixel values share the same signal-to-noise ratio {\\em above the background}.  We apply {\\sc asmooth} to both real observational data (an X-ray image of clusters of galaxies obtained with the Chandra X-ray Observatory) and to a simulated data set. We find the {\\sc asmooth}ed images to be fair representations of the input data in the sense that the residuals are consistent with pure noise, i.e.\\ they possess Poissonian variance and a near-Gaussian distribution around a mean of zero, and are spatially uncorrelated. ",
        "introduction": "Smoothing of two-dimensional event distributions is a procedure routinely used in many fields of data analysis. In practice, smoothing means the convolution \\begin{eqnarray} I'(\\vec{r}\\,) \\equiv I(\\vec{r}\\,) \\otimes {\\cal K}(\\vec{r}\\,) = \\int_{{\\sf I\\hspace*{-0.25ex}R}^2} {\\cal K}(\\vec{r}-\\vec{r}\\,')\\, I(\\vec{r}\\,') \\, d\\vec{r}\\,'  \\label{convol} \\\\ \\left(  \\int_{{\\sf I\\hspace*{-0.25ex}R}^2} {\\cal K}(\\vec{r}\\,)\\,d\\vec{r}\\,' = 1 \\right) \\hspace*{2cm} \\nonumber \\end{eqnarray} of the measured data $I(\\vec{r}\\,)$ with a kernel function $\\cal K$ (often also called `filter' or `window function'). Although the raw data may be an image in the term's common meaning [i.e. the data set can be represented as a function $I(x,y)$ where $I$ is some intensity, and $x$ and $y$ are spatial coordinates], the two coordinates $x$ and $y$ can, in principle, describe any two-dimensional parameter space.  The coordinates $x$ and $y$ are assumed to take only discrete values, i.e.\\ the events are binned into $(x,y)$ intervals. The only requirement on $I$ that we shall assume in all of the following is that $I$ is the result of a counting process in some detector, such that $I(x,y) \\in {\\sf I\\hspace*{-0.25ex}N}_0$.  An image, as defined above, is a two-dimensional histogram and is thus often a coarse representation of the underlying probability density distribution (e.g.\\ Merrit \\& Tremblay 1994, Vio et al.\\ 1994). However, for certain experiments, an unbinned event distribution may not even exist -- for instance if the $x$ and $y$ values correspond to discrete PHA (spectral energy) channels. Also, {\\em some}\\/ binning can be desirable, for instance, in cases where the dynamic range of the data under consideration is large. If the bin size is sufficiently small, the unavoidable loss of spatial resolution introduced by binning the raw event distribution may be a small price to be paid for a data array of manageable size.  Smoothing of high-resolution data is of interest whenever the signal (defined as the number of counts per pixel above the expected background) in the region of interest in $x-y$ space is low, i.e.\\ is less than or of the order of 10, after the raw event distribution has been sorted into intervals whose size matches approximately or exceeds the instrumental resolution. It is crucial in this context that the observed count statistics are not taken at face value but are corrected for background, which may be internal, i.e.\\ originating from the detector (more general: the instrumental setup), or external. If this correction is not applied, the observed intensity (counts) distribution $I(x,y)$ may be high across the region of interest, suggesting good count statistics, even if the signal above the background that we are interested in is actually low and poorly sampled. The statistics of the observed counts alone can thus be a poor indicator of the need for image smoothing.  Rebinning the data set into larger, and thus fewer, intervals improves the count statistics per pixel and reduces the need for smoothing.  This is also the basic idea behind smoothing with a kernel of the form \\begin{equation} {\\cal K}(\\vec{r}\\,',\\sigma) = \\left\\{ \\begin{array}{ll}  1/(\\pi \\sigma^2) & \\mbox{where \\,$|\\vec{r}\\,'| < \\sigma $} \\\\ 0   & \\mbox{elsewhere} \\end{array} \\right. \\label{tophat} \\end{equation} (circular `top-hat' or `box-car' filter of radial size $\\sigma$), the only difference being that smoothing occurs semi-continuously (the step size being given by the bin size of the original data) whereas rebinning requires an additional phase information [the offset of the boundaries of the first bin with respect to some point of reference such as the origin of the $(x,y)$ coordinate system]. However, when starting from an image binned at about the instrumental resolution, both rebinning and conventional smoothing share a well known drawback, namely that any improvement in the count statistics occurs at the expense of spatial resolution. ",
        "conclusions": "We describe {\\sc asmooth}, an efficient algorithm for adaptive kernel smoothing of two-dimensional image data. {\\sc Asmooth} determines the local size of the kernel from the requirement that, at each position within the image, the signal-to-noise ratio of the counts under the kernel, and above the background, must reach (but not exceed greatly) a certain preset minimum. Qualitatively, this could be called the `uniform significance approach'. As a consequence of this criterion, noise is heavily suppressed and real structure is enhanced without loss of spatial resolution.  Due to the choice of boundary conditions, the algorithm preserves the total number of counts in the raw data.  The {\\sc asmooth}ed image is a fair representation of the input data in the sense that the residual image is consistent with pure noise, i.e., it the residual possesses Poissonian variance and is spatially uncorrelated. As demonstrated by the results obtained for simulated input, the adaptively smoothed images created by {\\sc Asmooth} are fair representations of the true counts distribution underlying the noisy input data.  Since the background is accounted for in the assessment of any structure's s.n.r., a feature that distinguishes {\\sc asmooth} from most other adaptive smoothing or adaptive binning algorithms, the interpretation of the adaptively smoothed image is straightforward: all features in the {\\sc asmooth}ed image are equally significant at the scale of the kernel used in the smoothing. A map of these smoothing scales is returned by {\\sc asmooth} together with the smoothed image. Note that these are the {\\em smallest}\\/ scales at which real features reach the selected s.n.r.\\ threshold. Consequently, background regions never meet the s.n.r.\\ criterion and are smoothed at the largest possible scale for which the kernel size equals the size of the image. Such regions of insufficient signal are easily flagged using an s.n.r.\\ map which is also provided by our algorithm.  We emphasize that our definition of ``signal'' as meaning ``signal {\\em above the (local) background}'' was chosen because many, if not most, real-life appplications do not provide the user with a priori knowledge of the intensity of the background or any spatial background variations. Also, superpositions of features (such as the point sources on top of diffuse emission in our Chandra ACIS-S example) require a {\\em local}\\/ background estimate if the intrinsic signal of features on very different scales is to be assessed reliably. The approach taken by {\\sc Asmooth} is thus intrinsically very different from the one implemented in, for instance, the adaptive binning algorithm WVT (Diehl \\& Statler 2005) and the difference in the resulting output is accordingly dramatic (see Section~\\ref{simsec}).  {\\sc Asmooth} is being used extensively in the analysis of astronomical X-ray imaging data gathered in a wide range of missions from ROSAT to XMM-Newton, and has become the analysis tool of choice for the display of high-resolution images obtained with the Chandra X-ray observatory. Recent examples illustrating {\\sc asmooth}'s performance can be found in Ebeling, Mendes de Oliveira \\& White (1995), Brandt, Halpern \\& Iwasawa (1996), Hamana et al.\\ (1997), Ebeling et al.\\ (2000), Fabbiano, Zezas \\& Murray (2001), Krishnamurthi et al.\\ (2001), Karovska et al.\\ (2002), Bauer et al.\\ (2002), Gil-Merino \\& Schindler (2003), Heike et al.\\ (2003), Rasmussen, Stevens \\& Ponman (2004), Clarke, Blanton \\& Sarazin (2004), Ebeling, Barrett \\& Donovan (2004), or Pratt \\& Arnaud (2005), as well as in many Chandra press releases (\\mbox{\\tt{http://chandra.harvard.edu/press/}}).  Under development is an improved version of the algorithm which accounts for Poisson statistics using the analytic approximations of Ebeling (2003) to allow proper treatment of significant negative features, such as absorbed regions, detector chip gaps, or instrument elements (e.g., the window support structure of the ROSAT PSPC) obscuring part of the image.  {\\sc Asmooth} is written in the IDL programming language. The source code is available on request from ebeling{@}ifa.hawaii.edu.  A $C++$ version of an early version of the code, called {\\sc csmooth}, is part of CIAO, the official suite of data analysis tools for the Chandra X-ray observatory, and is not identical to the algorithm described here."
    },
    "0601/astro-ph0601130_arXiv.txt": {
        "abstract": "In order to trace the origin and evolution of carbon in the Galactic disk we have determined carbon abundances in 51 nearby F and G dwarf stars. The sample is divided into two kinematically distinct subsamples with 35 and 16 stars that are representative of the Galactic thin and thick disks, respectively.  The analysis is based on spectral synthesis of the forbidden [C\\,{\\sc i}] line at 872.7\\,nm using spectra of very high resolution ($R\\approx 220\\,000$) and high signal-to-noise ($S/N\\gtrsim 300$) that were obtained with the CES spectrograph on the ESO 3.6-m telescope on La Silla in Chile. We find that [C/Fe] versus [Fe/H] trends for the thin and thick disks are totally merged and  flat for sub-solar metallicities. The thin disk that extends to higher metallicities than the thick disk, shows a shallow decline in [C/Fe] from $\\rm [Fe/H]\\approx0$ and up to $\\rm [Fe/H]\\approx+0.4$. The [C/O] versus [O/H] trends are well separated between the two disks (due to differences in the oxygen abundances) and bear a great resemblance to the [Fe/O] versus [O/H] trends. Our interpretation of our abundance trends is that the sources that are responsible for the carbon enrichment in the Galactic thin and thick disks have operated on a time-scale very similar to those that are responsible for the Fe and Y enrichment (i.e., SN\\,Ia and AGB stars, respectively).  We further note that there exist other observational data in the literature that favour massive stars as the main sources for carbon. In order to match our carbon trends, we believe that the carbon yields from massive stars then must be very dependent on metallicity for the C, Fe, and Y trends to be so finely tuned in the two disk populations. Such metallicity dependent yields are no longer supported by the new stellar models in the recent literature. For the Galaxy we hence conclude that the carbon enrichment at  metallicities typical of the disk is mainly due to low and intermediate mass stars, while massive stars are still the main carbon contributor at low metallicities (halo and metal-poor thick disk). ",
        "introduction": "Next to hydrogen, helium and oxygen, carbon is the most abundant element in the Universe and plays an important r\\^ole in the chemical evolution of galaxies. Although the nuclear process (helium burning) that is responsible for the formation of carbon atoms is well-known \\citep[see, e.g., reviews by][]{wallerstein1997, arnett2004}, it is unclear which objects that contribute to the carbon enrichment of the interstellar medium. This has been much debated during recent years; \\cite{gustafsson1999} conclude that carbon is mainly contributed from massive, metal-rich stars, and not from low-mass stars; \\citet{chiappini2003, chiappini2003b} find strong indications that carbon is produced in low and intermediate mass stars; \\cite{shi2002} find that carbon is contributed by super-winds from massive stars in the early stages of disk formation in the Galaxy, while at later stages significant amounts of carbon are contributed by low-mass stars. Other studies that favour a mixture of low mass, intermediate mass, and high mass stars include \\cite{liang2001, gavilan2005, carigi2005} while \\cite{henry2000} and \\cite{akerman2004} suggest massive stars to be the main source.  Only few previous studies \\citep{andersson1994, gustafsson1999, takeda2005} have based their carbon abundances on the forbidden [C\\,{\\sc i}] line at 872.7\\,nm. This line is highly insensitive to errors in the stellar model atmosphere \\citep[see, e.g.,][]{asplund2005} and gives carbon abundances of high accuracy in solar-type stars, even when using one-dimensional (1-D) models under the assumption of local thermodynamic equilibrium (LTE).  Unfortunately, the line is very weak and is blended with a weak feature that is likely to be an Fe\\,{\\sc i} line \\citep[e.g.][]{lambert1967}.  So, in order to achieve accurate abundances from this line it is desirable to have spectra of high-resolution and high signal-to-noise ($S/N$) and it is essentially only the \\cite{gustafsson1999} study that has reliable carbon trends based on the forbidden line (see discussion in Sect.~\\ref{sec:systematicerrors}). They did, however, not include the blending Fe\\,{\\sc i} line in their analysis.  As a step to achieve a deeper understanding of the sites of carbon formation and the chemical histories of the Galactic disks we have obtained high-resolution, high S/N spectra of the forbidden [C\\,{\\sc i}] line at 872.7 nm for 51 dwarf stars in the solar neighbourhood. These stars are already well studied, have well determined oxygen and iron abundances, and have known kinematics \\citep{bensby2003, bensby2004, bensby2005}.  That the kinematics are known is important as this makes it possible to properly investigate the carbon trends in the two disks, independent of each other, and to perform a differential analysis of the carbon trends in the thin and thick disks in order to investigate the differences and similarities in the disks' chemical histories. The latter investigation could, combined with other evidence, be important for understanding the time-scales involved in the chemical enrichment of the gas that today make up what we observe as the stars in the Galactic thick disk.  In Sect.~\\ref{sec:sample} we shortly describe the stellar sample and the observations. Sect.~\\ref{sec:reduction} discusses the data reduction and associated problems and Sect.~\\ref{sec:abundances} discusses the abundance analysis and estimate random and systematic errors in the derived carbon abundances. In Sect.~\\ref{sec:discussion} we describe our resulting carbon trends and we use our results to trace the possible sources of carbon as well as discuss the chemical evolution of the Galactic disks. Finally, Sect.~\\ref{sec:summary} concludes with a summary of our findings.  \\begin{figure} \\centering \\resizebox{0.9\\hsize}{!}{\\includegraphics[bb=100 0 370 80, clip]{FIG_1a.eps}} \\resizebox{0.9\\hsize}{!}{\\includegraphics[bb=100 0 370 80, clip]{FIG_1b.eps}} \\resizebox{0.9\\hsize}{!}{\\includegraphics[bb=100 0 370 80, clip]{FIG_1c.eps}} \\caption{ {\\bf Top:} A normalised flatfield frame. {\\bf Middle:} A ``raw\" science frame (in this case a solar spectrum) {\\bf Lower:} The flatfielded science frame. The strong Si\\,{\\sc i} line at 872.8\\,nm, in whose left wing the [C\\,{\\sc i}] line is located, has been marked. The full CCD frames are 4096x1024 pixels. Here we show the central parts (approximately 2000x500 pixels in size) from $\\sim 871.6$\\,nm to $\\sim 874.0$\\,nm. } \\label{fig:ces} \\end{figure} \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{FIG_2.eps}} \\caption{The average profile (based on 4096 rows) perpendicular to the dispersion direction. Note that the background level is higher on the right hand side (pixel value $\\sim 240$) than on the left hand side (pixel value $\\sim 220$). } \\label{fig:profile} \\end{figure} ",
        "conclusions": "Before we discuss our carbon results and where carbon might be produced we first, briefly, recap some of the most important features of our sample selection, i.e., the kinematic definition of the two samples. We then turn to a description of the abundance trends found and their implications for where carbon can have formed.  \\subsection{The kinematic definition of our samples} \\label{sec:kinematic}  \\begin{table} \\centering \\caption{Kinematic properties of the stellar populations that were used when calculating probabilities to select thin and thick disk stars. The properties of the Hercules stream have been adopted from Famaey et al.~(2005) and the properties of the other populations are the same as in Bensby et al.~(2005) except for the thin disk that has a lower number density ($X$) due to the now included Hercules stream. All velocities and velocity dispersions are given in km\\,s$^{-1}$.} \\label{tab:kinematics} \\centering% \\setlength{\\tabcolsep}{2mm} \\begin{tabular}{lrrrrrrr} \\hline & $X$    & $\\sigma_{\\rm U}$ & $\\sigma_{\\rm V}$ & $\\sigma_{\\rm W}$ & $U_{\\rm lag}$ & $V_{\\rm lag}$ & $W_{\\rm lag}$ \\\\ \\hline Thin disk      & 0.84   &  35 & 20 & 16 &    0  &  $-15$ &   0   \\\\ Thick disk     & 0.10   &  67 & 38 & 35 &    0  &  $-46$ &   0   \\\\ Hercules       & 0.06   &  26 &  9 & 17 & $-40$ &  $-50$ & $-7$  \\\\ Halo           & 0.015  & 160 & 90 & 90 &    0  & $-220$ &   0   \\\\ \\hline \\end{tabular} \\end{table}  The stars in our sample were selected, based on their kinematics, to either be typical representatives of the thin or the thick Galactic disk.  The kinematic selection is further described  in \\cite{bensby2003,bensby2005} including an extended discussion of how e.g. the local normalisation of the thick disk influences the categorisation of an individual star.  In Fig.~\\ref{fig:kinematics}a we show a so called Toomre diagram for our stars. The distribution of the stars in our two samples in this plot is essentially  the same as \\cite{fuhrmann2004} finds for his thin and thick disk samples.  Figure~\\ref{fig:kinematics}b shows the $W_{\\rm LSR}$ velocity (which is proportional to the maximum distance below/above the Galactic plane, $Z_{\\rm max}$, a star can reach) as a function of [Fe/H]. We see that even at $\\rm [Fe/H]=0$ or higher there are  stars with high $W_{\\rm LSR}$. The number of such stars appear to  be higher than one would expect from orbital heating by e.g. molecular  clouds \\citep{hanninen2002}. One should though remember that our sample is far from complete and we certainly are subject to various biases. However, our point  is that these stars do exist, are readily found in any large catalogue (hence not rare objects)  and they imply that the thick disk indeed extends to at least solar metallicities.  Finally, it is also worthwhile to note, as pointed out by e.g. \\cite{nordstrom2004}, that there is a lot of kinematic structure when studying large samples of nearby stars. One such structure, or group of stars, of particular interest here is the Hercules stream which has kinematic properties very similar to some of the thick disk stars. \\cite{famaey2005} studied a sample of  $\\sim 6700$ nearby K and M giants and were able to identify several such structures (including the Hercules stream) for which they determined kinematic properties as well as local number densities. From their results we see that it is especially Hercules stream stars with high $U_{\\rm LSR}$ velocities in the direction  away from the Galactic centre that could  be erroneously classified as thick disk stars. Using a compilation of chemical abundances \\cite{soubiran2005} find that the abundance trends for the Hercules stream mostly follow those in the thin disk.  We have checked for possible contamination from Hercules stream  in our samples. This was done by recalculating the probabilities that we used when selecting our samples  \\citep[see][]{bensby2003, bensby2005}, and now including the properties for the Hercules stream (see Table~\\ref{tab:kinematics}). Using the same equations as in \\cite{bensby2003, bensby2005} (but now  modified to also include $U_{\\rm lag}$ and $W_{\\rm lag}$) we find that none of our stars are more likely to belong to the Hercules stream than to the thin or thick  disks.  For the discussion of our carbon results from this paper we will assume that our stellar samples are representative for the thin and thick disks, or the stellar components that can be kinematically associated with the two disks disks, and that the thick disk does extend up and until solar metallicities.  \\subsection{Observed trends} \\label{sec:trends}  \\begin{figure*} \\resizebox{\\hsize}{!}{\\includegraphics[bb=18 144 592 485,clip]{FIG_11a.eps} \\includegraphics[bb=18 144 592 485,clip]{FIG_11b.eps}} \\resizebox{\\hsize}{!}{\\includegraphics[bb=18 144 592 530,clip]{FIG_11c.eps} \\includegraphics[bb=18 144 592 530,clip]{FIG_11d.eps}} \\caption{ {\\bf a)} and {\\bf b)} show the carbon and oxygen abundances, respectively, relative to iron. {\\bf c)} and {\\bf d)} show the carbon and iron abundances, respectively, relative to oxygen. All carbon abundances are from this work and the oxygen and iron abundances have been taken from Bensby et al.~(2003, 2004, 2005), see also Table~\\ref{tab:sampleA}.  Thin and thick disk stars are marked by open and filled circles, respectively.  The few thick disk stars from Nissen et al.~(2002) included in the oxygen plots are marked by filled triangles. } \\label{fig:trender} \\end{figure*}  \\subsubsection[]{[C/Fe] versus [Fe/H]}  In Fig.~\\ref{fig:trender}a we show the trend of [C/Fe] versus [Fe/H]. The thin and thick disk [C/Fe] trends are fully merged and at sub-solar metallicities the [C/Fe] values all fall within $\\rm 0 \\lesssim [C/Fe] \\lesssim 0.2$ with no particular slope. At super-solar metallicities the [C/Fe] values decrease with increasing [Fe/H] (independent of if the Fe\\,{\\sc i} blend is included or not, compare Figs.~\\ref{fig:chdiff_noblend}d and e).  The flat [C/Fe] trend that we see at $\\rm [Fe/H]<0$ is what generally is found, within the uncertainties, in the literature  \\citep[see, e.g., compilation by][]{gavilan2005}.  \\cite{gustafsson1999}, however, find an increase in [C/Fe] for decreasing [Fe/H]. This is somewhat discomforting since both our and their carbon abundances are based on the forbidden line at 872.7\\,nm.  However, at a closer look at the [C/Fe] versus [Fe/H] trend in  \\cite{gustafsson1999} it is possible that ours and theirs trends are not that different. That they see a slope seem to be mainly based on a few ($\\sim 5$) stars with $\\rm [C/Fe]>0.2$ at $\\rm [Fe/H]\\lesssim -0.6$ (see their Fig.~4). So, {\\it if} those stars are disregarded, our  and their [C/Fe] versus [Fe/H] trends are, within the uncertainties,  similar. Good agreement to \\cite{gustafsson1999} is further strengthened from the detailed comparison in  Sect.~\\ref{sec:systematicerrors}.  \\subsubsection[]{[C/O] versus [O/H]}  Fig.~\\ref{fig:trender}c shows the [C/O] versus [O/H] trend for our stellar sample. The thin and thick disks show trends that are clearly separated.   The thin disk shows a shallow increase in [C/O] with [O/H]  whilst the thick disk first has a flat [C/O] trend that increases sharply at  $\\rm [O/H]=0$. The great resemblance with the observed [Fe/O] versus [O/H] trends shown in Fig.~\\ref{fig:trender}d could indicate that C and Fe indeed originate from objects that evolve on similar time scales. As our stellar sample spans a rather limited range in [O/H] it is difficult to use this sample only to say how the [C/O] trend for the thick disk would have evolved from very low oxygen abundances (i.e. $\\rm [O/H]\\lesssim-0.4$).  In Fig.~\\ref{fig:akerman} we therefore show our [C/O] versus [O/H] trend together with the stellar sample (consisting of stars with halo kinematics)  from \\cite{akerman2004}. They based their carbon and oxygen abundances on permitted C\\,{\\sc i} and O\\,{\\sc i} lines. As they assumed that the NLTE corrections for the C and O abundances for these lines are of similar size they did not apply any NLTE corrections. However, recent NLTE  calculations by \\cite{fabbian2005} have shown that these ``permitted''  [C/H] values should be decreased by as much as $\\sim 0.4$\\,dex  at  $\\rm [O/H]\\approx-2.3$ and less at higher  [O/H]. The NLTE corrections for [O/H] are slightly less ($-0.1$ to  $-0.3$\\,dex, depending on stellar parameters) resulting in the [C/O] being overestimated by $\\sim 0.2$\\,dex at the lowest [O/H] in the \\cite{akerman2004} data \\citep[see also][]{asplund2005araa}. As a consequence the upturn seen in [C/O] with decreasing [O/H] for $\\rm [O/H]\\lesssim-1$ in the \\cite{akerman2004} data is overestimated - their trend should be flatter.  \\subsection{Where is carbon made?}  The carbon yields for various objects differ significantly.  The yields are sensitive to a number of factors such as: stellar winds, treatment of convection, and the $^{12}$C($\\alpha$,$\\gamma$)$^{16}$O rate.  Earlier  models of massive stars found them to give high yields of C \\citep[e.g.,][]{maeder1992} whilst these were adjusted downwards once rotation was taken into account \\citep{meynet2002}. Several studies have used these yields to model galactic chemical evolution and compare it with observed  abundance trends (both in local stars as well as in Galactic and  extra-galactic H\\,{\\sc ii} regions). \\cite{gustafsson1999} gave an exhaustive list and discussion of  possible sites for carbon. These include: supernovae, novae, Wolf-Rayet stars, low and intermediate mass stars in the planetary nebula phase or by super-winds at the end or the red-giant phase. Through an  analysis of Galactic dwarf stars  (a mixture of thin disk stars, metal-poor thick disk stars, as well as stars with intermediate kinematics from Edvardsson et al. 1993)  and data from dwarf galaxies they arrived at the conclusion that the main contribution of carbon to the galactic chemical evolution comes from super-winds of metal-rich massive stars, that low mass stars were ruled out, and that the  r\\^ole of the intermediate mass stars remained unclear. \\cite{henry2000} modelled a large set of Galactic and extra-galactic H\\,{\\sc ii} regions  and reached a similar conclusion. Other studies find that the low  and intermediate mass stars are the essential contributors to the  chemical enrichment of the Galaxy. By adopting the new yields by \\cite{meynet2002}, \\cite{chiappini2003b} conclude that massive stars do contribute to the carbon enrichment but that the main contribution at higher metallicities must be due to low and intermediate mass stars.  Using the same set of yields \\cite{chiappini2003} are also able to explain the abundance ratios in other galaxies, such as irregulars. The recent study by \\cite{carigi2005} arrives at a  similar complex explanation for the observed abundance patterns where  metallicity dependent yields are necessary both for the massive stars (increasing yields with metallicity) as well as for low and intermediate mass stars (decreasing yields with  increasing metallicity). In their best-fitting model  the massive stars dominate the  C enrichment at early times and the low mass and intermediate mass starting to contribute later. The contributions at later times are roughly equal for the two sources. As \\cite{carigi2005} had to invoke the yields by \\cite{maeder1992} \\citep[that no longer should be used as they are superseded by][]{meynet2002} in order to reproduce the solar [C/O] ratio the importance they attribute to massive stars at higher metallicities could be overestimated.  \\begin{figure*} \\resizebox{0.8\\hsize}{!}{ \\includegraphics[bb=40 170 610 385,clip]{FIG_12_new.eps}} \\caption{[C/O] versus [O/H]. Our thin and thick disk stars are marked by open and filled circles, respectively. Halo stars from Akerman et al.~(2004) are marked by crosses. Note that the Akerman et al.~(2004) [C/O] values are overestimated by $\\sim 0.2$\\,dex at the lowest [O/H] due to neglect of NLTE effects. The effect decreases as [O/H] increases. } \\label{fig:akerman} \\end{figure*}   Our [C/Fe] vs [Fe/H] trend in the thick disk is flat whilst  our [O/Fe] vs [Fe/H] trend (see Figs.~\\ref{fig:trender}a and b, respectively)  shows a clear break and subsequent downturn at $\\rm [Fe/H] = -0.5$ (this corresponds to $\\rm [O/H]\\approx0$).  This type of  downturn is normally interpreted as the onset of SN\\,Ia \\citep[see, e.g.,][]{tinsley1979, matteucci2001}  and since oxygen is only produced in SN\\,II the increased Fe production by SN\\,Ia results in a downward trend as [Fe/H] increases.  For [C/Fe] vs [Fe/H] we see no such trend. Hence, we may infer from the C abundances in our thick disk sample that enrichment from carbon happens on the same time scale as the enrichment from SN\\,Ia. Hence the increase in Fe production is matched by enrichment of C.  The fact that our [C/O]\\,-\\,[O/H] and [Fe/O]\\,-\\,[O/H] trends are indeed very similar further strengthens the importance of low and intermediate mass stars as contributors to the carbon enrichment at higher metallicities. This is also supported by \\cite{chiappini2003b} who predict, if low and intermediate mass stars are important, that there first would be a flat [C/O] plateau for the thick disk that then (at roughly solar [O/H]) should sharply increase, to be followed by a shallow thin disk [C/O] trend at higher [C/O] ratios. And this is just what our data show.  Further, \\cite{akerman2004} note that the amplitude of the rise that they see in the [C/O] ratio around $\\rm [O/H]=-1$ in their Milky Way halo stars (see Fig.~\\ref{fig:akerman}) is well matched by similar trends for H\\,{\\sc ii} regions in nearby spiral and irregular galaxies. This they interpret as favourable for their findings that massive stars are the main contributor to the chemical enrichment of carbon as it would be very unlikely that the abundance trends in galaxies that have experienced different types of chemical histories should be so similar if carbon was mainly produced in sites where a substantial time-delay had to be invoked, e.g. low-mass stars. Given the uncertainties relating to NLTE corrections of the \\cite{akerman2004} data (see discussion above) and that \\cite{chiappini2003} are able to explain the abundance ratios in other galaxies as well (including irregulars), using the same set of yields that were used for the Milky Way, this indicates that the r\\^ole of massive stars might be overestimated by \\cite{akerman2004}.  If we compare our C abundances with Y the  trend of [C/Y] versus [Y/H] shows a flat trend for $\\rm [Y/H]\\lesssim 0$ which turns to a steadily declining trend for $\\rm [Y/H]\\gtrsim 0$ (see Fig.~\\ref{fig:cyfeo}a) which is very similar to what we see for [C/Fe] versus [Fe/H] (see Fig.~\\ref{fig:trender}a).  [C/Y] vs [Fe/H] is somewhat different, being, essentially,  flat for both the thin and the thick disk samples at all metallicities (see Fig.~\\ref{fig:cyfeo}b).  Approximately 70\\% of Y is produced in the s-process in AGB stars and  $\\sim 30$\\% comes from massive stars ($\\sim 10$\\% from a weak ``secondary''  s-process component and $\\sim 20$\\% from a primary component \\citealt{travaglio2004}).  The most straightforward interpretation of the lack of trends would  be that the two elements, C and Y, are made in objects that enrich the interstellar medium on the same time scale. Most naively we would  assume that their major components  are indeed made in the same objects, namely low and intermediate mass stars in the AGB phase.  However, if we inspect the  trend of [C/Y] versus [Y/H] in addition to the flat trend below sub-solar  [Y/H] we see a declining trend of [C/Y] for [Y/H]$>0$ indicating that as the Y increases there is a decrease in C production. This could be due to a metallicity effect in the production of s-process elements that favours light s-process elements (such as Y) over heavy s-process elements (such as Ba) at high metallicities \\citep{busso2001,bensby2005}.  So, we have shown that the carbon production must balance the the Fe production by SN\\,Ia and the Y production by AGB stars. With this observable in mind it appears that either the massive stars must have (very) metallicity dependent yields that are finely tuned to the [Fe/H] timescale of SN\\,Ia as observed in the Milky Way thick disk or that less massive stars are significant contributors too.  In summary we thus find that the carbon enrichment at low metallicities (i.e. the Galactic halo and metal-poor thick disk) is due to massive stars but as more and more low and intermediate mass stars start to contribute to the general enrichment they also more and more dominate the carbon enrichment of the interstellar medium as it is seen in the thin disk and the metal-rich thick disk.  \\begin{figure} \\resizebox{\\hsize}{!}{\\includegraphics{FIG_13.eps}} \\caption{ {\\bf a)} [C/Y] versus [Y/H] and {\\bf b)} [C/Y] versus [Fe/H]. Yttrium abundances are taken from Bensby et al.~(2005). Thin and thick disk stars are marked by open and filled circles, respectively. } \\label{fig:cyfeo} \\end{figure}    We present carbon abundances in 51 F and G dwarf stars (16 thick disk stars and 35 thin disk stars)in the solar neighbourhood. The analysis is based on the forbidden [C\\,{\\sc i}] line at 872.7\\,nm which is an abundance indicator that is insensitive to errors in the stellar atmosphere parameters. Combining these data with our previously published oxygen abundances, based on the forbidden [O\\,{\\sc i}] line at 630.0\\,nm \\citep[see][]{bensby2004}, we can form very robust [C/O] ratios that we then used to investigate the origin of carbon and the chemical evolution of the Galactic thin and thick disks. Our findings concerning carbon abundance trends in the Galactic thin and thick disks include: \\smallskip\\\\\\noindent$\\bigstar$ At sub-solar [Fe/H] we find that the [C/Fe] versus [Fe/H] trends for the thin and thick disks are fully merged and essentially flat. For the thin disk, that extends to higher metallicities, [C/Fe] is flat until $\\rm [Fe/H]=0$ when a slow decline starts (see Fig.~\\ref{fig:trender}a). \\smallskip\\\\\\noindent$\\bigstar$ Our abundance trends indicate that the sources that contribute to the carbon enrichment of the interstellar medium do so on the same time-scale as those that produce most of the iron, i.e. SN\\,Ia. The production of C and Y (coming mainly from AGB stars) also seem to work on the same time-scale. \\smallskip\\\\\\noindent$\\bigstar$ In light of our own as well as other studies in the literature \\citep[notably][]{gustafsson1999, chiappini2003, chiappini2003b, akerman2004, carigi2005} we feel that the source(s) of carbon is not yet settled but that there is growing evidence that a complicated, and finely tuned, set of objects contribute to the enrichment of carbon in galaxies. As discussed, based on our own results only we would conclude that the main source for carbon in the Galaxy is low and intermediate mass stars. However, for the Milky Way galaxy it appears that massive stars played a significant r\\^ole for the carbon enrichment at low metallicities (i.e. halo and metal-poor thick disk) whereas low and intermediate mass stars dominate more and more at higher metallicities, i.e. that they have been the major contributors to the carbon enrichment in the thin disk and the metal-rich thick disk. \\smallskip\\\\\\indent From our analysis of the forbidden [C\\,{\\sc i}] line at 872.7\\,nm to determine carbon abundances in solar-type stars we also found a few details that we think are noteworthy: {\\bf 1)} By examining how the neglect of the blending Fe\\,{\\sc i} line effects the derived carbon abundances from the forbidden [C\\,{\\sc i}] line we find that it is mainly for stars with effective temperatures less than $\\sim 5700$\\,K that differences are large (see Fig.~\\ref{fig:chdiff_noblend}a); {\\bf 2)} Currently, the only available source for the oscillator strength of the blending Fe\\,{\\sc i} line is a theoretical one \\citep{kurucz1993}. Therefore, it is possible that the strength of the Fe\\,{\\sc i} could under- or overestimated. By working strictly relative to the Sun the effects of an erroneous treatment of the blend will be reduced. {\\it It would, however, be very valuable if a $\\log gf$-value for the Fe\\,{\\sc i} line could be measured in the laboratory}; {\\bf 3)} For the Sun we find a carbon abundance of $\\rm \\epsilon_{\\odot}(C)=8.41$ when including the Fe\\,{\\sc i} blend and $\\rm \\epsilon_{\\odot}(C)=8.44$ when neglecting it; {\\bf 4)} The general appearance of our carbon trends in the two disks is similar whether we include the Fe\\,{\\sc i} blend in the abundance analysis or not (see Figs.~\\ref{fig:chdiff_noblend}d and e)."
    },
    "0601/astro-ph0601289_arXiv.txt": {
        "abstract": "This paper summarizes a three night observing campaign aimed at achieving milli-magnitude precision photometry of bright stars (V $<$ 9.0) with the 1-meter Swope telescope at Las Campanas Observatory. The test targets were the main sequence stars HD205739 and HD135446. The results show that, by placing a concentric diaphragm in front of the aperture of the telescope, it is possible to avoid saturation and to achieve a photometric precision of 0.0008--0.0010 mag per data point with a cadence of less than 4 minutes. It is also possible to reach an overall precision of less that 0.0015 mags for time series of 6 hours or more. The photometric precision of this setup is only limited by scintillation. Scintillation could be reduced, and therefore the photometric precision could be further improved, by using a neutral density filter instead of the aperture stop. Given that the expected median depth of extrasolar planet transits of about 0.01 mags, and their typical duration of several hours, the results of this paper show that 1-m telescopes equipped with standard CCDs can be used to detect planet transits as shallow as 0.002 mags around bright stars. ",
        "introduction": "\\label{sec:intro}  High precision photometry of bright stars is becoming a subject of increasing interest on extrasolar planet studies. Precision photometry complements the Doppler observations by establishing whether the observed radial velocity variations are caused by the reflex motion induced by a planetary companion or by stellar magnetic activity (e.g., Henry et al. 2000a, Queloz et al. 2001, and Paulson et al., 2004). Precision photometry can also determine the presence or absence of transits in stars with known planets. The presence of a transit provides additional information about the planet by allowing the determination of its radius and its absolute mass (e.g., Charbonneau et al. 2000, Henry et al., 2000b, Sato et al. 2005, and Charbonneau et al. 2005), and also more precise orbital parameters of the system.  Photometric precision of a few milli-magnitudes is routinely achieved using ground-based telescopes and standard CCDs for stars fainter than V = 9.0--10.0. Gilliland et al.(1991) showed that, using a 1-m telescope and a standard CCD, it is possible to achieve a precision of 0.0008 mags for 13th magnitude stars, in time series of about 9 hours with 2-minute candence. Everett \\& Howell (2001) have more recently demonstrated that it is possible to reach precisions of 0.00019 mags in a 4.5-hour time series for V = 14.0 stars with a 0.9-m telescope and a mosaic CCD. However, for stars V = 9.0 mag or brighter milli-mag precision photometry basically remains the realm of photomultiplier tubes\\footnote{The most successful ground-based experiment using photomultipliers is the automatic photometric telescopes (APTs) at Fairborn Observatory (Henry 1995a, Henry 1995b, Stammeir et al. 1997, Henry 1999, and Eaton, Henry \\& Fekel 2003). The APTs routinely achieve photometric precision of 0.001 mag for single observations of stars brighter than V = 8.5. They are also the only instruments that have maintained that precision for more than ten years, what makes them a superb tool.}. A typical CCD inmmediatly saturates when we place it in a 1-m class telescope and try to image a bright star. Saturation can be avoided by using a smaller telescope, but in that case atmospheric scintillation limits the photometric precision to 3--5 millimags.   Efforts to reach milli-mag precision photometry of bright stars with CCDs on ground based 1-m class telescopes are now beginning to produce results comparable to those of photomultipliers tubes. For example, Bouchy et al. (2005) recently published the detection of an extrasolar planet transit around HD189733 with the 1.2 m telescope at the Haute-Provence Observatory in France, using a 1024 x 1024 SITe back-illuminated CCD. The depth of this transit is 0.03 mags and their photometric precision per data point varies between 0.002 and 0.003 mags. Similarly, Charbonneau et al. (2006) have confirmed the 0.003 mag deep transit of HD149026b (Sato et al. 2005) using a 2-chip 2048 x 4608 CCD installed at the 48-inch (1.2m) telescope of the F. L. Whipple Observatory at Mount Hopkins, Arizona. Most efforts have been focused on the Northern Hemisphere, where small telescopes are more at hand. However, at least half of the sample of stars currently searched for planets lie at southern latitudes. In this paper I present the results of tests performed at the 40-inch (1m) telescope in Las Campanas Observatory, Chile. The tests were intended to achieve milli-mag precision photometry of southern bright currently being searched for planets by the radial velocity surveys. The selected targets were HD205739 (V = 8.5) and HD135446 (V = 8.2). Section 2 describes the instrumental setup and the observations. Section 3 summarizes the data analysis and presents the resulting light curves. Section 4 compares our results to the three currently known exoplanet transits around bright stars. Finally the summary of this work and future plans are given in section 5. ",
        "conclusions": "This paper shows the results of a three-day test campaign aimed at achieving milli-magnitude precision photometry of bright southern stars (V $<$ 9.0) with a 1-m telescope and a standard CCD. The targets for these tests were HD205739 and HD135446, two 8th magnitude stars with no previous record of photometric variability.  Stars that bright would inevitably saturate given the collecting area of the telescope and the minimum exposure time of 5 sec imposed by the displacement rate of the CCD shutter. In the absence of neutral density filters at Swope, saturation can be avoided by implementing an aperture stop to reduce the aperture of the telescope, dimming the stars by two magnitudes, and making it possible to image objects as bright as V = 6.5 without saturating.  The observations consisted of follow-up photometry of HD205739 over three nights, for a total of 16.95 hours. HD135446 was also monitored over two nights, for a total of 6.2 hours. Special care was taken to keep all sources of noise below the aimed photometric precision goal of 0.001 mags. The total number of counts in both the targets and the comparison stars was at least $10^6$ photo-electrons. The level of counts per pixel never exceeded the linearity limit of the CCD. The level of counts in the combined bias and master flat-field calibration frames was also at least $10^6$ photo-electrons. The stars were monitored only at airmasses under 1.65 to reduce differential extinction effects. The stars were also placed and kept over the same pixels each night to minimize errors caused by the pixel-to-pixel sensitivity variation of the CCD. Finally, we chose as comparison stars nearby objects with magnitudes and colors similar to the targets to minimize the effects of extinction.  The aperture photometry was performed by iteratively generating differential photometry light curves for different combinations of apertures of the target and the comparison star. The best light curves were then selected as the ones resulting from the combination of apertures that gave the smallest standard deviations.  The resulting light curves show that it is possible to achieve precisions between 0.0008 and 0.0010 mags in individual points with a cadence of less than 4 minutes. The results also show that it is possible to keep the level of precision of the light curves below 0.0013-0.0015 for extended periods of time (at least 6 hours).  The precision of the light curves is limited by atmospheric scintillation, which with the diaphramed aperture of the telescope and the current exposure times (20-30 sec) introduces an average photometric variation of 0.0012 mags. Scintillation depends heavily on the diameter of the telescope. Its effect can be therefore reduced to 0.0008 mags if we replace the aperture stop by an inexpensive neutral density filter. The neutral density filter will have the same advantage of dimming the stars by 2 magnitudes, as the aperture stop, but will not have the inconvenience of the lose in telescope aperture that enhances the atmospheric scintillation.  These tests demonstrate that the Swope telescope can be used to obtain high precision photometry of southern bright stars, with very small alterations to its current setup. In particular Swope can be used to detect transits of close-in giant planets around bright planet-hosting stars. For example, Swope would had easily detected the transits of the planets HD189733b, with a depth of 0.030 mags, HD209458b, with a depth of 0.015 mags, and even the only 0.003-mag deep transit of HD149026b. There are currently no facilities in the Southern Hemisphere pursuing a targeted planet transit search among known planet-bearing stars. Swope is becoming therefore the first one of those facilities.  Future plans include the replacement of the aperture stop by neutral density filters and the extension of these tests to other wavelengths, mainly UBVRI Johnson filter passbands. Another technical goal is to try to improve the duty cycle of the observations. In addition, we are starting a targeted search for transits around the southern stars with known planets that have not been monitored yet."
    },
    "0601/astro-ph0601240_arXiv.txt": {
        "abstract": "{ Spectroscopic observations of hot stars belonging to the young cluster LMC-NGC\\,2004 and its surrounding region were carried out with the VLT-GIRAFFE facilities in MEDUSA mode. We determine fundamental parameters (\\teff, \\logg, \\vsini, and radial velocity), for all B and Be stars in the sample thanks to a code developed in our group. The effect of fast rotation (stellar flattening and gravitational darkening) are taken into account in this study. We also determine the age of observed clusters. We compare the mean \\vsini~obtained for field and cluster B and Be stars in the Large Magellanic Cloud (LMC) with the ones in the Milky Way (MW). We find, in particular, that Be stars rotate faster in the LMC than in the MW, in the field as well as in clusters. We discuss the relations between \\vsini, metallicity, star formation conditions and stellar evolution by comparing the LMC with the MW. We conclude that Be stars begin their Main Sequence life with an initial rotational velocity higher than the one of B stars. It is probable that only part of the B stars, with a sufficient initial rotational velocity, can become Be stars. This result may explain the differences in the proportion of Be stars in clusters with similar ages. ",
        "introduction": "The study of physical properties of B stars with respect to the relative frequency of Be stars in young open clusters and their surrounding field can provide important insights into the origin of the Be phenomenon. Indeed, our knowledge of mass-loss conditions, which lead to the formation of an anisotropic envelope around a fraction of B stars, is still very poor.  Several physical processes could be involved, such as rapid rotation, non-radial pulsations, magnetic fields, evolutionary effects, binarity... Although the Be phenomenon has frequently been considered as a rapid rotation-related phenomenon in OB stars, it is not yet established whether only a fraction or all of the rapidly rotating early-type stars evolve into Be stars.  Several recent studies tackled the major question of the influence of metallicity and evolution on rapid rotators.  The Be phenomenon could be favoured in stars of low metallicity (Maeder et al. 1999). Preliminary results by Royer et al. (2004) could confirm this metallicity effect, although they use limited stellar samples.  The appearance of the Be phenomenon could also depend on stellar ages (Fabregat \\& Torrej\\'on 2000). Accounting for the effects of fast rotation and of gravitational darkening, Zorec et al. (2005) concluded  that Be stars spread over the whole Main Sequence (MS) evolutionary phase. However, they find that in massive stars the Be phenomenon tends to be present at smaller \\ttms age ratios than in less massive stars. Moreover, Keller (2004) found that young clusters host more rapid rotators than their surrounding field. This assessment is valid for the LMC, as well as in the Galaxy. A similar trend was found by Gies \\& Huang (2004) from a spectroscopic survey of young galactic clusters.  Up to now, spectroscopic surveys were only obtained with a resolution power R$<$5000 and/or a low signal to noise ratio (S/N). The new instrumentation FLAMES-GIRAFFE installed at the VLT-UT2 at ESO is particularly well suited to obtain high quality spectra of large samples needed for the study of stellar populations. Thus, we have undertaken the determination of fundamental parameters for a large sample of B and Be stars in regions of different metallicity: (i) to check whether the low metallicity favours the formation of rapid rotators and in particular of Be stars; and (ii) to investigate the evolutionary status of Be stars. In a first paper (Martayan et al. 2005, hereafter Paper I), we reported the identification of numerous B-type stars, the discovery of new Be stars and spectroscopic binaries in the young cluster LMC-NGC\\,2004 and its surrounding field  with the help of medium resolution spectra obtained with the FLAMES instrumentation. The present paper deals with fundamental parameters and evolutionary status of a very large fraction of those objects, taking into account rotational effects (stellar flattening, gravitational darkening) when appropriate. ",
        "conclusions": "With the VLT-GIRAFFE spectrograph, we obtained spectra of a large sample of B and Be stars in the LMC-NGC\\,2004 and in its surrounding field. We determined fundamental parameters for B stars in the sample, and apparent and parent non-rotating counterpart (pnrc) fundamental parameters for fast rotators such as Be stars.  From the \\vsini~ study for B and Be stars in the LMC and its comparison with the MW, we conclude that Be stars begin their life on the MS with a stronger initial velocity than B stars. Moreover, this initial velocity is sensitive to the metallicity. Consequently, only a fraction of B stars can become Be stars. This result may explain the differences in the proportion of Be stars in clusters with similar age.  Our results support Stepien'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.  We find no clear influence of the metallicity on rotational velocity in B-type stars.  The low metallicity may favour the PMS evolution of high velocity stars by minimizing the braking due to magnetic field interactions with the disk, but the influence of metallicity during the life of B-type stars in the MS is not preponderant. As Be stars are not critical rotators, an additional process, such as magnetic field by transfering momentum to the surface or non-radial pulsations (see Rivinius et al. 1998), must provide additional angular momentum to eject material from the star.  The effects of metallicity, the star formation conditions and the evolutionary status of B and Be stars discussed in this paper will be investigated in a forthcoming paper in the Small Magellanic Cloud, which has a lower metallicity than the Large Magellanic Cloud, in order to enlarge the results presented here."
    },
    "0601/astro-ph0601595_arXiv.txt": {
        "abstract": "We present deep echelle spectrophotometry of the Galactic {\\hii} regions M16, M20 and NGC~3603. The data have been taken with the Very Large Telescope Ultraviolet-Visual Echelle Spectrograph in the 3100 to 10400 \\AA\\ range. We have detected more than 200 emission lines in each region. Physical conditions have been derived using different continuum and line intensity ratios. We have derived He$^{+}$, C$^{++}$ and O$^{++}$ abundances from pure recombination lines as well as abundances from collisionally excited lines for a large number of ions of different elements. We have obtained consistent estimations of the temperature fluctuation parameter, {\\ts}, using different methods. We also report the detection of deuterium Balmer lines up to D$\\delta$ (M16) and to D$\\gamma$ (M20) in the blue wings of the hydrogen lines, which excitation mechanism seems to be continuum fluorescence. The temperature fluctuations paradigm agree with the results obtained from optical CELs  and the more uncertain ones from far IR fine structure CELs in NGC~3603, although, more observations covering the same volume of the nebula are necessary to obtain solid conclusions. ",
        "introduction": "Spectrophotometric studies of Galactic {\\hii} regions provide paramount information for the study of the chemical evolution of our galaxy. The use of echelle spectrographs at large aperture telescopes is a step ahead in the knowledge of the physics and chemical composition of these objects. In the last years, our group and others have obtained deep intermediate and high resolution optical spectra of bright Galactic {\\hii} regions \\citep[e.g.][]{peimbertetal93, estebanetal98,estebanetal99a,estebanetal99b,estebanetal04,garciarojasetal04,garciarojasetal05,tsamisetal03}, and extragalactic {\\hii} regions \\citep[e.g.][]{estebanetal02, apeimbert03, tsamisetal03}. These observations have permitted to obtain accurate measurements of very faint recombination lines (hereinafter RLs) of heavy element ions (especially {\\cii} and {\\oii} lines) avoiding the problem of line blending. A common result of all the studies commented above is that the abundance determinations from RLs are systematically larger than those obtained using the standard methods based on the standard analysis of collisionally excited lines (hereinafter CELs). The discrepancies can be of the order of a factor 2-3. This problem may be related to the so-called temperature fluctuations suggested to be present in ionised nebulae \\citep{torrespeimbertetal80}. In fact, the intensity of CELs is much more strongly dependent on the temperature than RLs. This implies that RLs should be, in principle, more precise indicators of the true chemical abundances of the nebula. Nonetheless, this is still a controversial matter that has not been solved; in fact, alternative causes have been proposed recently, such as e.g. small-scale chemical inhomogeneities \\citep{tsamisetal03, tsamispequignot05}.  This paper presents results of Very Large Telescope (VLT) spectrophotometry obtained with the Ultraviolet-Visual Echelle Spectrograph (UVES) of three bright {\\hii} regions: M16, M20 and NGC~3603. Our dataset, of an unprecedented quality, allows us to derive, with high precision, physical conditions and ionic abundances of many heavy elements and, for the first time in the three {\\hii} regions studied in this paper, C$^{++}$ and O$^{++}$ abundances from pure recombination lines.  M16 --{\\it the Aquila nebula}, NGC~6611--  is a well known {\\hii} region of our galaxy. Several works about star formation in the ``elephant trunks'' have appeared in the last years, most of them related to the evaporating gaseous globules (EGGs) discovered in this object \\citep{hesteretal96,thompsonetal02,mccaughreanandersen02} In particular, recent near-infrared observations have led some authors to suggest that the epoch of star formation in M16 may be near its endpoint \\citep{thompsonetal02,mccaughreanandersen02}. Surprisingly, there are only a few optical spectrophotometric works on the chemical composition of this object, computing the most common abundance ratios \\citep{hawley78, rodriguez99b,deharvengetal00}.  M20 --{\\it the Trifid nebula}, NGC~6514-- is a nearby, small and symmetrical {\\hii} region ionised by the O7 V star HD 164492A. Although many works have been carried out to understand its kinematics \\citep[e.g.][]{bohuski73a,rosadoetal99} and its temperature and density structure \\citep[e.g.][]{bohuski73b,dopita74,copettietal00}, there is only a handful of studies devoted to the chemical composition of this nebula. \\citet{hawley78} and \\citet{rodriguez99b} derived heavy element abundances for M20 from spectrophotometric data. \\citet{lyndsoneil85} used long-slit spectroscopic data and narrow-band photometry to measure the intensities of several bright emission lines across the nebula, and derived He, N, O and S abundances for M20 through the computing of a dusty model. Additionally, because of the particular shape, size and dust distribution of this nebula some works were carried out to explore the interaction between gas and dust in this region \\citep[e.g.][]{krishnaswamyodell67}  Finally, NGC~3603 is the only optically visible, giant {\\hii} region in our Galaxy \\citep{gossradhakrishnan69,balicketal80}. The study of the physical properties of this object is crucial for the knowledge of physical processes in large {\\hii} regions. Many efforts have been developed on the study of the star formation and the stellar content of NGC~3603 \\citep[e.g.][]{brandletal99,stolteetal04} and on the study of its kinematics \\citep[e.g.][]{balicketal80,clayton86, clayton90, nurnbergeretal02}, but few works have studied the chemical properties of the ionised gas of this object using optical data \\citep{melnicketal89,girardietal97,tapiaetal01} and mid and far-infrared data \\citep{lacyetal82, simpsonetal95}.  In \\S\\S~\\ref{obsred} and~\\ref{lin} we describe the observations, the data reduction procedure and the measurement and identification of the emission lines. In \\S~\\ref{phiscond} we calculate electron temperatures and densities using several diagnostic ratios and discuss the {\\ts} results. In \\S~\\ref{abundances} ionic abundances are determined based on both kinds of lines: CELs and RLs. Total abundances are determined in \\S~\\ref{totabun}. Deuterium Balmer lines and their excitation mechanism are discussed in \\S~\\ref{deut}. The detection of velocity components in NGC~3603 is reported in \\S~\\ref{veloc}. Finally, in \\S\\S~\\ref{discus} and ~\\ref{conclu} we present the general discussion and the conclusions, respectively. ",
        "conclusions": "We present deep echelle spectroscopy in the 3100-10400 ${\\rm \\AA}$ range of bright zones of the Galactic {\\hii} regions M16, M20 and NGC~3603. We have measured the intensity of about 250 lines per object. This is the most complete set of emission lines ever obtained for these three objects.  We have derived the physical conditions of each nebula making use of several line intensities and continuum ratios. The chemical abundances have been derived using the intensity of collisionally excited lines (CELs) for a large number of ions of different elements. We have determined, for the first time in the three objects, the C$^{++}$ and O$^{++}$ abundances from recombination lines (RLs) and, finally we have also determined the abundance of O$^{+}$ from RLs for the first time in M20.  We have obtained consistent estimations of the temperature fluctuations parameter, {\\ts}, applying different methods: by comparing the O$^{+}$ (when available) and O$^{++}$ ionic abundances derived from RLs to those derived from CELs; by applying a chi-squared method which minimizes the dispersion of He$^+$/H$^+$ ratios from individual lines; and by comparing the electron temperatures derived from CELs to those derived from Balmer and Paschen continua. The adopted average value of {\\ts} has been used to correct the ionic abundances derived from CELs.  We report the detection of several deuterium Balmer lines in the spectra of M16 and M20. The properties of these lines indicate that fluorescence is their most  probable excitation mechanism.  We have compared the results obtained for optical CELs in NGC~3603 with those obtained from far-infrared fine-structure CELs finding an apparent agreement, in spite of the high uncertainties of the abundances derived from far IR data, if the temperature fluctuations paradigm is assumed. However, IR and optical spectrophotometry covering the same volume of the nebula is necessary to make a reliable and conclusive comparison between optical and IR CEL abundances."
    },
    "0601/astro-ph0601076_arXiv.txt": {
        "abstract": "{We analyzed archival X-ray spectra of MXB 1728-34 obtained in 1996-99 by the Proportional Counter Array on board of the RXTE satellite. X-ray spectra were fitted to our extensive grids of model atmosphere spectra to determine the effective temperature $T_{\\rm eff}$ on the neutron star surface, logarithm of surface gravity $\\log{g}$, and the gravitational redshift $z$ simultaneously. We have chosen fitting by numerical model spectra plus broad Gaussian line, modified by interstellar absorption and the absorption on dust. We arbitrarily assumed either hydrogen-helium chemical composition of a model atmosphere, or H-He-Fe mixture in solar proportion. The statistically best values of $\\log{g}$, and $z$ subsequently were used to determine mass and radius of the neutron star. We obtained the best values of the parameters for the neutron star in X-ray burst source MXB 1728-34: mass either $M=0.40$ or $0.63 M_\\odot$ (for H-He or H-He-Fe models, respectively), radius $R=4.6$ or $5.3$ km, $\\log g=14.6$ or $14.6$ and the gravitational redshift $z=0.14$ or $0.22$. All the above parameters have very wide 1-$\\sigma$ confidence limits. Their values strongly support the equation of state for strange matter in MXB 1728-34.  {\\bf Key words} {\\it Stars: fundamental parameters -- stars: neutron -- X-rays: bursts -- stars: individual (MXB 1728-34) }  } ",
        "introduction": "X-ray source MXB 1728-34 was discovered in 1976 by Forman, Tananbaum \\&Jones (1976) in the survey observations of the Uhuru satellite. Type I X-ray bursts from this source were identified in the same year by Lewin, Clark \\& Doty (1976) in SAS-3 satellite data. Optical counterpart still remains unidentified probably due to large extinction in the direction to this source. Recently Mart{\\'\\i} et al. (1998) has identified an infrared counterpart  to MXB 1728-34, but the connection with the X-ray source has not yet been confirmed. MXB 1728-34 was observed in a wide energy range from radio  (Mart{\\'\\i} et al. 1998) to $\\gamma$-ray energies (cf. Claret et al. 1994).  The determination of mass and radius of a neutron stars is a very important and interesting problem, because it allows one to determine or to constrain the equation of state of superdense matter. Distance to the source is also necessary to localize the source relatively to the Galactic center. For MXB 1728-34 the first estimate of its distance $d$ and true radius $R$ was given by van Paradijs (1978), $d=4.2\\pm 0.2$ kpc and $R=6.5\\pm 0.4$ km (see also Foster et al. 1986). Both parameters were estimated assuming that the peak flux of a given burst from this source approached the Eddington limiting flux. It became immediately clear that such a radius $R$ was smaller than the minimum radius expected for a neutron star of the canonical mass of $1.4 M_\\odot$.  MXB 1728-34 was frequently observed in X-rays after that year, and parameters of the neutron star were estimated in several papers. Kaminker et al. (1989) interpreted spectra of this source using simplified emission models. They have derived the following values of the neutron star parameters: mass $M=1.4-2.0\\, M_\\odot$ and radius $R=6.5-12$ km assuming the distance $d=6$ kpc and pure helium atmosphere.  Di Salvo et al. (2000) and Galloway et al. (2003) also assumed, that luminous bursts from this source are standard candles of the bolometric luminosity at the Eddington limit. They concluded that the distance to this X-ray burster is in the range 4.4 -- 5.1 kpc.  On the other hand, Shaposhnikov et al. (2003) obtained other values: $d=4.5-5.0$ kpc, $M=1.2-1.6 M_\\odot$, and $R=8.7-9.7$ km. They also concluded, that the helium mass abundance ${\\rm Y} >0.9$ in atmosphere of the neutron star in MXB 1728-34. These authors used semianalytical models of expanding atmospheres by Titarchuk (1994), and Titarchuk \\& Shaposhnikov (2002).  A different approach to this problem was used by Li et al. (1999), who interpreted power spectra of the X-ray light curve and compared them to the model by Osherovich \\& Titarchuk (1999) and Titarchuk \\& Osherovich (1999). As the result Li et al. (1999) constrained the area on the ($M$, $R$) plane, where the neutron star of MXB 1728-34 is located. Constrained localization of MXB 1728-34 was characteristic for a star built up of strange matter.  We present here for the first time the estimation of mass and radius of MXB 1728-34, obtained by fitting of archival RXTE X-ray spectra of this source to the new extensive grids of model atmospheres and theoretical spectra of neutron stars. This technique also allowed us to obtain the upper limit of the distance to MXB 1728-34. Theoretical model spectra were obtained with the {\\sc ATM21} code, which computes model atmospheres of hot neutron stars with the account of Compton scattering on free electrons. The code takes into account angle-averaged Compton scattering of X-ray photons with initial energies approaching the electron rest mass.  Detailed description of the equations and numerical methods were given in a long series of earlier papers (Madej 1991a,b; Joss \\& Madej 2001; Majczyna et al. 2002; Madej, Joss \\& R\\'o\\.za\\'nska 2004; Majczyna et al. 2005). \\clearpage ",
        "conclusions": ""
    },
    "0601/astro-ph0601279.txt": {
        "abstract": "We have studied a sample of 809 Mg II absorption systems with 1.0$\\le z_{abs}\\le$1.86 in the spectra of Sloan Digital Sky Survey QSOs, with the aim of understanding the nature and abundance of the dust and the chemical abundances in the intervening absorbers. Normalized, composite spectra were derived, for abundance measurements, for the full sample and several sub-samples, chosen on the basis of the line strengths and other absorber and QSO properties. Average extinction curves were obtained for the sub-samples by comparing their geometric mean spectra with those of matching samples of QSOs without absorbers in their spectra. There is clear evidence for the presence of dust in the intervening absorbers. The 2175 {\\AA} feature is not present in the extinction curves, for any of the sub-samples. The extinction curves are similar to the SMC extinction curve with a rising UV extinction below 2200 {\\AA}. The absorber rest frame colour excess, $E(B-V)$, derived from the extinction curves, depends on the absorber properties and ranges from $<$0.001 to 0.085 for various sub-samples. The column densities  of Mg II, Al II, Si II, Ca II, Ti II, Cr II, Mn II, Fe II, Co II, Ni II, and Zn II do not show such a correspondingly large variation. The overall depletions in the high $E(B-V)$ samples are consistent with those found for individual damped Lyman alpha systems, the depletion pattern being similar to halo clouds in the Galaxy. Assuming an SMC gas-to-dust ratio we find a trend of increasing abundance with decreasing extinction; systems with N$_{\\rm H\\;I} \\sim 10^{20}$ cm$^{-2}$ show solar abundance of Zn. The large velocity spread of strong Mg II systems seems to be mimicked by weak lines of other elements. The ionization of the absorbers, in general appears to be low: the ratio of column densities of Al III to Al II is always less than 1/2. QSOs with absorbers are, in general, at least three times as likely to have highly reddened spectra as compared to QSOs without any absorption systems in their spectra. ",
        "introduction": "\\subsection{The nature of QSO absorbers} Absorption lines in spectra of quasi-stellar objects have been known since shortly after the discovery of QSOs (Sandage 1965; Schmidt 1966; Arp, Bolton \\& Kinman 1967; Burbidge, Lynds \\& Burbidge 1966; Stockton \\& Lynds 1966). While many of the first detections were close to the redshifts of the QSOs, what are now termed associated systems, Bahcall, Peterson \\& Schmidt (1966) found one of the earliest systems which could be termed intervening and Bahcall (1968) established the reality of such systems. Bahcall and Spitzer (1969) presented plausibility arguments that foreground galaxies with extended halos could produce some of the intervening systems. If such were the origin of the QSO absorption line systems (QSOALSs), they would provide an excellent way to understand various aspects of structure formation and evolution in the universe, in particular, the buildup of the elements over cosmic time and evolution of element flows into and out of galaxies as they form (Hermann et al. 2001; Adelberger et al. 2003, 2005; Tumlinson \\& Fang 2005).  There are, indeed, specific examples of QSO absorption lines coming from specific galaxies (Cohen et al. 1987; Yanny, York \\& Gallagher 1989; Bergeron 1988; Chen, Kennicut \\& Rauch 2005). On the other hand, there are numerous cases in which a QSO sightline passes near a visible galaxy without producing absorption lines (Bechtold \\& Ellingson 1992). Some of the galaxies are almost certainly dwarf galaxies (York et al. 1986; Rao et al. 2003; Haehnelt, Stenimetz \\& Rauch 1998). Others may have an outer morphology that resembles the Milky Way (MW) distribution of 21 cm clouds (York 1982; Wakker \\& van Woerden 1997). Such a patchy configuration would explain why absorbers are sometimes found near galaxies and why some QSO/near galaxy coincidences do not produce absorbers. In some cases extended, turbulent regions around star-forming galaxies (Adelberger et al. 2005) may show up as QSOALSs.  Despite numerous observations of absorption line systems in spectra of QSOs, there may be some selection effects in the samples. Why, for instance are QSOALSs with fully or near solar abundances seldom seen (Prochaska et al. 2003a; Khare et al. 2004; Kulkarni et al. 2005a)? Only a few systems with metallicity above solar are definitely known: SDSSJ1323-0021, $z_{abs}$=0.716 (Khare et al. 2004), Q0058+019, $z_{abs}$=0.612 (Pettini et al. 2000).  There may be some selection effects that keep such absorbers out of the observed sample.  For example, dwarf galaxies may constitute a large part of the sample; these usually have sub-solar abundances. Also, absorbers with solar abundances may have more dust, making any background objects fainter and keeping those QSOs out of magnitude-limited samples (e.g. Fall \\& Pei 1993; Boisse et al. 1998;  Vladilo \\& Peroux 2005). Ostriker \\& Heisler (1984) argued that foreground absorbers were probably absorbing light from background QSOs, making them redder. The effect is possibly strong enough to eliminate the QSOs completely from magnitude-limited samples, in some cases, since the attenuation in the rest frame UV is much stronger than that in the rest frame optical range. Akerman et al. (2005) find no evidence for such an effect, but testing on a larger sample could yield different results.  The presence of and the nature of dust in QSOALSs is therefore of interest, as it may shed light on the evolution of dust particles in cosmic time, may allow the evaluation of selection effects and may allow a discrimination between different kinds of dust in the intergalactic medium and in the circum-QSO environment. \\subsection{Inference of dust from gas abundances} Dust has long been suspected to be a constituent of the QSO absorbers given the general connections of galaxies with such systems, the presence of dust in known galaxies and the evidence of depletion of gas phase elements onto dust grains in the QSOALSs (Meyer, Welty \\& York 1989; Khare et al. 2004 and references therein), similar to that seen in the MW (Hobbs et al. 1993; Sembach \\& Savage 1996; Jenkins 2004), Small Magellanic Cloud (SMC, Welty et al. 2001) and Large Magellanic Cloud (LMC, Welty et al. 1999a). The signature of dust depletion is often characterized by the ratio of column densities of dominant ionization stages of a refractory (depleted) species to that of a volatile (weakly depleted) species. Examples include N$_{\\rm Cr\\;II}$/N$_{\\rm Zn\\;II}$, N$_{\\rm Fe\\;II}$/N$_{\\rm Zn\\;II}$ (Khare et al. 2004 and references therein), or N$_{\\rm Fe\\;II}$/N$_{\\rm Si\\;II}$ (Prochaska et al. 2003b), N being the column density, in units of number of ions per square centimeter. More quantitatively, N$_{\\rm Cr\\;II}$/N$_{\\rm Zn\\;II}$ and N$_{\\rm Fe\\;II}$/N$_{\\rm Zn\\;II}$ are smaller than their cosmic abundance ratios by factors of 2-5 in QSOALSs, whereas, in the Galaxy, the values range from 2-100 depending on the kind of gas. Akerman et al. (2005) do not find statistically significant difference in Zn metallicity as well as depletions between their sample of 15 DLAs in the complete optical and radio absorption line systems and the values in optically selected sample published in the literature. \\subsection{Extinction in QSO absorption line systems} Attempts to quantitatively detect dust by its differential extinction of a given object have been difficult. Pei, Fall \\& Bechtold (1991) found higher extinction in a sample of 13 QSOs pre-selected to have damped Ly\\, $\\alpha$ systems (DLAs), compared to a sample of 15 QSOs that had no DLAs and concluded that extinction existed in the first sample, but not the second: a dust-to-gas ratio of 10\\% of the local interstellar medium was found with the absorber rest frame colour excess, hereafter, $E(B-V), < $0.03. Ellison, Hall \\& Lira (2005) similarly find a slight reddening ($E(B-V) <$ 0.04) in 14 out of 42 QSOs with DLAs. The Sloan Digital Sky Survey (SDSS) archive has been used for several studies of extinction in sightlines to QSOs. Richards et al. (2003) showed that an observer frame colour excess, hereafter, $\\Delta(g-i)$, which is the difference between the actual colours of a QSO and the median colours of QSOs at that redshift, could be defined for the SDSS QSO spectra and could be used to form templates of objects with apparent degrees of extinction. However, they could not discern if the extinction so detected was from the QSO itself or from intervening systems. Hopkins et al.  (2004) used composite QSO spectra to show that the extinction towards QSOs is dominated by SMC-like extinction, which they argued was predominantly located at QSO redshifts.  Recently, Wild \\& Hewett (2005) have found evidence of dust in QSOALSs with detected Ca II lines, with $E(B-V)$=0.06, while Wild, Hewitt \\& Pettini (2005) find strong Ca II absorbers to have $E(B-V)>0.1$. Murphy \\& Liske (2004) studied the spectral indices of QSOs in SDSS Data Release 2 and found no evidence for the presence of dust in DLAs at a redshift $\\sim$ 3. They derived an upper limit of 0.02 on $E(B-V)$. \\subsection{Searches for the 2175 {\\AA} bump} The 2175 {\\AA} bump is a feature of the extinction curve of gas in the Local Group. It was discovered in the 1960s using sounding rockets (Stecher 1965, 1969). The origin of the feature is unknown, but it is thought to be made by carbonaceous material in dust or large molecules (Draine 2003). It is well correlated with the total density of hydrogen (atomic and molecular) and is, by inference, an unsaturated signature of material in dust grains or molecules. Since it is known to vary in strength with the steepness of the general extinction (Valencic, Clayton \\& Gordon 2005), knowledge of its behavior is an indicator of different sizes and types of grains. Its absence indicates, in models, a dearth of certain types or sizes of grains in a given locale. Understanding its behavior in other galaxies at different times or under different conditions reveals what is generic and what is particular about the dust in space.  Malhotra (1997) reported detection of the 2175 {\\AA} bump, in the extinction curves of QSOALSs, by averaging 92 QSO spectra in the absorber rest frames (See section 9.2 of this paper for a discussion of this result).  Ostriker, Vogeley \\& York (1990) showed a possible detection of the 2175 {\\AA} bump in a single object.  Recently, the presence of the bump has been recognized in a few QSOALSs (Motta et al. 2002; Wang et al. 2004; Junkkarinen et al. 2004).  Yet another ISM constituent, the diffuse interstellar bands (Herbig 1995) are seen in one system (AO 0235+164, Junkkarinen et al. 2004). \\subsection{Overview of the paper} In this paper, we proceed to use all three techniques noted above to understand dust in QSOALSs: looking for the presence of reddening, searching for the 2175 {\\AA} bump,  and studying the relative abundances of elements in a large and homogeneous sample of QSOs provided by the SDSS (York et al. 2000). It is of particular interest to derive the relationship between these three aspects of dust discernment, to test the universality of the characteristics of dust throughout the Universe. It is also of interest to distinguish dust associated with QSOALSs from the dust near QSOs. Since the presence of dust may affect the selection of QSOs for such a study, it is important to use samples with well-defined selection criteria that can be modeled for the purpose of generalizing conclusions. Finally, by surveying well-defined samples of large numbers of absorbers, albeit at 170 km s$^{-1}$ resolution, we can hope to separate true trends from selection effects that may arise in small samples by accident.  The contents of the paper, by section, are briefly outlined here. Section 2 includes a description of the SDSS survey and spectra, as they relate to the present paper. The selection of 809 absorption systems for the study of the 2175 {\\AA} bump is outlined in section 3. The same sample is used for our derivation of extinction and of line strengths. Averaging procedures used to obtain composite spectra, along with the line identifications, using the full sample as an example, are given in section 4. Two different procedures for finding average extinction, one using QSO spectra and the other using the SDSS colours are given in section 5. Definition of selection of sub-samples, based on entities that can be measured in individual QSOs, along with the physical motivation of those criteria appear in section 6. Table 1, found in section 6, summarizes the sub-samples and their properties, including the color excesses determined as described in section 4. Section 7 contains the central observational results of this study: the effects of extinction are correlated with the absorption line strengths and the derived column densities; the strength of the Mg II absorption lines is defined with regard to extinction and shown not to be a dominant factor; and control samples are used to show that there are only small effects (on colour excess) in redshift, i-magnitude and beta. The extinction is shown to be most like the SMC extinction, with no bump and no inflection below 2000 {\\AA}. The abundances and depletions for our sub-samples are given in section 8, where we identify our highest extinction samples with traditional DLAs (sample of $\\sim$ 100 absorbers) and suggest, assuming an SMC extinction curve, that the largest sub-sample (698 absorbers) has, on average, a sub-DLA (systems with N$_{\\rm H\\;I}$ between 10$^{19}$ to 2$\\times 10^{20}$ cm$^{-2}$) value of N$_{\\rm H\\;I}\\sim 10^{20}$ cm$^{-2}$. With these assumptions, the latter sample shows solar metallicity in the absorbers.  Uncertainties in column densities, abundances and extinction values are explored in section 9. Our conclusions are summarized in section 10. ",
        "conclusions": "We have studied a sample of 809 absorption systems, in the spectra of SDSS non-BAL QSOs, with W$_{\\rm Mg\\;II\\lambda2796}>$0.3 {\\AA}, 1$ \\le z_{abs}<$1.86 and $\\beta>$0.01 to determine the average extinction and the average strengths of absorption lines in QSOALSs.  The extinction curves, obtained by comparing the geometric mean average spectra of the QSOs in 27 sub-samples of absorbers (constructed on the basis of the absorber as well as QSO properties) with those of matching samples of QSOs without any absorbers in their spectra, show clear evidence of extinction by the intervening absorbers. The absorber rest frame colour excess $E(B-V)$ for the sub-samples varies from $<$0.001 to 0.085. The extinction curves uniformly share the main properties of the typical SMC extinction curve: no 2175 {\\AA} bump and rising extinction below 2200 {\\AA}. An average MW curve, similar to those of local interstellar clouds, is excluded. A mix of up to 30\\% of MW curves cannot, however, be ruled out, but the MW curves are not a dominant contributor.  Adapting the SMC extinction curve, the observer frame colours excesses ($\\Delta(g-i)$) were  used, with the absorption redshift, to derive average absorber frame colour excesses, ($<E(B-V)_{\\Delta(g-i)}>$), which are consistent with the values of $E(B-V)$ derived from the extinction curves.  While about 14\\% of the absorber sample has high values of $\\Delta(g-i)$ ($>$0.2), 5\\% of the sample of QSOs with no absorbers also shows such high values. These high values, for the non-absorber QSOs (and some of the absorber QSOs as well) could be caused by strong absorbers with $z_{abs}<$0.5, where SDSS can not detect Mg II doublets or else, there is a class of QSOs with anomalously red continua or there are regions with large amounts of dust associated with gas that is too hot to produce the near UV absorption lines.  The average strengths of absorption lines between 1800 and 3300 {\\AA} in the absorber rest frame, are found to increase as the value of W$_{\\rm Mg\\;II\\lambda2796}$ increases. While Mg II line strengths are generally correlated with reddening, there is large scatter: in particular, of 145 systems with W$_{\\rm Mg\\; II\\lambda 2796}$ above 2.5 {\\AA} in rest frame equivalent width, 2/3 show little reddening ($\\sim$ 0.01 in $E(B-V)$, on average), while the reminder have an average $E(B-V)$ that is eight times higher.  The column densities for Si II, Ca II, Ti II, Cr II, Mn II, Fe II, Ni II and Zn II show a strong trend, increasing slowly with our derived values of $E(B-V)$. If we adopt the SMC, $E(B-V)$ - N$_{\\rm H\\; I}$ relation to derive the mean N$_{\\rm H\\;I}$ for our sub-samples, the abundances derived from the average column densities and average N$_{\\rm H\\;I}$ show a consistent trend of lower abundances with increasing $E(B-V)$. Our largest sub-sample, with $E(B-V)=0.002$ and N$_{\\rm H\\;I} \\sim 10^{20}$ cm$^{-2}$, has near solar zinc and titanium. Systems in this sub-sample (which contains 85\\% of our full sample) are thus sub-DLAs or Lyman-limit systems. For samples with $E(B-V)\\sim$0.08, we find gas depletions with respect to Zn of up to 10 compared to solar abundances. Depletions of Fe, Mn and Cr measured with respect to Zn are similar to those for DLAs as published in the literature. The depletion patterns are similar to those found in the halo gas in the MW.  In general, the ionization corrections appear to be small. The ratio of the column densities of Al III to Al II is smaller than 0.5. The Al III does not seem to be a part of the Al II gas and C IV appears to be separated from both Al II and Al III.  Typical QSOALS have wide blends of velocity components in Mg II $\\lambda\\lambda$2796, 2802, from below our instrumental width of 170 km s$^{-1}$ up to 450 km s$^{-1}$ or more, independent of extinction. These wide systems contribute to the apparent fact that W$_{\\rm Mg\\; I\\lambda 2852}$ is near the linear portion of the curve of growth even for equivalent width $\\sim$ 500 m{\\AA}. Many elements seem to participate in these wide flows and the effects of gas depletion seem to be selective, with Ti mildly depleted amounts in low reddening samples and highly depleted in samples of modest reddening. Since Ti is not yet well observed in studies of individual systems, this effect needs to be explored in more detail.  Studies of systems with low but well defined N$_{\\rm H\\; I}$ are needed, using high $S/N$ and high resolution, to determine what fraction of QSOALS have solar or greater abundances of Zn."
    },
    "0601/astro-ph0601656_arXiv.txt": {
        "abstract": "We present evidence that stars with planets exhibit statistically significant silicon and nickel enrichment over the general metal-rich population.  We also present simulations which predict silicon enhancement of planet hosts within the context of the core-accretion hypothesis for giant planet formation.  Because silicon and oxygen are both $\\alpha$-elements, [Si/Fe] traces [O/Fe], so the silicon enhancement in planet hosts predicts that these stars are oxygen-rich as well. We present new numerical simulations of planet formation by core accretion that establish the timescale on which a Jovian planet reaches rapid gas accretion, $t_{\\rm rga}$, as a function of solid surface density $\\sigma_{\\rm solid}$: $(t_{\\rm rga} / 1 \\; {\\rm Myr}) = ( \\sigma_{\\rm solid} / 25.0 \\; {\\rm g \\: cm}^{-2} )^{-1.44}$.  This relation enables us to construct Monte Carlo simulations  that predict the fraction of star-disk systems that form planets as a function of [Fe/H], [Si/Fe], disk mass, outer disk radius and disk lifetime.  Our simulations reproduce both the known planet-metallicity correlation and the planet-silicon correlation reported in this paper.  The simulations predict that $16\\%$ of Solar-type stars form Jupiter-mass planets, in agreement with $12\\%$ predicted from extrapolation of the observed planet frequency-semimajor axis distribution. Although a simple interpretation of core accretion predicts that the planet-silicon correlation should be much stronger than the planet-nickel correlation, we observe the same degree of silicon and nickel enhancement in planet hosts.  If this result persists once more planets have been discovered, it might indicate a complexity in the chemistry of planet formation beyond the simple accumulation of solids in the core accretion theory. ",
        "introduction": "Since the first recognition that metal-rich stars are more likely to harbor planets (Gonzalez 1997), there have been tantalizing suggestions that planet hosts undergo a different process of chemical development than planetless field stars.  \\cite{gonzalez97} proposed that stars with planets are metal-rich because they self-pollute by planetesimal accretion during the planet-formation epoch.  \\cite{sandquist98} found that accretion of a few tens of earth masses of planetesimals would account for all the metals present in the convective envelope of solar-type stars, but the giant-planet migration mechanism for scattering such a substantial planetesimal mass into the host star is only possible if $90\\%$ of close encounters result in the outward ejection of planetesimals.  \\cite{mc02} confirmed that stars with planets are iron-rich, and proposed that accretion of $5 M_{\\earth}$ of iron, in addition to high primordial metallicity, would explain the trend.  The self-pollution scenario leads to the expectation that planet hosts, although being iron-rich, have reduced abundances of volatiles such as C, N and O, because these would be minimally present in the accreted planetesimals (Smith et al. 2001).  Finding no evidence of volatile depletion among stars with planets (Ecuvillon et al. 2004, Takeda \\& Honda 2005), most investigators have concluded that the iron enrichment of planet hosts is primordial: stars hosting giant planets form preferentially in metal-rich molecular clouds.  This assessment concurs with the finding of \\cite{FV05} that there is no relation between metallicities of planet hosts and convection zone depth, either while the star is on the main sequence or after it evolves to the subgiant branch and its convection zone deepens.  If pollution of stars by planetesimal accretion were responsible for the planet-metallicity correlation, planet hosts with the most shallow convection zones would have the highest metallicity (Laughlin \\& Adams 1997).  Barring planet host self-pollution as the chief cause of the planet-metallicity correlation, there is still the possibility that stars with planets form from molecular clouds with anomalous metal enrichment histories.  \\cite{gchem} calculated the enrichment pattern over time of metals from Li to Zn in sequences [X/Fe] vs. [Fe/H].  For most elements, these theoretical sequences closely match observations. If planet hosts do not lie on the same [X/Fe] vs. [Fe/H] sequences of chemical evolution as other Pop I stars, the enhanced iron abundance seen in planet hosts could be an artifact of this unusual chemical evolution.  If the ability to form planets were uniformly the result of a particular type of chemical event--a nearby supernova, for example--planet hosts would show markedly different abundance distributions than field stars in both volatile and refractory elements. Yet analysis of the trends in [X/Fe] vs. [Fe/H] for $\\alpha$- and Fe-peak elements by \\cite{bodaghee03} reveals that abundance distributions of planet hosts are the high-metallicity extensions of the trends governing chemical evolution in planetless stars.  A similar analysis by \\cite{huang05} suggests that, although S \\& Mg may be enhanced in planet hosts (contrary to Ecuvillon et al. 2004, who found no evidence of sulfur enhancement), they still follow the same slope in the [S/H] vs. [Fe/H] relation as the comparison sample of planetless field stars.  We are left, then, with the conclusion that planet-bearing stars are indistinguishable from other Population I stars in their chemical enrichment histories, implying that no extraordinary chemical events are necessary to stimulate planet formation.  Working with the assumption that planet hosts belong to the same stellar population as other metal-rich field stars, is it still worth checking for patterns in the abundances of stars with planets?  The studies mentioned above use comparison samples of planetless field stars with markedly lower median [Fe/H] than that of the planet hosts. \\cite{bodaghee03} and \\cite{ecuvillon04} have no stars in their comparison sample in the range $0.2 \\leq {\\rm [Fe/H]} \\leq 0.5$, which is where more than $1/3$ of their planet hosts are found.  Using a comparison sample with no stars of ${\\rm [Fe/H]} \\geq 0.1$, \\cite{huang05} could not tell if the Mg and S enhancement in their planet hosts was particular to stars with planets, or a characteristic of metal-rich stars in general.  While comparison samples of this type allow the investigator to comment on the general chemical evolution trends of planet hosts, they cannot elucidate whether planet hosts systematically possess more or less of a certain element than field stars {\\it of the same metallicity.}  Comparison between planet hosts and field stars of the same metallicity probes a third type of relationship between element abundances and the presence or absence of planets: stars with planets do not self-enrich at birth, are not born in molecular clouds with chemically tumultuous histories; but they nevertheless may be enhanced or deficient in some element that aids the formation of planets, {\\it simply due to the natural Galactic variation in stellar abundances.}  The chemical compositions of metal-rich stars are strikingly uniform: in the sample of metal-rich stars observed by \\cite{SPOCS}, [X/Fe] at a given value of [Fe/H], where X = Na, Si, Ti and Ni, is observed to vary by 0.4 dex at the very most (Valenti \\& Fischer 2005).  However, \\cite{FV05}, in their study of the planet-metallicity correlation, find that the probability of planet detection increases by a factor of five if when iron abundance is enhanced by a factor of two (0.3 dex).  If a correlation of this magnitude exists between planet detectability and another element besides iron--or if iron abundance derives its power as a predictor of planet presence from its correlation with the abundance of another element of more physical importance to the planet-formation process--it will be detectable in the 0.4 dex spread of [X/Fe] in stars with the same value of [Fe/H].  In this work, we present our statistical method for assessing the relationship between planet presence and the observed abundances of individual refractory elements. As specific examples of the method, in \\S \\ref{observed}, we present statistical evidence of possible silicon and nickel enhancement in planet hosts, and we show additionally that titanium abundances display no such correlation.  In \\S \\ref{expsim}, we empirically determine the exponent of a power law giving probability of planet formation as a function of [Si/Fe], among stars of the same [Fe/H].  Finally, in \\S \\ref{model}, we construct a simulation based on the core accretion theory that reproduces both the planet-metallicity correlation and our observed correlation between silicon abundance and presence of planets. ",
        "conclusions": "We find that there is statistical evidence for silicon and nickel enrichment of planet hosts in the SPOCS data of \\cite{SPOCS}.  Although the planet hosts do not exhibit any anomalous abundance patterns that indicate that they might be members of a different stellar population than other field stars, they do have values of [Si/Fe] and [Ni/Fe] that are systematically enhanced over other stars of the same [Fe/H].  We have constructed Monte Carlo simulations that predict the likelihood of forming planets by core accretion in disks of differing mass, lifetime, outer radius and chemical composition.  Our simulations reproduce both the planet-metallicity correlation and the planet-silicon correlation reported in this work.  The simulation demonstrates that the abundance patterns of planet hosts in the SPOCS data are consistent with planets having formed by core accretion.  According to the core accretion theory, planets form most easily in ice-rich disks.  We predict that the observed silicon enhancement of planet hosts arises primarily from the correlation in the abundances of silicon and oxygen, and secondarily from silicon's own importance as a grain-forming material.  The planet-silicon correlation would naturally arise if the core of Jupiter is made mainly of water ice frozen on silicate and/or iron grains.  (See, however, \\cite{tar}, which argues that Jupiter's core may have formed from a tar-like mixture of organic compounds.)  More observations are needed to determine whether the planet-silicon correlation is robust, and whether it truly takes the form predicted by our simulations.  Despite its large size, the SPOCS catalog is seriously affected by small-sample statistics when the planet hosts are separated into [Fe/H] bins to have their abundance patterns analyzed.  The tantalizing suggestion that a silicon correlation may exist in the data underscores the need for a larger pool of stars from which to draw planet hosts and control sets. Our simulation predicts that the silicon effect should persist in a larger sample.  Although we do not directly simulate variations in nickel abundance, the idea that nickel enhancement increases likelihood of planet formation is consistent with the core accretion theory.  However, if all solids form protoplanetary cores equally well, the planet-silicon correlation should be much stronger than the planet-nickel correlation.   The most immediate test of our simulation is whether planet hosts show oxygen enhancement.  Our simulation relies on the assumed correlation between oxygen and silicon abundances, so our analytical framework would be invalidated if a SPOCS-type survey shows no evidence for oxygen enrichment in planet hosts.  Other $\\alpha$-elements should also be correlated with the presence of planets, particularly the important core-forming materials C and Mg.  The question of what fraction of Pop I stars have at least one giant planet has been much debated recently.  Our prediction of planetary systems around $16\\%$ of FGK dwarfs roughly agrees with the \\cite{marcy05} extrapolation of the planet frequency-semimajor axis relation in the Lick/Keck/AAT planet search data.  According to our calculation of the fraction of giant planets that experience Type II migration, planets on Jupiter-like orbits account for at least $4 \\%$ of giant planet systems.  A Solar system-like configuration of planets, while not particularly common, is nevertheless a viable outcome of the core accretion-migration scenario of planet formation.   Detection of planets by transit searches and direct imaging will aid in solving the problem of completeness in discoveries of planets at 3-20 AU orbits, since these survey methods are sensitive to planets in a complementary part of the mass and semimajor axis parameter space to current Doppler surveys.  This work was supported by the National Science Foundation Graduate Research Fellowship awarded to S. R., by the NASA Origins of Solar Systems program grant NNG04GN30G and a National Science Foundation Career Grant to G. L., and by NASA grant NAG-5-13285 and NSF grant AST-0507424 to P. B.  \\appendix"
    },
    "0601/astro-ph0601460_arXiv.txt": {
        "abstract": "We test the accuracy of the ALI method, widely used in stellar atmosphere modelling, by solving exactly a standard radiative transfer problem in plane-parallel geometry. Some recommendations are given for a practical use of this method. ",
        "introduction": "We are interested in solving the radiative transfer equation in a homogeneous and isothermal plane-parallel atmosphere in local thermodynamical equilibrium, assuming isotropic and monochromatic scattering. If no radiation is incident on the boundary layers, the thermal source function is $\\epsilon B(T)$, where $\\epsilon$ is the photon destruction probability per scattering and $B(T)$ is the Planck function at temperature $T$. The source function $S$, normalized to the Planck function, satisfies the following integral equation \\begin{equation} \\label{eq_source} S(\\tau)=\\epsilon + (1-\\epsilon) (\\Lambda S)(\\tau)\\,, \\end{equation} where \\begin{equation} \\label{eq_lambda} (\\Lambda S)(\\tau)=\\frac{1}{2}\\int_0^{\\tau^*}E_1(\\vert \\tau-\\tau^{\\prime}\\vert)S(\\tau^{\\prime}) \\,\\mathrm{d}\\tau^{\\prime}\\, . \\end{equation} $E_1(\\tau)=\\int_0^1\\exp(-\\tau/\\mu)\\,\\mathrm{d}\\mu /\\mu$ is the first exponential integral function, and the optical depth variable $\\tau$ covers the range $[0, \\tau^*]$, $\\tau^*$ being the optical thickness of the atmosphere at a given frequency. Equation (\\ref{eq_source}) models the multiple scattering of photons in a continuum or in a spectral line with a rectangular profile.  Chevallier and Rutily (2003) have solved recently Eq. (\\ref{eq_source}) with an accuracy better than $10^{-10}$ for any value of $\\epsilon$, $\\tau^*$ and $\\tau$. Their solution can safely be used as a reference solution for testing the accuracy of a numerical code, e.g., the ALI code widely used in stellar atmosphere modelling. Our focus is to present some tests for this code, which we first describe in the next section. Some additional tests can be found in Chevallier et al. (2003). ",
        "conclusions": "In this particular problem (\\ref{eq_source}), our code was able to converge in all situations, its accuracy being bounded by $5\\times 10^{-3}$ and $10^{-2}$ when using the stopping criterion $R_c = 10^{-4}$. We intend to extend this work to more realistic problems (e.g., non isothermal atmospheres), replacing the Jacobi scheme by Gauss-Seidel or Successive Over-Relaxation schemes (see Trujillo Bueno \\& Fabiani Bendicho 1995)."
    },
    "0601/astro-ph0601111_arXiv.txt": {
        "abstract": "Two-dimensional hydrodynamic simulations are performed to investigate explosive nucleosynthesis in a collapsar using the model of MacFadyen and Woosley (1999). It is shown that $\\rm ^{56}Ni$ is not produced in the jet of the collapsar sufficiently to explain the observed amount of a hypernova when the duration of the explosion is $\\sim$10 sec, which is considered to be the typical timescale of explosion in the collapsar model. Even though a considerable amount of $\\rm ^{56}Ni$ is synthesized if all explosion energy is deposited initially, the opening angles of the jets become too wide to realize highly relativistic outflows and gamma-ray bursts in such a case. From these results, it is concluded that the origin of $\\rm ^{56}Ni$ in hypernovae associated with GRBs is not the explosive nucleosynthesis in the jet. We consider that the idea that the origin is the explosive nucleosynthesis in the accretion disk is more promising. We also show that the explosion becomes bi-polar naturally due to the effect of the deformed progenitor. This fact suggests that the $\\rm ^{56}Ni$ synthesized in the accretion disk and conveyed as outflows are blown along to the rotation axis, which will explain the line features of SN 1998bw and double peaked line features of SN 2003jd. Some fraction of the gamma-ray lines from $\\rm ^{56}Ni$ decays in the jet will appear without losing their energies because the jet becomes optically thin before a considerable amount of $\\rm ^{56}Ni$ decays as long as the jet is a relativistic flow, which may be observed as relativistically Lorentz boosted line profiles in future. We show that abundance of nuclei whose mass number $\\sim 40$ in the ejecta depends sensitively on the energy deposition rate, which is a result of incomplete silicon burning and alpha-rich freezeout. So it may be determined by observations of chemical composition in metal poor stars which model is the proper one as a model of a gamma-ray burst accompanied by a hypernova. ",
        "introduction": "There has been growing evidence linking long gamma-ray bursts (GRBs; in this study, we consider only long GRBs, so we call long GRBs as GRBs hereafter for simplicity) to the death of massive stars. The host galaxies of GRBs are star-forming galaxies and the position of GRBs appear to trace the blue light of young stars~\\citep{vreeswijk01,bloom02,gorosabel03}. Also, 'bumps' observed in some afterglows can be naturally explained as contribution of bright supernovae~\\citep{bloom99,reichart99,galama00,garnavich03}. Moreover, direct evidences of some GRBs accompanied by supernovae have been reported such as the association of GRB 980425 with SN 1998bw~\\citep{galama98,iwamoto98} and that of GRB 030329 with SN 2003dh~\\citep{hjorth03,price03,stanek03}.   It should be noted that these supernovae are categorized as a new type of supernovae with large kinetic energy ($\\sim 10^{52}$ ergs), nickel mass ($\\sim 0.5M_{\\odot}$), and luminosity~\\citep{iwamoto98,woosley99}, so these supernovae are sometimes called as hypernovae. Also, since GRBs are considered to be jet-like phenomena~\\citep{rhoads99,stanek99}, it is natural to consider the accompanying supernova to be jet-induced explosion~\\citep{macfadyen99, khokhlov99}. It is radioactive nuclei, $\\rm ^{56}Ni$ and its daughter nuclei, $\\rm ^{56}Co$, that brighten the supernova remnant and determine its bolometric luminosity. $\\rm ^{56}Ni$ is considered to be synthesized through explosive nucleosynthesis because its half-life is very short (5.9 days). So it is natural to consider explosive nucleosynthesis in jet-induced explosion to understand GRBs accompanied by hypernovae.   Nagataki et al. (1997) have done a numerical calculation of explosive nucleosynthesis taking account of effects of jet-induced explosion in the context of normal core-collapse supernova explosion whose explosion energy is set to be $10^{51}$ erg. They found that $\\rm ^{56}Ni$ is much produced in the jet region, which means that much explosive nucleosynthesis occurs around the jet region. Also, it was found that velocity distribution of iron becomes double peaked due to the asperical explosion and explosive nucleosynthesis~\\citep{nagataki98,nagataki00}. It was also found that the velocity distribution, which will be observed as a line profile, depends on the angle between our line of sight and rotation axis~\\cite{nagataki00}. Maeda et al. (2002) have done a numerical calculation of explosive nucleosynthesis taking account of effects of jet-induced explosion in the context of hypernovae whose explosion energy is set to be $10^{52}$ erg. They have shown that sufficient mass of $\\rm ^{56}Ni$ enough to explain the observation of hypernovae ($\\sim 0.5 M_{\\odot}$) can be synthesized around the jet region when the explosion energy is set to be $10^{52}$ erg. They have also calculated line profiles of [$\\rm Fe_{\\rm II}$] blend and of [$\\rm O_{\\rm I}$] from one-dimensional non-LTE nebular code~\\cite{mazzali01} and found that these line profiles depend on the angle between our line of sight and the direction of the jet. Recently, it was reported that an asymmetric hypernovae, SN 2003jd reveal double-peaked profiles in the nebular lines of neutral oxygen and magnesium~\\cite{mazzali05}.    However, here is one question: whether the difference between explosion energies of $10^{51}$ erg and $10^{52}$ is important or not? What does the different explosion energy mean? The answer is 'Yes'. It is impossible to overemphasize its importance because the scenario of explosion has to be dramatically changed to explain the energetic explosion energy of $10^{52}$ and to realize a GRB.   Nagataki et al. (1997) investigated explosive nucleosynthesis taking account of effects of jet-induced explosion in the context of normal core-collapse supernova explosion, because there is a possibility that normal core-collapse supernova becomes jet-like when effects of rotation are taken into account~\\citep{yamada94,shimizu94,kotake03}. In this scenario, the typical timescale of core-collapse supernovae is as short as $\\sim 500$ ms~\\cite{wilson85}. So the surrounding layers of iron core does not collapse so much. As a result, the progenitor outside the iron core can not be deformed due to the collapse. This fact supports the treatment of using a spherical progenitor when explosive nucleosynthesis is investigated. Note that the central iron core collapses to a neutron star and it is enough to calculate explosive nucleosynthesis in a spherical outer layer such as Si-, O-, and He-rich layers in the case of normal core-collapse supernovae~\\citep{nagataki97,nagataki98,nagataki98b,nagataki00}. Also, to initiate the explosion, the explosion energy is deposited around the Si-rich layer as an initial condition in their works. This is justified because the timescale of explosion is very short.    On the other hand, the central engine of GRBs accompanied by hypernovae is not known well. But it is generally considered that normal core-collapse supernovae can not cause an energetic explosion of the order of $10^{52}$ erg. So another scenario has to be considered to explain the system of GRBs associated with hypernovae. One of the most promising scenario is the collapsar scenario~\\cite{woosley93}. In the collapsar scenario, a black hole is formed as a result of gravitational collapse. Also, rotation of the progenitor plays an essential role. Due to the rotation, an accretion disk is formed around the equatorial plane. On the other hand, the matter around the rotation axis falls into the black hole. It was pointed out that the jet-induced explosion along to the rotation axis occurs due to the heating through neutrino anti-neutrino pair annihilation that are emitted from the accretion disk. MacFadyen and Woosley (1999) demonstrated the numerical simulations of the collapsar, showing that the jet is launched $\\sim 7$ sec after the gravitational collapse and the duration of the jet is about 10 sec, which is comparable to the typical observed duration of GRBs~\\citep{mazet81,kouveliotou93,lamb93}. This timescale is much longer than the typical timescale of normal core-collapse supernovae. As a result, the progenitor becomes deformed even at the Si-rich and O-rich layer in the collapsar model~\\cite{macfadyen99}. In particular, the density around the rotation axis becomes low because considerable amount of the matter falls into the black hole, which is a good environment to produce a fire ball~\\citep{woltjer66,rees67}.   Maeda et al. (2002) investigated explosive nucleosynthesis taking account of effects of jet-induced explosion in the context of hypernovae whose energy is $10^{52}$ erg using the spherical progenitor model and depositing explosion energy at the inner most region initially. However, this treatment seems to be incompatible with the collapsar scenario. The importance of the duration of explosion, $\\sim 10$ sec, is investigated in some papers~\\citep{nagataki03,maeda03} and it was concluded that the abundance of $\\rm ^{56}Ni$ synthesized during explosion depends sensitively on the duration of explosion (i.e. energy deposition rate) and $\\rm ^{56}Ni$ is not produced sufficiently to explain the observed amount of $\\sim 0.5 M_{\\odot}$ when the timescale of explosion becomes as long as 10 sec. In fact, MacFadyen and Woosley (1999) discussed that enough $\\rm ^{56}Ni$ should not be synthesized in the jet in the collapsar model. Rather, they pointed out the possibility that a substantial amount of $\\rm ^{56}Ni$ is produced in the accretion disk and a part of it is conveyed outwards by the viscosity-driven wind~\\citep{macfadyen99,pruet02}. There is another question. Does all of $\\rm ^{56}Ni$ produced in the jet of the collapsar model brighten the supernova remnant? If the jet becomes optically thin before $\\rm ^{56}Ni$ decays into $\\rm ^{56}Co$ and $\\rm ^{56}Co$ decays into $\\rm ^{56} Fe$, these nuclei should result in emitting gamma-rays rather than brightening the supernova remnant.   Let us summarize our motivation. We want to understand how collapsars produce a GRB jet and how collapsars eject sufficiently enough $\\rm ^{56}Ni$ to explain the luminosity of hypernovae. We want to seek the self-consistent theory of GRB/Hypernova connection. As a first step, we want to consider in this work the consistency between the collapsar model of MacFadyen and Woosley (1999) and explosive nucleosynthesis in a hypernova jet. Can a hypernova jet cause a GRB jet and a sufficiently enough explosive nucleosynthesis to explain the luminosity of hypernova? Rather, should we consider the GRB jet is different from the hypernova jet? Moreover, should we consider $\\rm ^{56}Ni$ comes from the explosive nucleosynthesis in the hypernova jet? Rather, should we consider $\\rm ^{56}Ni$ comes from a different site? We want to know the answer. This is our motivation of this work.  Due to the motivation mentioned above. We investigate explosive nucleosynthesis in the context of the collapsar model. We use the collapsar model of MacFadyen and Woosley (1999) in which effects of rotation is included. As a result, the progenitor becomes deformed significantly as mentioned above. We show that $\\rm ^{56}Ni$ is not produced sufficiently to explain the observed amount when the duration of the explosion is $\\sim$10 sec, which is consistent with the previous works~\\citep{nagataki03,maeda03}. A fine tuning is required to explain the amount of $\\rm ^{56}Ni$ by the explosive nucleosynthesis in the jet. This result bring us to the conclusion that the origin of $\\rm ^{56}Ni$ in hypernovae associated with GRBs is not the explosive nucleosynthesis in the jet but the one in the accretion disk. We also show that the explosion becomes bi-polar naturally due to the effect of deformed progenitor. This fact suggests that the $\\rm ^{56}Ni$ synthesized in the accretion disk and conveyed as outflows are blown along to the rotation axis, which can explain the line features of SN 1998bw and double peaked line features of SN 2003jd~\\cite{mazzali05}. Also we predict that some fraction of gamma-ray lines from $\\rm ^{56}Ni$ decays in the jet may show relativistically Lorentz boosted line profiles, which might be observed in future. ",
        "conclusions": "We have performed 2-dimensional hydrodynamic simulations to investigate explosive nucleosynthesis in a collapsar using the model of MacFadyen and Woosley (1999). We have shown that $\\rm ^{56}Ni$ is not produced in the jet sufficiently to explain the observed amount of a hypernova such as SN 1998bw when the duration of the explosion is $\\sim$10 sec (the standard model, E51). Even though a considerable amount of $\\rm ^{56}Ni$ is synthesized if all explosion energy is deposited initially (the extreme models, E52 and E52S), the opening angles of the jets become too wide to realize highly relativistic outflows and a GRB in such a case. From these results, it is concluded that the origin of $\\rm ^{56}Ni$ in hypernovae associated with GRBs is not the explosive nucleosynthesis in the jet. We consider that the idea that the origin of $\\rm ^{56}Ni$ in hypernovae is the explosive nucleosynthesis in the accretion disk is more promising. We have also shown that the explosion becomes bi-polar naturally due to the effect of the deformed progenitor. This fact suggests that the $\\rm ^{56}Ni$ synthesized in the accretion disk and conveyed as outflows are blown along to the rotation axis, which will explain the line features of SN 1998bw and double peaked line features of SN 2003jd, and will help the idea of the accretion disk mentioned above.    Also, some predictions are presented in this study. Some fraction of the gamma-ray lines from $\\rm ^{56}Ni$ decays in the jet will appear without losing their energies because the jet becomes optically thin before a considerable amount of $\\rm ^{56}Ni$ decays as long as the jet is a relativistic flow. So it has been predicted that some fraction of $\\rm ^{56}Ni$ synthesized in the jet may show relativistically Lorentz boosted line profiles. That is, highly blue shifted (or red shifted) broad line features might be observed in future. It has been also shown that there is an enhancement of nuclei whose mass number $\\sim 40$ in Models E52 and E52S compared with Model E51 as a result of incomplete silicon burning and alpha-rich freezeout. So it may be determined which model is the proper one as a model of a gamma-ray burst accompanied by a hypernova by observations of chemical composition in metal poor stars.   Of course there is still uncertainty of observations and theories on the formation of GRBs accompanied by hypernovae. Further investigation should still be required to understand the central engine of GRBs and origin of $\\rm ^{56}Ni$ in hypernovae."
    },
    "0601/astro-ph0601327_arXiv.txt": {
        "abstract": "The color-magnitude relation has been determined for the RDCS~J0910+5422 cluster of galaxies at redshift z = 1.106.  Cluster members were selected from HST Advanced Camera for Surveys (ACS) images, combined with ground--based near--IR imaging and optical spectroscopy.  Postman et al. (2005) morphological classifications were used to identify the early-type galaxies.  The observed early--type color--magnitude relation (CMR) in $(i_{775} - z_{850})$ versus $z_{850}$ shows an intrinsic scatter in color of $0.060 \\pm 0.009$~mag, within $1\\arcmin$ from the cluster X--ray emission center. Both the ellipticals and the S0s show small scatter about the CMR of $0.042 \\pm 0.010$~mag and $0.044 \\pm 0.020$~mag, respectively.  From the scatter about the CMR, a mean luminosity--weighted age $\\overline t > 3.3$~Gyr ($z_f > $~3) is derived for the elliptical galaxies, assuming a simple stellar population modeling (single burst solar metallicity). This is consistent with a previous study of the cluster RDCS\\,1252.9-292 at $z{=}1.24$ (Blakeslee et al.).   Strikingly, the S0 galaxies in RDCS~J0910+5422 are systematically bluer in $(i_{775}{-}z_{850})$ by $0.07 \\pm 0.02$~mag, with respect to the ellipticals. The blue S0s are predominantly elongated in shape; the distribution of their ellipticities is inconsistent with a population of axisymmetric disk galaxies viewed at random orientations, suggesting either that they are intrinsically prolate or there is some orientation bias in the S0 classification. The ellipticity distribution as a function of color indicates that the face-on S0s in this particular cluster have likely been classified as elliptical. Thus, if anything, the offset in color between the elliptical and S0 populations may be even more significant.  The color offset between S0 and E corresponds to an age difference of $\\approx$~1~Gyr, for a single-burst solar metallicity model. Alternatively, it could be the result of a different star formation history; a solar metallicity model with an exponential decay in star formation will reproduce the offset for an age of 3.5 Gyr, i.e. the S0s have evolved gradually from star forming progenitors. The color offset could also be reproduced by a factor of $\\sim\\,$2 difference in metallicity, but the two populations would each need to have very small scatter in metallicity to reproduce the small scatter in color.  The early--type population in this cluster appears to be still forming. The blue early-type disk galaxies in RDCS~J0910+5422 likely represent the direct progenitors of the more evolved S0s that follow the same red sequence as ellipticals in other clusters.  Thirteen red galaxy pairs are observed and the galaxies associated in pairs constitute $\\sim$40\\% of the CMR galaxies in this cluster. This finding is consistent with the conclusions of van Dokkum and Tran et al. that most of the early--type galaxies grew from passive red mergers. ",
        "introduction": "The Advanced Camera for Surveys (ACS; Ford et al. 2002), by virtue of its high spatial resolution and sensitivity, allows us to study galaxy clusters in great detail up to redshifts of unity and beyond. At these redshifts, galaxy clusters are still assembling and galaxies are evolving towards the populations that we observe today.    Recent results from our ACS Intermediate Redshift Cluster Survey (Blakeslee et al. 2003a; Lidman et al. 2004; Demarco et al. 2005; Goto et al. 2005; Holden et al. 2005a; Holden et al. 2005b; Homeier et al. 2005; Postman et al. 2005) have shown that galaxy clusters at redshift around unity show many similarities with local clusters, in terms of galaxy populations and their distribution, but also significant differences in galaxy morphology, ellipticity, and mass--luminosity ratios. The strongest evolution observed in the early--type population is a deficit of a S0 population in this sample when compared to lower redshift samples (Postman et al. 2005). This would give evidence that the formation of the S0 population is still under way in clusters at redshift unity.  One of the most striking similarities is that the tight relation between early-type galaxy colors and luminosities that applies locally (the color--magnitude relation; CMR) is already in place at redshifts as high as $z \\sim 1.3$ (e.g. Stanford et al.\\ 1997; Mullis et al.\\ 2005). The CMR in local samples of galaxy clusters presents universal properties, in terms of scatter and zero point (Bower et al. 1992; van Dokkum et al. 1998, Hogg et al. 2004; L\\'opez--Cruz et al. 2004; Bell et al. 2004; Bernardi et al. 2005; McIntosh et al. 2005) that evolve back in time in agreement with passively evolving models (Ellis et al.\\ 1997; Stanford, Eisenhardt, \\& Dickinson 1998; van Dokkum et al. 2000, 2001; Blakeslee et al. 2003a; Holden et al. 2004; De Lucia et al. 2004; Blakeslee et al. 2005). ACS enables accurate measurement of the scatter around the CMR, with enough precision to seriously constrain galaxy formation age, which is impossible to obtain from ground--based data (see for example Holden et al. 2004). The measurement of the CMR scatter of the first cluster in our ACS cluster survey, RXJ1252.9-292, permitted us to constrain the mean luminosity--weighted age for the ellipticals to be $> 2.6$~Gyr ($z > 2.7$) (Blakeslee et al. 2003a), based on simple modeling.  In this paper, we extend the results obtained in Blakeslee et al. (2003a) to RXJ~0910+5422.   RXJ~0910+5422 is part of the ACS cluster survey (guaranteed time observation, GTO, program \\#9919), that includes eight clusters in the redshift range at $0.8 < z < 1.3$, selected in the X--ray, optical and near--IR (Ford et al.\\ 2004). RXJ~0910+5422 was selected from the ROSAT Deep Cluster Survey (Rosati et al. 1998) and confirmed with near-IR and spectroscopic observations by Stanford et al. (2002). Extensive followup spectroscopy at the Keck Observatory has been carried out in a magnitude limited sample reaching $K_s = 20.0$~mag in the central $3$~arcmin (Stanford et al.\\ 2005; in preparation).  The mean redshift of the cluster was measured to be $z=1.106$ (Stanford et al. 2002).  In this paper, we combine ACS imaging with ground--based spectroscopy and near--IR imaging to constrain galaxy ages and formation histories from the study of their color--magnitude relation. We discuss the properties of the elliptical (E) and lenticular (S0) populations separately in the light of simple galaxy formation scenarios. ",
        "conclusions": "In this paper we have studied the color--magnitude relations of galaxies in RXJ\\,0910+5422 to constrain their ages and formation histories. Our results show that the  color--magnitude relation for the elliptical galaxies is consistent both in slope and scatter with that of RXJ\\,1252.9-292 (Blakeslee et al. 2003a, Lidman et al. 2004) and recent results from Holden et al. (2004) and De Lucia et al. (2004), confirming that elliptical galaxies in galaxy clusters show a universal color--magnitude relation consistent with an old passively evolving population even at $z \\sim 1$. From the color--magnitude relation of the ellipticals, we derive a mean luminosity--weighted age $\\overline t > 3.3$~Gyr ($z_f > $~3).  We find that the S0s in RXJ~0910+5422 define a color--magnitude sequence with a scatter similar to that found for the ellipticals, but shifted bluer by  $0.07 \\pm 0.02$~mag.   This is peculiar with respect to previous cluster studies, which more typically found that the S0s followed the same CMR as the ellipticals, but with somewhat larger scatter (Bower et al. 1992; Ellis et al.\\ 1997; Stanford et al.\\ 1997; L\\'opez--Cruz et al. 2004; Stanford, Eisenhardt, \\& Dickinson 1998; van Dokkum et al. 2000, 2001; Blakeslee et al. 2003a; Holden et al. 2004). Only one earlier study, van Dokkum et al. (1998) found a significantly bluer S0 population. We examine this population of blue S0s in some detail, noting that there is a strong predominance of flattened systems with axial ratios $\\frac{b}{a} > 0.7$, and conclude that the face-on members of the population have likely been classified as ellipticals.  If so, the color offset between the two classes would become even more significant, and the true CMR scatter for the ellipticals would be slightly lower than we have estimated. This peculiarity is not observed in other clusters of our ACS Intermediate Redshift Cluster Survey, and its amplitude is smaller than the uncertainties adopted in Postman et al. (2005).  If the observed color difference between the ellipticals and S0s is mainly due to metallicity at the same age, this would imply that the redder ellipticals were able to retain more metals than the S0s, i.e., they are more massive. However, current data suggest that the CMR is mainly the result of a metallicity--mass (i.e. metallicity--magnitude) relation (e.g., Kodama \\& Arimoto 1997; Kauffman \\& Charlot 1998 Vazdekis et al. 2001; Bernardi et al. 2005). This implies that we do not expect large metallicity variations at a given magnitude. If, instead, the offset is mainly due to age, then the implied age difference would be $\\sim$1~Gyr for single-burst solar metallicity BC03 models.  It could also result from different star formation histories, with the S0s experiencing a more extended period of star formation.  A model with solar metallicity and with an exponential decay of the star formation reproduces the offset at a galaxy mean age of $\\approx$~3.5~Gyr.  The blue S0s may comprise a group infalling from the field onto a more evolved red cluster population, or they may be a transitional cluster population not yet evolved all the way onto the elliptical red sequence (van Dokkum \\& Franx 2001). Assuming passive evolution, they will reach this red sequence after about 1~Gyr. High fractions of faint blue late type galaxies were observed in substructures infalling in a main cluster (e.g. Abraham et al. 1996; Tran et al. 2005a), and proposed as the progenitors of faint S0s in clusters. The view of this cluster as a structure still in formation is supported by X--ray observations of the cluster temperature structure (Stanford et al. 2002), the lack of a cD galaxy, and its filamentary structure that suggests merging of substructures. However, we derive a small cluster velocity dispersion, unusual for merging substructures. Moreover, the blue S0s in this sample span the same luminosity range of the bright ellipticals, are distributed towards the center of the cluster, and some of the faintest ones are physically associated with brighter ellipticals belonging to the central filamentary structure. These elements would argue against the bluer S0s being a young group merging with an existing red cluster population, and  support the hypothesis that we are observing a transitional blue S0 population in a cluster core that is still evolving onto the elliptical red sequence. This result is also consistent with the deficit of S0s observed in our ACS cluster sample, when compared to lower redshift samples, that implies that the S0 population is  not yet in place but still forming in clusters at redshifts around unity (Postman et al. 2005). Interestingly, we also observe in this cluster potential progenitors for bright S0 galaxies: four bright spirals (spectroscopically confirmed cluster members) with $z_{850}$ brighter than 22.5~mag and $(i_{775}-z_{850})$ between 0.5 and 1.3~mag.  Red galaxy pairs are also observed. A triplet and three red galaxy pairs have projected distances less than $10 h^{-1}_{70}$~kpc, and of those, the triplet and one pair show zero relative velocity. This would be the evidence of red galaxy mergers at z$\\sim$~1. van Dokkum (2005) and Tran et al. (2005b) have observed mergers of red galaxies in a nearby elliptical sample and in  MS 1054-03 at z=0.83, respectively. They suggested a scenario in which most of the early--type galaxies were formed from passive red galaxy--galaxy mergers, called {\\it dry} mergers, because they involve gas--poor early--type galaxies.  Future papers will analyze the ages and masses of the cluster members using our optical spectroscopy along with newly obtained Spitzer IRAC imaging. A larger sample would be needed to draw firmer conclusions about the formation of S0s."
    },
    "0601/astro-ph0601057_arXiv.txt": {
        "abstract": "We investigate properties of the interstellar medium (ISM) in galaxies hosting long-duration $\\gamma$-ray bursts (GRBs) from an analysis of atomic species (Mg$^0$, Fe$^0$) and excited fine-structure levels of ions (e.g.\\ Si$^+$).  Our analysis is guided primarily by echelle observations of GRB~050730 and GRB~051111. These sightlines exhibit fine-structure transitions of O$^0$, Si$^+$, and Fe$^+$ gas that have not yet been detected in intervening quasar absorption line systems.  Our results indicate that the gas with large \\ion{Mg}{1} equivalent width (e.g.\\ GRB~051111) must occur at distances $\\gtrsim 50$\\,pc from GRB afterglows to avoid photoionization.  We examine the mechanisms for fine-structure excitation and find two processes can contribute: (1) indirect UV pumping by the GRB afterglow provided a far-UV intensity in excess of 10$^6$ times the Galactic radiation field; and (2) collisional excitation in gas with electron density $n_e > 10^4 \\cm{-3}$.  The observed abundances of excited ions are well explained by UV pumping with the gas at $r\\,\\sim$\\,a few hundred pc from the afterglow for GRB~051111 and $r<100$ pc for GRB~050730, without invoking extreme gas density and temperature in the ISM.  We show that UV pumping alone provides a simple explanation for all reported detections of excited ions in GRB afterglow spectra. The presence of strong fine-structure transitions therefore may offer little constraint for the gas density or temperature. We discuss additional implications of UV pumping including its impact on chemical abundance measurements, new prospects for observing line-strength variability, and future prospects for studying the gas density and temperature.  Finally, we list a series of criteria that can distinguish between the mechanisms of UV pumping and collisional excitation. ",
        "introduction": "Optical afterglows of long-duration $\\gamma$-ray bursts (GRBs) are almost certainly signposts of extreme star-forming regions at high redshift, because these bursts originate in the catastrophic death of massive stars \\citep[e.g.][]{woo93,pac98a,bkp+02,smg+03}.  A detailed study of metal absorption features identified in the circumburst environment can, in principle, yield strong constraints on the progenitor \\citep[e.g.][]{pl98a,rgs+05,vlg05} and offer important insights for understanding mass-loss and chemical feedback during the final evolution stages of massive stars \\citep{wij01,vink05}. The benefit of GRBs, which presumably are embedded in environments of Wolf-Rayet stars, is that their optical afterglows provide extremely bright lighthouses at moderate to high redshift, allowing redshifted UV lines to be (albeit transiently) observable with high-resolution spectroscopy.  Past studies based on low-resolution afterglow spectra have been limited to a few diagnostic measurements such as the \\ion{H}{1} column density, the gas metallicity, and the dust-to-gas ratio \\citep[e.g.][]{sff03,vel+04}.  In contrast, high-resolution spectroscopy of GRB afterglows has uncovered detailed kinematic signatures, population ratios of excited ions, and chemical compositions of the interstellar medium (ISM) and, perhaps, the circumstellar medium (CSM) of the progenitor star of the GRB \\citep{vel+04,cpb+05}.  These quantities, in principle, permit direct comparisons with simulations describing the evolution history of GRB progenitors \\citep[e.g.][]{vlg05}.  An emerging feature of GRB progenitor environments is the presence of strong fine-structure transitions from excited states of C$^+$, Si$^+$, O$^0$, and Fe$^+$ \\citep{cpb+05}.  Together with their resonance transitions, these fine-structure lines can in principle provide robust constraints on the temperature and density of the gas, as well as the ambient radiation field, independent of the ionization fraction, relative abundances, or metallicity of the gas \\citep{bw68}. In particular, detections of \\ion{Fe}{2} fine structure transitions have only been reported in rare places such as broad absorption-line (BAL) quasars \\citep{has+02}, $\\eta$ Carinae \\citep{gull05,nielsen05} and the circumstellar disk of $\\beta$ Pictoris \\citep{lvf88}.  The presence of strong \\ion{Fe}{2} fine-structure lines therefore suggests extreme gas density and temperature in the GRB progenitor environment. In the case of circumburst medium, however, the greatly enhanced radiation field due to the GRB afterglow may also contribute significantly to the formation of these excited ions.  \\begin{figure*} \\begin{center} \\includegraphics[width=6.5in]{f1.eps} \\caption{Velocity profiles for transitions arising from gas in the host galaxy of GRB~051111.  The vertical dashed line at $v=0$ corresponds to $z=1.5495$ and the dotted lines indicate blends with coincident transitions. } \\label{fig:051111} \\end{center} \\end{figure*}  In this paper we present a comprehensive study of the physical properties of the ISM in GRB host galaxies, using absorption lines identified in the echelle spectra of two GRB afterglows, GRB\\,051111 and GRB\\,050730 as case studies.  We first consider the enhanced radiation field in the circumburst environment due to the progenitor star as well as the optical afterglow, and constrain the distance of the observed neutral gas cloud based on the presence/absence of various atomic lines such as Mg\\,I and Fe\\,I.  The dominant species of these elements is in the singly ionized state, because their first ionization potential is less than 1\\,Ryd.  The presence of strong atomic lines therefore places a lower limit to the distance of the cloud from the afterglow.  Next, we investigate the excitation mechanism of the absorbing gas, again taking into account the presence of an intensified radiation field known from afterglow light-curve observations.  We compare the predicted population ratios with observations of different ions under photon pumping and collisional excitation scenarios.  Although the analysis is based on echelle data obtained for two GRB hosts, we note that the results are applicable to the majority of long-duration GRB's discovered to date.  Finally, we discuss additional framework for studying the ISM gas of future events.  The paper is organized as follows.  The observations and analysis of echelle spectra for GRB~051111 and GRB~050730 are briefly summarized in $\\S$~\\ref{sec:obs}.  In $\\S$~\\ref{sec:rad} we describe the radiation field produced by GRB afterglows.  In $\\S$~\\ref{sec:neut} we analyze the non-dominant atomic transitions observed in optical afterglow spectra, while $\\S$~\\ref{sec:fine} investigates the mechanisms for populating fine-structure states in the dominant ions. In $\\S$~\\ref{sec:disting} we consider tests to distinguish between UV pumping and collisional excitation as the principal mechanism for population of fine-structure levels.  Finally, $\\S$~\\ref{sec:discuss} presents a discussion of the results for the specific cases of GRB~051111 and GRB~050730 and, also, generic results for observations of long-duration GRBs. Throughout the paper we assume a cosmology with $\\Omega_\\Lambda = 0.7$, $\\Omega_m = 0.3$ and $H_0 = 70 {\\rm km \\, s^{-1} \\, Mpc^{-1}}$.   \\begin{figure}[ht] \\begin{center} \\includegraphics[width=3.5in]{f2.eps} \\caption{Velocity profiles for \\ion{Fe}{2} transitions arising from gas in the host galaxy of GRB~050730.  The vertical dashed line at $v=0$ corresponds to $z=3.96855$ and the dotted lines indicate blends with coincident transitions. } \\label{fig:050730} \\end{center} \\end{figure} ",
        "conclusions": "\\label{sec:discuss}   \\subsection{Implications for GRB~051111 and GRB~050730}  Despite different spectral coverage in the rest frame of GRB\\,051111 and GRB\\,050730, the echelle spectra of these afterglows exhibit ground-state Si$^+$ and Fe$^+$ features from all of their excited levels.  In the case of GRB\\,050730, excited O$^0$ is also present. But due to measurement uncertainties, their population ratios may be explained by Boltzmann distribution with $T_{Ex} = 2600 - 6000$K (Figures 10 \\& 11), as well as by UV pumping provided $G/G_0 \\gtrsim 10^{6.5}$.  We have investigated variations in the column densities and set {\\it upper limits} to the line-strength variability of $< 10\\%$ per 1800s for GRB~051111 and $< 30\\%$ per 1800s for GRB~050730. It appears difficult to unambiguously rule out either UV pumping or collisional excitation as the dominant excitation mechanism. The only strong argument against collisional excitation is the absence of Fe$^0$ and C$^0$ atomic gas for GRB~051111 and GRB~050730, respectively.  Under collisional excitation, the excited ions occur irrespective of the presence of the afterglow and the gas is required to have $n_e > 10^{5} \\cm{-3}$ and $T \\sim 2000-6000$ K to populate the \\ion{Fe}{2} excited levels.  If the excited gas arises from a layer of high density ($n_e > 10^5 \\cm{-3}$), then based on the absence of \\ion{Fe}{1} we can place an upper limit on the ratio of the electron density to far-UV intensity $n_e\\times (G/G_0)^{-1}$ (Equation~\\ref{eqn:nefe}).  Adopting the best-fit Boltzmann distribution, we calculate the total Fe$^+$ column density $N_{\\rm tot}{Fe^+} = 10^{14.4} \\cm{-2}$, including all levels.  Adopting a conservative upper limit to the atomic state $\\N{Fe^0} < 10^{12} \\cm{-2}$ (Table~\\ref{tab:051111}) which accounts for continuum and statistical uncertainty, we therefore find from Equation~\\ref{eqn:nefe} that $n_e\\times (G/G_0)^{-1} < 1$ for $T < 10000$K.  For $n_e > 10^5 \\cm{-3}$, we find that $G/G_0 > 10^5$.  Therefore, we conclude that we would only observe significant populations of the excited \\ion{Fe}{2} levels for gas in the presence of a strong far-UV radiation field.  To satisfy $G/G_0 > 10^5$, it requires the gas to lie within $\\approx 0.1$pc of a $T=40000$K O star with $R = 10 R_\\odot$.  While the GRB progenitor star is expected to reside in a star-forming region and therefore near other massive stars, we consider it highly unrealistic that the gas is at $>100$ pc from the GRB but also within $0.1$ pc of an O star.  A similar conclusion can be drawn for GRB~050730 using \\ion{C}{1}.  We infer a C$^+$ column density by scaling $\\N{Fe+} = 10^{14.8} \\cm{-2}$ from the excited region with solar relative abundances: $\\N{C^+} = 10^{15.8} \\cm{-2}$.  The observed upper limit $\\N{C^0}/\\N{C^+} < 10^{-2.3}$ requires $G/G_0 > 5 \\sci{4}$, again assuming a conservative limit to the temperature ($T < 10000$K).  For GRB~050730, we have no significant constraint on the distance of the gas from the afterglow. But, if collisional excitation dominates indirect UV pumping (and that $n_e < 10^{9} \\cm{-3}$), then this gas must also be experiencing a strong local source of far-UV radiation.  We conclude based on the null detections of atomic gas that {\\it indirect UV pumping by GRB afterglow radiation is the dominant excitation process.}  We also note that there are additional challenges to a collisional excitation scenario.  First, the gas would require a very large heating source to maintain the inferred temperatures for collisional excitation.  For the implied physical conditions (e.g.\\ $T=2600$ K and $n_e > 10^5 \\cm{-3}$), [\\ion{Fe}{2}] cooling dominates.  Using the coefficients given by \\cite{hollenbach89}, we calculate a rate of $n \\lfeII = 4 \\sci{-23} {\\rm ergs \\, s^{-1} \\, H^{-1}}$ assuming 1/10 solar metallicity and that $90\\%$ of the iron is locked into dust grains.  This implies a cooling rate $t_{cool} \\sim kT / (n \\lfeII) < 500 {\\rm yr^{-1}}$.  Including the cooling by meta-stable states of \\ion{O}{1} would imply an even shorter cooling time.  Second, gas with $T \\approx 3000$ K is generally not stable to thermal perturbation. In the case considered here, the [\\ion{Fe}{2}] cooling curve is relatively insensitive to temperature variations but very sensitive to density changes.  For example, an increase in $n_H$ increases $n \\lfeII$ and drives the gas to higher density and lower temperature until $T \\ll 1000$\\,K.  Third, the implied pressure $nT$ is extremely high.  One would require a strong mechanical process (e.g.\\ shock) to confine the gas and it would be surprising to therefore find the quiescent kinematic characteristics that are observed (Figures~\\ref{fig:051111},\\ref{fig:050730}).   Adopting UV pumping as the the excitation mechanism, we can constrain the distance of the gas from the GRB based on known afterglow radiation field.  For GRB~051111, the gas must arise at $r \\approx 200$\\,pc.  This distance is large enough to prevent the ionization of Mg$^0$ but also close enough that the \\ion{Fe}{2} and \\ion{Si}{2} levels would be UV pumped and also show the relative populations that are observed. In the case of GRB~050730, where the excited \\ion{Fe}{2} levels are populated according to their degeneracies $g_J$, the gas must lie within $\\approx 100$\\,pc of the afterglow (Table~\\ref{tab:rlim}).  The ionizing column of photons at 1\\,pc is $6\\sci{22} \\cm{-2}$.  The gas is therefore not likely at $< 3$\\,pc, where the far-UV radiation field from the GRB afterglow will photoionize all of the C$^+$, Si$^+$, O$^0$, and Fe$^+$ gas.  \\subsection{Implications of Fine-Structure Lines in Other GRB Afterglows}  The detection of \\ion{Si}{2}* transitions is nearly ubiquitous in GRB afterglow spectra \\citep{vel+04,cpb+05,bpck+05}. This stands in stark contrast with intervening quasar absorption line systems where only \\ion{C}{1}* and \\ion{C}{2}* fine-structure lines have been observed \\citep[e.g.][]{srianand00,wolfe03a,howk05}. To date, the observation of \\ion{Si}{2}* absorption for GRB have been interpreted as the result of collisional excitation and therefore evidence for large volume densities in the gas. We have demonstrated in this paper, however, that UV pumping is the most likely excitation mechanism if the gas is located within a distance of $\\approx 100$\\,pc.  This conclusion holds for all of the GRBs with reported \\ion{Si}{2}* detections (Table~\\ref{tab:rlim}) except for GRB~050408 \\citep{foley06}. And, for sightlines exhibiting \\ion{Fe}{2}* absorption, we contend UV pumping is likely the only mechanism.  In summary, we conclude that in the presence of an intensified UV radiation field from the afterglow {\\it indirect UV pumping alone offers a simple explanation for all the excited fine-structure transitions observed in GRB host environment to date without invoking extreme ISM properties}, such as high gas density.  If we take the gas density allowed by the observed relative abundances of atomic species, $n_{\\rm H}=100-1000\\,cm^{-3}$ (\\S\\ 4.2), the observed large \\nhi\\ \\citep[e.g.][]{cpb+05} implies a cloud thickness of $3-30$ pc which is typical of giant molecular clouds (Turner 1988).  In lieu of diagnostics from the fine-structure populations, one must rely on comparisons of the atomic states with higher ionization levels to constrain the gas density.  We contend that the principal quantity inferred from the fine-structure lines is the distance to the GRB afterglow.  At present, there is no obvious diagnostic for the gas temperature.  \\subsection{Implications for Gas Within 100\\,pc of GRB Afterglows}  Our analysis reveals a number of implications for the study of gas located within 100\\,pc of the GRB afterglow.  The first is that this gas may be identified by the absence of atomic absorption, in particular \\ion{Mg}{1} gas.  Although the sightline is likely to penetrate \\ion{Mg}{1} gas at distances far from the GRB (e.g.\\ within the galactic halo), differences in the line-of-sight velocities should distinguish these clouds.  Therefore, a study of the kinematic differences between \\ion{Mg}{1} and \\ion{Mg}{2} profiles in GRB spectra is warranted.  In the absence of atomic gas (e.g.\\ \\ion{Mg}{1} absorption) that coincides kinematically with the low-ion species, the gas may have small distance from the GRB afterglow and therefore may be circumstellar material from the progenitor.  Unfortunately, the afterglow radiation washes out the most readily available diagnostics for studying the density and temperature of this gas.  We note, however, that self-shielding of UV photons matching various excitation energies may be significant, especially for the majority of other GRB sightlines where the fine-structure column density is signficantly lower than the ground-state \\citep[e.g.][]{vel+04,savaglio04}.  We expect this may also be the case for GRB~051111 and GRB~050730.  It is therefore possible that a significant portion of the gas is not excited by the afterglow and we could infer physical conditions for the bulk of the medium.  Another implication is that one must consider the fine-structure levels in the chemical abundance analysis.  This will be especially important for sightlines with low total metal-line column density. The correction to $\\N{Fe^+}$ in GRB~050730, for example, is 33$\\%$ and we estimate the correction to $\\N{Si^+}$ is greater than 50$\\%$. If one cannot account for the excited-states in the analysis (e.g.\\ because the lines are too saturated/blended for accurate measurement), then analysis of S$^+$ is preferred because it does not have fine-structure levels near the ground-state.  Finally, we emphasize again that to observe gas at 10\\,pc one requires a large pre-existing column density of neutral or partially ionized gas prior to the GRB event.  This includes highly ionized species like \\ion{C}{4}.  Consider, as an example, the relatively spectacular \\ion{C}{4} gas observed toward GRB~021004 which several authors have interpreted to arise from gas associated with the GRB progenitor \\citep{sgh+03,Mir03,vlg05,fdl+05}.  At $t_{obs}=1$ hr after the GRB, we estimate an \\ion{H}{1} ionizing column density of $3 \\sci{21} \\, {\\rm photons} \\, \\cm{-2}$ at 10\\,pc assuming the afterglow properties given in Table~\\ref{tab:rlim}.  This column significantly exceeds the \\nhi\\ value observed for the gas toward GRB~021004, therefore we assume that all of the low-ion gas within 10\\,pc was either ionized prior to the GRB or prior to the observations.  This helps explain the absence of strong low-ion absorption\\footnote{The observation of \\ion{Al}{2} absorption is an interesting exception.}.  More importantly, the column density of C$^{+3}$ ionizing photons with energy $h\\nu = 64.5$ to 70eV at 10\\,pc is $\\approx 10^{20} \\cm{-2}$. If this flux is unattenuated by gas within 10\\,pc then it would significantly ionize all of the C$^{+3}$ ions within a few tens of pc from the GRB.  At these energies, the principal source of opacity is He$^+$, but we estimate that this gas will also have been burned away by the GRB and its progenitor.  We contend, therefore, that the majority of C$^{+3}$ gas observed along the GRB~021004 sightline is either associated with the larger star forming region and/or halo gas from the GRB host galaxy and/or its galactic neighbors."
    },
    "0601/astro-ph0601507_arXiv.txt": {
        "abstract": "s{ 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$. Given that the GRBs range of redshifts is rather high, it turns out that they are not very sensitive to the dark energy component, being however, fairly good estimators of the amount of dark matter in the Universe. } ",
        "introduction": "Recently, there has been a flurry of activity about the prospect use of 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 a redshift 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}$) and variability ($V$). The isotropic equivalent luminosity is the inferred luminosity (energy emitted per unit 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. 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 was 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 corrected collimation energy emitted in gamma-rays, $E_\\gamma$. This collimation energy is a measure of the energy released by a GRB taking into account that the energy 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, but is dependent on poorly constrained quantities related to the properties of the medium around the 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 a luminosity distance for $z<1.5$ that is measured by SNe Ia, 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 now 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.  We have analyzed the use of these several correlations in the study of the 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 total amount of matter is fixed, there are only two free variables, $A$ and $\\alpha$, although it is more customary 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, 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 parameter, $\\omega$, of the  dark energy equation of state $p = \\omega \\rho$, and the fraction of of dark matter, $\\Omega_m$. The test of these models is particularly interesting since it is known that they are degenerate for redshifts $z<1$~\\cite{bertolami04,bento05b}. ",
        "conclusions": ""
    },
    "0601/astro-ph0601488_arXiv.txt": {
        "abstract": "We present the first scattered light images of debris disks around a K star (HD 53143) and an F star (HD 139664) using the coronagraphic mode of the Advanced Camera for Surveys (ACS) aboard the {\\it Hubble Space Telescope} (HST).  With ages 0.3 - 1 Gyr, these are among the oldest optically detected debris disks. HD 53143, viewed $\\sim$45$\\degr$ from edge-on, does not show radial variation in disk structure and has width $>$55 AU.  HD 139664 is seen close to edge-on and has  belt-like morphology with a dust peak 83 AU from the star and a distinct outer boundary at 109 AU. We discuss evidence for significant diversity in the radial architecture of debris disks that appears unconnected to stellar spectral type or age. HD 139664 and possibly the solar system belong in a category of narrow belts 20$-$30 AU wide.  HD 53143 represents a class of wide disk architecture with characteristic width $>$50 AU. ",
        "introduction": "The configuration of our solar system is perhaps the most significant starting point for our understanding of planet formation.  Therefore a fundamental question is whether or not the architecture of our solar system is common relative to other planetary systems. One point of comparison is the structure of our Kuiper Belt relative to other systems, which are typically seen as debris disks in scattered light or thermal emission. In scattered light, some debris disks, such as $\\beta$ Pic and AU Mic, have central holes, but are radially extended to hundreds of AU radii \\citep{smith84, kalas04}.  Other debris disks, such as HR 4796A and Fomalhaut consist of relatively narrow rings with sharp inner and outer boundaries \\citep{schneider99, kalas05a}. However, a narrow-belt architecture has not previously been detected in scattered light among stars similar in spectral type and age to the Sun.  HD 53143 (K1V) and HD 139664 (F5V) are two stars $\\sim$18 parsec from the Sun known to have circumstellar dust due to excess thermal emission at far-infrared wavelengths (Aumann 1985; Stencel \\& Backman 1991; Table 1). Various indicators place the age of HD 53143 at $1.0 \\pm 0.2$ Gyr \\citep{decin00, song00, nord04}, whereas HD 139664 may be a younger system with age  $0.3^{+0.7}_{-0.2}$ Gyr \\citep{lach99, montes01, mallik03, nord04}. For these two stars the infrared excess corresponds to a dust mass 3 - 10 times smaller than that of $\\sim$10 Myr-old systems such as AU Mic and $\\beta$ Pic \\citep{kalas04}. Direct imaging of debris disks with masses this small is observationally challenging, but it is now feasible using the optical coronagraph in ACS. ",
        "conclusions": "Taking a census of eight debris disks resolved in scattered light and the predicted dust distribution in our Kuiper Belt, we observe two basic architectures that are not correlated with stellar mass and luminosity, but must depend on other environmental factors (Table 2). Debris systems are either narrow belts or wide disks.  Both types have central dust depletions, but the distinguishing characteristic is the presence or absence of a distinct outer edge.  HR 4796A, Fomalhaut, HD 139664 and the Sun are examples of narrow-belt systems.  The belt systems appear to have radial widths ranging between 20 and 30 AU, and the inner edges may begin as close as 25 AU (Sun), or as far as 133 AU (Fomalhaut). HD 32297, $\\beta$ Pic, HD 107146, HD 53143 and AU Mic  are examples of disks with sensitivity-limited outer edges that imply disk widths $>$50 AU.  The F, G, and K stars are interesting because {\\it a priori} we might expect to find planetary systems similar to our own, yet we discover significant diversity in the outer regions that correspond to Neptune and our Kuiper Belt. With age $\\sim$1 Gyr, HD 53143 is among the oldest known extrasolar debris disks, yet its wide-disk architecture resembles that of the $\\sim$10 Myr old systems of $\\beta$ Pic and AU Mic.   The lingering dust mass (Table 1) throughout the system could signal the absence of giant planets that otherwise sweep clear the parent bodies (comets and asteroids) responsible for the producing the dust disk.  Yet the presence of a dust disk out to at least 110 AU radius shows that the primordial circumstellar disk probably contained the prerequisite mass of gas and dust to form giant planets. By contrast, the younger system HD 139664 has already developed a narrow-belt architecture.  Narrow-belt architectures for the underlying population of planetesimals may originate from early stochastic dynamical events, such as a close stellar flyby, that strip disk mass and dynamically heat the surviving disk \\citep{ida00, adams01}.  Theoretical simulations show that a reduction in disk mass, combined with dynamical heating, produces a less stable planetary system that is more likely to eject giant planets from their formation site to much larger radii, as has been proposed for the origin of Neptune \\citep{thommes99,tsiganis05}.  Therefore, planetesimal belts not only evolve into a narrow structures because of external stochastic events, but they may be found at large distances from the central star due to the subsequent outward migration of interior planets.  The collisionally replenished dust population will spread away from any narrow belt of planetesimals.  If the scattered light appearance continues to manifest a narrow structure, then both the inner and outer edges are probably maintained by other gravitational perturbers such as stellar or sub-stellar companions.   However, only HR 4796 has a known stellar companion that may truncate the outer radius of the dust belt \\citep{aug99}. If there is no confinement mechanism for the outer radius, then the architecture will manifest as a wide-disk.  For example, $\\beta$ Pic and AU Mic may have narrow belts of planetesimals \\citep{aug01, strubbe05}, but the observed dust disk widths extend to hundreds of AU. Though the predicted structure of our Kuiper Belt places the solar system among the narrow disk architectures, it is conceivable that the dust component extends to greater radii \\citep{trujillo01}.  Thus the Sun's classification as a narrow-disk system is tentative."
    },
    "0601/astro-ph0601394_arXiv.txt": {
        "abstract": "It is an observational fact that bulges of spiral galaxies contain a high fraction of old and metal-rich stars. Following this observational fact, we have investigated colors of 21 bulges hosted by a selected sample of high surface brightness spirals and low surface brightness galaxies observed in $B$ and $R$ optical bands and in $J$ and $K_s$ near-IR bands. Using structural parameters derived from these observations we obtain evidence that bulges could be formed by pure disk evolution (secular evolution), in agreement with the suggestion by some authors. The color profiles, especially the near-IR ones show null or almost null color gradients, supporting the hypothesis that the disk stellar populations are similar to those present in the bulge, and/or some bulges can be understood as disks with enhanced stellar density (or pseudobulges). In the optical, half of the galaxies present an inverse color gradient, giving additional evidence in favor of secular evolution for the sample investigated. The comparison of the observed colors with those obtained from spectrophotometric models of galaxy evolution suggests that bulges of the selected sample have solar and subsolar metallicity, and are independent of the current stellar formation rate. Also, we obtain evidence suggesting that galaxies hosting small bulges tend to be systematically metal poor compared to those with larger bulges. These results are being checked more carefully with high S/N spectroscopy. ",
        "introduction": "Low surface brightness (LSB) galaxies have become important members of our extragalactic zoo. Since \\citet{disney1976} showed that the central surface brightness limit of $\\mu_0(B) = 21.65$ \\mucentral ~ in spirals claimed by \\citet{freeman1970} was just a selection effect due to the narrow surface brightness sensitivity of photographic plates, the discovery of LSB galaxies in huge numbers in the universe gave rise to a new field of study in astronomy: the low surface brightness universe. Since then, a wealth of observational data \\citep{longmore1982, bothun1985, bothun1991, mcgaugh1995, oneil2000a} and theoretical work  \\citep{mcleod1995, sprayberry1997, deblok1997, dalcanton1997a, mcgaugh1998a, chiueh2000, bell2002, mcgaugh2003, swaters2003} has been invested not only to describe the most relevant features of these galaxies, but also to analyze whether these features are consistent with models of galaxy formation and evolution. In the observational front, LSB galaxies have been investigated in the optical \\citep{dalcanton1997b, bergvall1999, burkholder2001} in the near-IR \\citep{romanishin1982, bergvall1999, bell2000, galaz2002, monnier2003}, in radio wavelengths, including studies of the HI distribution \\citep{hoeppe1994, longmore1982, deblok1996, matthews2001a}, detected and mapped in CO \\citep{oneil2000b, matthews2001b}, and also studied spectroscopically \\citep{impey2001, bergmann2003}, including the analysis of rotation curves \\citep{mcgaugh2001, deblok2001} and the Tully-Fisher relation \\citep{sprayberry1995, chung2002}.  These references strongly indicate that LSB galaxies are unevolved systems with low stellar formation rate (SFR), low metallicity, small stellar density, relatively high gas fractions and, from dynamical information, large amounts of dark matter. However, several challenging questions still remain. First, in spite of the evidence indicating that LSB galaxies are young when compared with spectrophotometric models, the spread in color is wide, implying that they have probably gone through diverse evolutionary paths. Moreover, the fact that only small amounts of blue stars are needed to make bluer the integrated colors, due to the small stellar density, makes some LSBs appear spuriously younger than they really are. Or, equivalently, small changes in the SFR make galaxies look much younger what they actually are. Second, while emission line analysis of a few spectra of LSB galaxies have revealed the metallicity of the current {\\em gas phase} \\citep{impey2001, bergmann2003}, the metallicity of the {\\em underlying stellar population}, as measured for example using absorption lines, has not been possible to be determined, except in some few cases \\citep{bergmann2003}. Third, the total stellar mass of LSB galaxies is still poorly estimated, mostly due to the large uncertainties in the stellar mass-to-light ratios at optical wavelengths. Only recently, less scattered M/L ratios have been derived using near-IR bands, where the extinction due to dust has a smaller effect than in optical wavelengths \\citep{galaz2002}. Finally, recent modeling of LSB galaxies predict that, under the assumption that blue LSB galaxies are currently undergoing a period of enhanced stellar formation, {\\em there should exist} a population of {\\em red}, non bursting, quiescent LSB galaxies \\citep{gerritsen1999, oneil2000a}.  In the context of the old stellar content of disk and LSB galaxies, the bulge or nuclear component plays a fundamental role as it also affect any analysis where rotation curves are studied \\citep{valenzuela2003,rhee2003}. Although by the proper Hubble morphological definition late type spirals tend to be bulgeless, recent studies with the HST have shown that many late type, face-on spirals apparently bulgeless, have in fact small, tiny bulge \\citep{boker2003} missed by studies where the spatial resolution was not good enough to resolve such small bulges. One of the recurrent questions is whether many of these bulges hosted by LSB galaxies or disks with low stellar density, are in fact high surface brightness bulges embedded in LSB disks. In such a case the bulge could evolve independently with respect to the disk, and one should expect that bulges in LSB disks present more or less the same metallicity than bulges of HSB galaxies.  On the other hand, the secular evolution of galaxies, i.e. the idea that the bulge is formed from disk evolution, allowing the direct formation of bulges from disks in isolated galaxies, implies that the stellar population of disks are similar to those in bulges. For example, one can expect that color gradients are less violent, or even inverse (red to blue, from the center to the outskirt of the galaxy). Also, one can expect that light profiles are fitted only by a combination of exponential profiles, instead of the classic de Vaucouleurs + exponential profile \\citep{courteau1996, kormendy2005}. In this sense, secular evolution could produce what is called a pseudobulge, structurally similar to a bulge, although more comparable to disk when its stellar population is analyzed and/or its light and color profile is studied.  In this paper we explore possible evolutionary scenarios for a set of selected face-on spirals and LSB galaxies, using a combination of optical and near-IR colors. In particular, we investigate the observables and structural parameters in scope of the secular disk to bulge evolution, as stated for example in \\citet{courteau1996}. Most galaxies investigated here are {face-on \\em nucleated} LSB galaxies; all of them are part of the catalogue published by \\citet{impey1996}. We also investigate the stellar content of the bulges. The goal of this paper is to determine what ranges of metallicities and ages are consistent with bulge and disk formation scenarios, by comparing observed $B$, $R$, $J$ and $K_s$ magnitudes and colors with spectrophotometric models of galaxy evolution. We have selected a subsample with prominent bulges, for which we have carefully studied their structural parameters (bulge and disk scale lengths, surface brightnesses). For this subsample, we present radial brightness profiles and color gradients, both in optical and near-IR bands. Secular evolution is analyzed in scope of these results. In \\S 2 the sample selection is presented. Observations and data reduction are described in \\S 3. In \\S 4 we present results and a basic analysis. In \\S 5 we present a discussion in scope of the bulge formation in HSB and LSB galaxies. We conclude in \\S 6. ",
        "conclusions": "We have presented an analysis of a sample of 21 face-on nucleated HSB and LSB spirals selected from the catalogue of \\citet{impey1996}. We have studied their bulges colors and structural parameters in scope of bulge and disk formation, obtaining the following relevant results.  First, bulge and disk scale lengths appear to be marginally correlated, weakly supporting a secular evolutionary process between these two components. However, the relationship between the bulge and disk sizes for HSB seems (B/D $\\sim 0.1$) to be identical to that observed for LSB galaxies, which can be understood in a general context of bulge and disk formation. Also, our results agree with the analysis both from observations presented by \\citet{courteau1996}, and from models by \\citet{mayer2004}, suggesting that bulges could be formed by bar instabilities in early stages of disk evolution. The recently discovered red LSBs, with a notable non-axisymmetric structure and bulge components, are in agreement with many of the LSB models which incorporate massive disks.  Second, structural analysis shows that LSB galaxies follow the same trends (bulge size, disk size, scale lengths, etc.) observed in HSB galaxies, which constitutes evidence that, at least in their structure, LSB galaxies are not a different type of galaxy. On the other hand, we rule out the possibility that HSB galaxies are progenitors of LSBs, basically using the fact that the colors of their bulges are redder compared to those observed in LSBs.  Third, almost half of our sample shows inverse gradients in their $B - R$ radial color profiles (bluer color toward center). Although there is not a straightforward explanation for this, this analyze is key to understand correctly the possible chronology of events which lead to the formation of disk and bulge. Our near-IR surface brightness profiles, do not show inverse color gradients, but present null color gradients which, with the evidence that almost all the light profiles are well fitted by exponential forms (in the optical and in the near-IR), support the secular evolution for these face-on nucleated galaxies and the existence of pseudobulges.  Fourth, using a spectrophotometric model of galaxy evolution, and using the fact that the color index $J - K_s$ is metal sensitive, we find evidence suggesting that bulges of LSBs are metal poor compared to those hosted by HSBs, extending the overall result from emission line analysis that stellar formation regions of LSBs are also metal poor compared to  HSB ones \\citep{deblok1998, galaz2005}. Using the synthesis model PEGASE, we have shown that a possible small amount of dust in the bulges does not play an important role on the bulges color and structural parameter properties, and that the extinction in $B - R$ is smaller than 0.3 mags. We arrive to the same conclusions of \\citet{bell2000}. Moreover, we obtain that bulges hosted by LSB galaxies tend to be small ones. In turn, small bulges are metal poor, compared to larger bulges. This agree with the secular evolution model, in which bulges are built up secularly from the disk, where small bulges have been subject of weak stellar formation processes and slower accretion, leading to a metal poor stellar population in average."
    },
    "0601/astro-ph0601677_arXiv.txt": {
        "abstract": "{The surface magnetic field configuration of the Ap star HD~118022 (78 Vir) has been reconstructed in the framework of the magnetic charge distribution (MCD) method from the analysis of Stokes $IQUV$ spectra obtained using the MuSiCoS spectropolarimeter at Pic du Midi Observatory. Magnetically-sensitive Fe~{\\sc ii} lines were primarily employed in the analysis, supposing that iron is evenly distributed over the stellar surface. We show that the Stokes $IQUV$ profile shapes and variations of 78 Vir can be approximately fit assuming a global magnetic field configuration described by a slightly decentered, inclined magnetic dipole of polar surface intensity approximately 3.3~kG. The derived inclinations of the stellar rotational axis to the line of sight $i=24\\pm 5\\degr$ as well as to the magnetic dipole axis $\\beta=124\\pm5\\degr$ are in good agreement with previous estimations by other authors, whereas the sky-projected position angle\\thanks{ $\\Omega$ increases clockwise from the axis to the North Celestial Pole and relates to the azimuth angle $\\Theta$ specified by Landolfi at al.~(\\cite{Landolfi+93}) as $\\Omega=360\\degr-\\Theta$.} of the stellar rotation axis $\\Omega\\sim110\\degr$ is reported here for the first time. In addition, several lines of Cr~{\\sc ii} and Ti~{\\sc ii} were studied, yielding evidence for non-uniform surface distributions of these elements, and magnetic field results similar to those derived from Fe. ",
        "introduction": "}   78 Virginis (HD~118022) is a bright Ap star, and the first star other than the sun in which magnetic field was discovered (Babcock~\\cite{Babcock47}). The longitudinal magnetic field of 78 Vir, as first observed by Babcock, is variable, and constantly negative. 78 Vir is also marginally variable in broadband light, in radial velocity, and shows variations of its line profiles (Preston~\\cite{Preston69}). As suggested by Preston (\\cite{Preston69}), the observed properties of 78 Vir can be explained in the framework of the oblique rotator model (Stibbs~\\cite{Stibbs50}), in which the star rotates with a period of approximately 3.7 days and has a rotation axis forming a small angle with respect to the line of sight.  The first comprehensive modelling of the surface magnetic field structure of 78 Vir was performed by Borra (\\cite{Borra80}). Borra obtained high-resolution measurements of circular polarisation across the Fe\\,{\\sc ii} $\\lambda$4520.2\\AA\\, line, which he attempted to reproduce using nine different magnetic field configuration models. As argued by Borra (\\cite{Borra80}), the profiles were well reproduced with an inclined dipole geometry that contains a moderate quadrupolar component. Nevertheless, he also pointed out that a decentered ($a$=0.2) dipole model provides {essentially} the same fit quality. Borra also concluded that the 78 Vir presents to an observer a remarkably uniform magnetic field over most of its visible disk, at most phases. It is unfortunate that due to the small inclination of the rotational axis (about $25\\degr$) a significant part of the stellar surface remains hidden - a part which might have a less uniform field geometry.  The analysis of broadband linear polarization (BBLP) measurements of several Ap stars (including 78 Vir) by Leroy~et~al.~(\\cite{Leroy+96}) suggests that the magnetic fields of many Ap stars exhibit departures from the standard oblique rotator model assuming a pure dipole field geometry. According to those authors, significant discrepancies between the BBLP observations and the ``canonical model'' results {(Landolfi~et~al. \\cite{Landolfi+93})} required the assumption of local departures from a dipolar field. In particular, Leroy~et~al. (\\cite{Leroy+96}) showed that the BBLP variations of most stars could be reproduced assuming a dipole magnetic field geometry with slightly expanded field lines over some parts of the magnetic equator. For 78 Vir, the observed BBLP variations do not show especially strong departures from the dipolar case, although according to Leroy~et~al.~(\\cite{Leroy+96}) a modified dipolar model provides a somewhat better fit to the BBLP curves for this star.   In order to better constrain the magnetic field configuration of 78 Vir, in this paper we undertake a more detailed modelling of the surface magnetic field structure based on the analysis of high resolution line profiles in all 4 Stokes parameters. In Sec.~\\ref{obs} we discuss the properties of obtained Stokes $IQUV$ spectra, and in Sec.~\\ref{fund} we summarise the fundamental characteristics of 78 Vir. In Sec.~\\ref{mod} we describe the features of the magnetic field modelling framework, and derive the global magnetic field configuration of the star, comparing observed and computed Stokes profiles for various spectral features. In Sec.~\\ref{discuss} we discuss the derived global magnetic field characteristics and their agreement with earlier studies. ",
        "conclusions": "\\label{discuss}  The abundance of Fe\\,{\\sc ii} derived in this study is $\\log{Fe/N_{\\rm tot}}=-3.16\\pm 0.20$ (see Table~\\ref{fe2sm}). From analysis of spectra of 78 Vir in the regions 3850-3870~\\AA\\ and 3868-4650~\\AA\\ Adelman~(\\cite{Adelman73b}) obtained an Fe\\,{\\sc ii} abundance of $\\log{Fe/N_{\\rm H}}=-2.89$dex. The differences in the iron abundances {may be due to the inclusion/exclusion of polarized transfer, microturbulence, effects of Balmer line wings, etc.}, as well as from the different effective temperatures that have been employed for the abundance analysis (for 78 Vir Adelman~(\\cite{Adelman73a}) applied $T_{\\rm eff}=9950$K, significantly higher than the temperatures employed in our model. %  For the 11 different Fe features studied here, we observe no systematic difference in strength between weak and strong lines. Therefore, the results of this study are consistent with the absence of important Fe stratification in the atmosphere of 78~Vir.  According to our results the global magnetic field structure of 78 Vir is well-described by a slightly decentered magnetic dipole. Usually the starting point of a simulation (in the free parameter hyperspace) is chosen to correspond to a central magnetic dipole, but for all the analysed lines the final model results in a slightly decentered magnetic dipole. This fact is in good qualitative agreement with the results previously obtained by Borra (\\cite{Borra80}). In that paper Borra considered the classical decentered dipole model, where the magnetic dipole size is insignificant in comparison with the stellar radius, and obtained $a_{\\rm 0}$=0.2. Our model provides a smaller $a_{\\rm 0}=0.006\\pm 0.002$ due to the non-zero dipole size ($2a=0.008\\pm 0.003$), that allows us to take into account the non-symmetrical (with respect to the stellar center) magnetic field configuration of the star. The magnetic field intensity and location of the positive and negative magnetic poles in the stellar rotational reference frame are:  \\begin{equation}\\label{poles} \\begin{array}{lll} B_{p}\\!=\\!3.4\\pm 0.7~{\\rm kG},&\\!\\lambda_{p}\\!=\\!-9^\\circ\\!\\pm 5^\\circ,& \\!\\delta_{p}\\!=\\!-33^\\circ\\!\\pm 12^\\circ; \\\\ B_{n}\\!=\\! -3.3\\pm 0.7~{\\rm kG},&\\!\\lambda_{n}\\!=\\!171^\\circ\\!\\pm 5^\\circ,& \\!\\delta_{n}\\!=\\!33^\\circ\\!\\pm 12^\\circ. \\end{array} \\end{equation}   \\noindent Due to the shift of the dipole from the stellar centre, the positive magnetic pole has a slightly stronger field intensity than the negative pole, but not all the analysed lines are equally sensitive to this difference. Nevertheless, the final model is hardly distinguishable from a symmetric central dipole model, which provides a slightly poorer agreement between the simulated and observed data. For Fe\\,{\\sc ii} lines % the centered dipole model results in a $\\chi^{2}$-function which exceeds by 3\\%$\\div$5\\% the best fit result from the decentered dipole model. This fact can be related to the quality of the data and to the model sensitivity. The dipole offset could be a result of the particular geometry of 78 Vir, because we never completely see the positive magnetic pole of the star. The modeled global magnetic field configuration of 78 Vir is illustrated in Fig.~\\ref{structure} according to the best fit results obtained for Fe\\,{\\sc ii} $\\lambda$5018\\AA. Five of the Fe\\,{\\sc ii} lines analysed using the entire Stokes vector result in a plane-of-sky orientation of the rotational axis of $\\Omega=110\\degr\\pm17\\degr$, while the last two lines (Fe\\,{\\sc ii} $\\lambda$6432\\AA\\, and $\\lambda$6516\\AA), with comparatively weaker variability of the observed Stokes $Q$ and $U$ profiles, provide higher values of this angle. Fig.~\\ref{XOmega} shows the dependence of the $\\chi^{2}_{\\rm Q}$ and $\\chi^{2}_{\\rm U}$ functions on the angle $\\Omega$ for the three strong Fe\\,{\\sc ii} lines (see Subsec.~\\ref{strong}). Both functions reach their minima at $\\Omega=109\\degr$ and $\\Omega=289\\degr$. The second value describes a model which is indistinguishable from the configuration specified in Table~\\ref{fe2sm} due to the definition of the Stokes Q and U parameters.   \\begin{figure}[th] \\includegraphics[scale=1.0]{4326f11.eps} \\caption{Illustration of the surface magnetic field geometry of 78 Vir. The location of the positive rotation pole is indicated by the open circle. We always see the negative magnetic pole and only for phase $\\varphi=0.0$ are the positive and negative magnetic poles visible simultaneously. The stellar image has been rotated in the plane of the page by the positional angle $\\Omega$=110\\dr. \\label{structure} } \\end{figure}  As discussed in Sect. 3, 78 Vir is a probable member of the Ursa Major stream. It approaches us with a mean radial velocity $V_{\\rm r}=-8.1\\pm1.0$ km~s$^{-1}$ (Wade~et~al.~2000a). This velocity varies by about $\\pm2$ km~s$^{-1}$ with the phase of stellar rotation (Preston~\\cite{Preston69}, this work). The possible variability of the radial velocity is not taken into account in the simulation procedure. Respectively, our estimations of the radial velocity for different Fe\\,{\\sc ii} lines are distributed around the aforementioned value within the range $\\pm1.5$ km~s$^{-1}$.   \\begin{figure*}[th] \\begin{tabular}{@{\\hspace{-0.07in}}c@{\\hspace{+0.07in}}c} a) & b) \\\\ \\includegraphics[width=3.5in]{4326f12a.eps} & \\includegraphics[width=3.5in]{4326f12b.eps} \\end{tabular} \\caption{Dependence of a) $\\chi^{2}_{\\rm Q}$ and  b) $\\chi^{2}_{\\rm U}$ on the angle $\\Omega$ calculated for three strong lines Fe\\,{\\sc ii} $\\lambda$4923.927\\AA\\, (solid line), $\\lambda$5018.44\\AA\\, (dashed line) and $\\lambda$5169.03\\AA\\, (dotted line). The calculations were performed under the condition of stability of all the best fit parameters for these lines (see Table~\\ref{fe2sm}) except the $\\Omega$. } \\label{XOmega} \\end{figure*}   Nevertheless, the $V_{\\rm r}$ variability through the rotational cycle has been determined from the strong Fe\\,{\\sc ii} lines $\\lambda$4923.927\\AA\\, $\\lambda$5018.44\\AA\\, and $\\lambda$5169.033\\AA\\, (see Subsec.~\\ref{strong}) and for the other Fe\\,{\\sc ii} lines as well. We show that the mean radial velocity derived from the simulation varies from line to line (see Table~\\ref{fe2sm}). The other Fe\\,{\\sc ii} lines provide a similar $\\Delta V_{\\rm r}$ variability. The precision of the mean radial velocity determination depends on the observational errors and on the quality of the description of the surface magnetic field structure. For the blended profiles it also depends on the accuracy of oscillator strengths.   The analysed lines show weak, coherent variations of $\\Delta V_{\\rm r}$ with the period of stellar rotation (see Fig.~\\ref{VelocityDiff}). This suggests a possible moderately non-uniform distribution of Fe\\,{\\sc ii}. Nevertheless, the non-uniform iron distribution will not lead to a significantly different surface magnetic field structure. The simulation of Cr\\,{\\sc ii} lines with the assumption of a non-uniform chromium distribution resulted in almost the same field structure that we obtained from the simulation of Fe\\,{\\sc ii} lines.   The other best fit parameters are the rotational axis inclination $i=24\\degr\\pm5\\degr$ and the dipole axis obliquity $\\beta=124\\degr\\pm5\\degr$, which are in a good agreement with the estimates of Leroy~et~al.~(\\cite{Leroy+96}). The derived $V_{\\rm e}\\sin{i}=12\\pm1~{\\rm km~s^{-1}}$ provides $V_{\\rm e}=29\\pm 4~{\\rm km~s^{-1}}$ and is consistent with that of Preston~(\\cite{Preston71}).       The derived surface magnetic field variability interval ranges from 2.1~kG to 3.2~kG and covers the value $B_s$=2.9~kG estimated by Preston~(\\cite{Preston71}). Besides, as Fig.~\\ref{Blon1} shows, the longitudinal magnetic field variation obtained from the two Fe\\,{\\sc ii} lines simulation is also in good agreement with the LSD $B_{\\rm l}$ data (Wade~et~al.~\\cite{Wade+00b}). These facts justify the applicability of the slightly decentered magnetic dipole model with the aforementioned values of free parameters for the global magnetic field structure description at 78 Vir.  Unfortunately in the case of 78 Vir the intensity of the Stokes $Q$ and $U$ profiles is similar to the observational errors. The data are unsuitable for performing an analysis of the small-scale structure of the magnetic field, using a more sophisticated technique such as Magnetic Doppler Imaging (MDI; Kochukhov~\\&~Piskunov~\\cite{K+P02}). However, it appears, given the relatively good agreement between the intensity, structure and variability of the observed and simulated Stokes $QU$ profiles, that the real magnetic field of 78 Vir does not depart strongly from the configuration derived here. At the same time, we have noted that the Stokes $QU$ profiles are not, at some phases, fit to within their errors. Therefore these data, notwithstanding their relatively low S/N, already suggest the limitations of the MCD model framework.    A similar analysis of the Stokes $IV$ profile variability was performed by Kochukhov et al.~(\\cite{Kochukhov+02}) for $\\alpha^2$~CVn. Those authors applied MDI using ``multipolar regularization'', analysing the Stokes $IV$ line profiles and obtained a good agreement between observed and computed profiles for Fe\\,{\\sc ii}, Cr\\,{\\sc ii} and Si\\,{\\sc ii} lines. The quality of the Stokes $I$ and $V$ profile fits obtained in this paper is only marginally poorer than that obtained by Kochukhov et al.~(\\cite{Kochukhov+02}), although remarkably the magnetic field model framework employed here is much simpler.     New spectro-polarimetric observations of 78 Vir with higher spectral resolution and signal-to-noise ratio (especially for the Stokes $Q$ and $U$ spectra) are required for further refinement of the global magnetic field structure. With such data, a model with additional parameters (such as MDI) could be applied, in order to study the local magnetic field topology and the detailed relationship between the magnetic field and the abundance distributions."
    },
    "0601/astro-ph0601349_arXiv.txt": {
        "abstract": "{In the central part of this paper, we revisit the classical study of the $H$-function defined as the unique solution, regular in the right complex half-plane, of a Cauchy integral equation. We take advantage of our work on the $N$-function published in the first article of this series. The $H$-function is then used to solve a class of Cauchy integral equations occurring in transfer problems posed in plane-parallel media. We obtain a concise expression of the unique solution analytic in the right complex half-plane, then modified with the help of the residue theorem for numerical calculations.} ",
        "introduction": " ",
        "conclusions": "We used the theory of Cauchy integral equations to the solution of equations in the form (1) which appear in the linear transport theory assuming isotropic scattering. We felt necessary a revisited description of the solution to this type of equation for the following reasons:  \\noindent 1) The pionneering studies of the sixties may be presented more simply than formerly by introducing the values on the cut of any function defined by a Cauchy-type integral.  \\noindent 2) In the fifties and sixties, the factorization relation (27) and the non-linear $H$-equation (31) were well-trodden starting points for the study of the $H$-function. However, no completely self-consistent approach was available starting from the linear $H$-equation (17) as constrained by Eq. (18). For instance the papers by Mullikin et al.$^{[3,9]}$ presuppose the factorization relation (27) usually established from the Wiener-Hopf method. In Case and Zweifel's mono\\-graph,$^{[10]}$ the technique of invariant imbedding is used to establish the relation connecting the $H$-function to the $N$-function.  \\noindent 3) The interest in studying the $H$-function from its definition (17)-(18) is not only of pedagogical essence: this approach can be enlarged to the study of finite media. The difficulty in solving Eq. (1) in a finite slab originates in the dependence of the free term $c(z)$ on the unknown function $f(z)$. It will be dealt with in the third paper of this series with the help of formula (50), since the solution has to be analytic in the right complex half-plane (and in the left one as well)."
    },
    "0601/astro-ph0601380_arXiv.txt": {
        "abstract": "Stellar atmosphere models of ionized accretion disks have generally neglected the contribution of magnetic fields to the vertical hydrostatic support, although magnetic fields are widely believed to play a critical role in the transport of angular momentum.  Simulations of magnetorotational turbulence in a vertically stratified shearing box geometry show that magnetic pressure support can be dominant in the upper layers of the disk.  We present calculations of accretion disk spectra that include this magnetic pressure support, as well as a vertical dissipation profile based on simulation. Magnetic pressure support generically produces a more vertically extended disk atmosphere with a larger density scale height.  This acts to harden the spectrum compared to models that neglect magnetic pressure support.  We estimate the significance of this effect on disk-integrated spectra by calculating an illustrative disk model for a stellar mass black hole, assuming that similar magnetic pressure support exists at all radii. ",
        "introduction": "Almost without exception, models of geometrically thin accretion disks that are used to predict radiation spectra in order to compare with observation are based on the alpha prescription introduced by \\citet{sha73}. This prescription enables one to compute the radial distribution of surface density in the disk.  When this is combined with additional assumptions about the vertical distribution of dissipation, complete models of the vertical structure at each radius and the local radiation spectrum can be computed.  In black hole accretion disks, for example, the most detailed of such models so far are those of \\citet{hub01} for supermassive black holes, \\citet{dav05} for stellar mass black holes, and \\cite{hui05} for intermediate mass black holes.  The models fully account for relativistic effects, and include a detailed non-LTE treatment of level populations of hydrogen, helium, and abundant metals.  Continuum opacities due to bound-free and free-free transitions, as well as Comptonization, are included.  In spite of this level of sophistication, the models are limited by other assumptions:  the disk is assumed to be stationary, the alpha-prescription is used to calulate the radial distribution of surface density, a no-torque inner boundary condition is assumed, the vertical structure at each radius is assumed to depend only on height and is symmetric about the midplane, the vertical distribution of dissipation per unit mass is assumed to be constant, and the disk is assumed to be supported vertically against the tidal field of the black hole by just gas and radiation pressure.  It is now widely believed that the anomalous stress in accretion disks is a result of sustained magnetohydrodynamical turbulence which is akin to that seen in the nonlinear development of the magnetorotational instability, or MRI \\citep{bal98}.  In principle, shearing box simulations of this turbulence which incorporate the vertical tidal field of the central mass \\citep{bra95, sto96, mil00} can be used to build models of the vertical structure of accretion disks that are free of the ad hoc assumptions that plague existing models.  The most recent of these simulations \\citep{tur04,hir05} are particularly useful for this purpose as they evolve the equations of radiation magnetohydrodynamics and have self-consistent thermodynamics.  While much of the dissipation of the turbulence in these simulations is numerical and happens on the grid scale, one hopes that this still mocks up the microscopic dissipation occuring at the smallest scales of the turbulent cascade, and in any case provides a better handle on the true spatial distribution of dissipation than the ad hoc assumptions used in alpha-disk models.  Moreover, in radiation pressure dominated simulations of MRI turbulence, dissipation of compressive motions by photon diffusion is fully resolved \\citep{tur03}.  \\citet{dav05} used the time and horizontally-averaged vertical dissipation profile from the simulation of \\citet{tur04} to compute the vertical structure and spectrum of an annulus, and compared it to a model based on the usual ad hoc assumption of constant dissipation per unit mass.  While the resulting profiles of density and temperature were very different, the emergent spectrum was very similar.  In fact, \\citet{dav05} found that the locally emergent spectra of individual annuli were remarkably robust to different vertical dissipation profiles and to different overall stress prescriptions, provided most of the dissipation occurs deeper than the effective photosphere.  Part of this robustness stems from the fact that the radiative flux becomes constant above the dissipation region, and so the radiative force per unit mass also becomes constant if the flux mean opacity is dominated by electron scattering in the surface layers.  Because the tidal gravity of the central mass increases with height, gas pressure gradients must therefore compensate, and this results in very steep density gradients near the photosphere.  The density and temperature at the effective photosphere then become roughly independent of the details of the vertical structure, resulting in a robust emergent spectrum.  In addition to a markedly different dissipation profile, vertically stratified MRI simulations also show that the surface layers of a disk will be {\\it magnetically} supported.  This additional support removes the need for steep density gradients, and in fact results in a much more extended atmosphere with a larger density scale height.  As qualitatively suggested by \\citet{hir05}, magnetic pressure support could therefore dramatically alter the emergent spectrum, much more than alterations of the dissipation profile deeper than the effective photosphere.  It is this new effect arising from MRI simulations that we investigate in this paper.  In section 2, we discuss the time and horizontally-averaged dissipation and magnetic pressure profiles that we computed from the simulations. We then used these profiles within the stellar atmosphere code TLUSTY \\citep{hub88,hub95} to compute the vertical structure and emergent spectrum of individual annuli.  We present the results of these calculations in section 3, including an illustrative whole disk spectrum for an accreting stellar mass black hole.  We then summarize our conclusions in section 4. ",
        "conclusions": "The advent of thermodynamically self-consistent simulations of MRI turbulence in vertically stratified shearing boxes \\citep{tur04,hir05} has finally opened the door to the construction of spectral models which are based on the physics of the turbulence itself.  Surprisingly, our investigation in this paper indicates that the thermodynamics itself, i.e. the vertical profile of dissipation, has little effect on the emergent spectrum compared to models based on the the standard assumption of a constant dissipation rate per unit mass.  This conclusion was also reached in the earlier investigation of \\citet{dav05}.  Unless an annulus in the disk is only moderately effectively thick, most of the dissipation occurs deeper than the effective photosphere, and its vertical profile therefore has little effect on the emergent spectrum.  The reader should bear in mind, however, that the numerical dissipation profiles, e.g. equation (\\ref{eqdfdmfit}), are based on simulations in which mechanical energy is simply lost at the grid scale and replaced by internal energy.  More work needs to be done to investigate whether these profiles are robust to changes in grid resolution.  In addition, the simulations were for gas pressure dominated annuli, and large uncertainties still remain to be fully investigated in the radiation pressure dominated case, which is more relevant for the innermost annuli of luminous accretion disks around black holes and neutron stars.  In contrast to the thermodynamics, the actual {\\it dynamics}, i.e. the contribution of magnetic pressure to the vertical support of the disk against gravity, does produce significant changes in the emergent spectrum when compared to models based on the standard ad hoc assumptions.  Magnetic pressure support at high altitude results in a much more extended density scale height, because gas pressure is not required to support the atmosphere against the increasing vertical gravity at high altitude.  Because absorption opacity generally increases with density, a larger density scale height results in a smaller density at the effective photosphere in order to produce the same effective optical depth of unity.  As anticipated by \\citet{hir05}, this enhances non-LTE effects in the spectrum because collisional processes in the plasma are then diminished compared to radiative processes.  Even in LTE, lower density at the effective photosphere generally implies that the plasma becomes more ionized, reducing the bound-free opacity.  Electron scattering is enhanced compared to absorption, driving the emergent radiation spectrum closer to a modified blackbody.  All of these effects mean that the emergent spectrum will be generically harder compared to an atmosphere where magnetic support is neglected.  We emphasize that these spectral implications will apply to all accretion disks in which the MRI is thought to act, including active galactic nuclei, cataclysmic variables, and X-ray binaries.  In the one whole disk model we constructed, around a stellar mass black hole, we found that the magnitude of this hardening was equivalent to choosing a color correction factor in a relativistic, multitemperature blackbody of 1.74, as opposed to 1.48 in the nonmagnetized case. This has important implications for recent attempts to measure black hole spins in X-ray binaries by continuum spectral fitting (e.g. \\citealt{sha05}). Because magnetic pressure support hardens the spectrum, the fitted black hole spin to the observed continuum will have to compensate by being smaller.  It is conceivable that magnetic pressure support may be exaggerated in the local shearing box simulations.  Long wavelength Parker instability modes might be suppressed because they cannot fit inside the box.  This is an issue which will need further investigation, but we did a preliminary check by comparing with global, non-radiative, general relativistic MRI simulations \\citep{dev03}.  In the shearing box simulations, magnetic pressure becomes dominant when the density falls below 0.03 times the midplane density.  In contrast, in a global simulation in Schwarzschild geometry, the magnetic pressure becomes dominant below 0.1 times the local midplane density.  Magnetic support therefore appears to be even more important in the global simulation which, however, was non-radiative.  The shearing box simulations show significant time-dependence, horizontal structure, and asymmetries above and below the disk midplane, effects which we have ignored by ruthless averaging.  More sophisticated radiative transfer techniques will have to be employed to investigate their effect on the emergent spectrum. Density irregularities in particular are expected to be even more prominent in the radiation pressure dominated inner regions of black hole accretion disks \\citep{tur03, tur04,tur05}.  Monte Carlo simulations of the photon transfer through such inhomogeneous structures suggest that the enhanced ratio of absorption to scattering opacity in the denser regions helps to thermalize (and therefore soften) the emergent spectrum \\citep{dav04}, at least when the inhomogeneities are treated as static structures.  Whether or not this persists in a time-dependent calculation, or whether or not the hardening due to magnetic pressure support turns out to be the more important effect, remains to be seen."
    },
    "0601/astro-ph0601455_arXiv.txt": {
        "abstract": "We present detailed optical, X-ray and radio observations of the bright afterglow of the short gamma-ray burst 051221a obtained with Gemini, {\\it Swift}/XRT, and the Very Large Array, as well as optical spectra from which we measure the redshift of the burst, $z=0.5464$.  At this redshift the isotropic-equivalent prompt energy release was about $1.5\\times 10^{51}$ erg, and using the standard afterglow synchrotron model we find that the blastwave kinetic energy is similar, $E_{K,\\rm iso}\\approx 8.4\\times 10^{51}$ erg.  An observed jet break at $t\\approx 5$ days indicates that the opening angle is $\\theta_j\\approx 7^\\circ$ and the total beaming-corrected energy is therefore $\\approx 2.5\\times 10^{49}$ erg, comparable to the values inferred for previous short GRBs.  We further show that the burst experienced an episode of energy injection by a factor of 3.4 between $t=1.4$ and 3.4 hours, which was accompanied by reverse shock emission in the radio band.  This result provides continued evidence that the central engines of short GRBs may be active significantly longer than the duration of the burst and/or produce a wide range of Lorentz factors.  Finally, we show that the host galaxy of GRB\\,051221a is actively forming stars at a rate of about 1.6 M$_\\odot$ yr$^{-1}$, but at the same time exhibits evidence for an appreciable population of old stars ($\\sim 1$ Gyr) and near solar metallicity.  These properties are intermediate between those of long GRB hosts and those of previous short bursts suggesting that the progenitor lifetimes may have a large spread. The lack of bright supernova emission and the low circumburst density ($n\\sim 10^{-3}$ cm$^{-3}$), however, continue to support the idea that short bursts are not related to the death of massive stars and are instead consistent with a compact object merger.  Given that the total energy release is a factor of $\\sim 10$ larger than the predicted yield for a neutrino annihilation mechanism, this suggests that magnetohydrodynamic processes may be required to power the burst. ",
        "introduction": "\\label{sec:intro} Following over a decade of speculation about the nature of the short-duration hard-spectrum gamma-ray bursts (GRBs), the recent discovery of afterglow emission from these bursts provided the first physical constraints on possible progenitor models.  We now know that short GRBs occur at cosmological distances, $z\\gtrsim 0.1$, have energy releases of about $10^{49}$ erg when corrected for relatively wide opening angles of $\\sim 10^\\circ$, are associated with both star-forming and elliptical galaxies, and are not associated with bright supernovae \\citep{bcb+05,bpc+05,bpp+05,ffp+05,gso+05,hsg+05,hwf+05,vlr+05}. These properties indicate that short GRBs are associated with an old stellar population, and are broadly consistent with expectations of the popular neutron star (NS) and/or black hole (BH) coalescence models (e.g., \\citealt{elp+89,npp92,kc96,rr02,ajm05}).  On the other hand, some of the afterglow observations, in particular the detection of soft X-ray tails and X-ray flares on a timescale $\\gtrsim 10^2$ s \\citep{bcb+05,vlr+05}, cannot be easily explained in the coalescence scenario.  These observations raise the possibility that short GRBs may have different or multiple progenitors systems (e.g., accretion-induced collapse of a neutron star; \\citealt{mrz05}), or that some key physical processes in the energy extraction mechanism or accretion disk hydrodynamics remain unaccounted for in the simple models \\citep{paz06}.  Secondary indicators such as the redshift and luminosity distributions have also been used to argue against the coalescence model (\\citealt{ngf05}, but see \\citealt{gp05}), but these rely on a very limited sample of events ($\\lesssim 4$ with reasonably robust redshifts) and strong constraints are likely premature.  In the absence of direct evidence for or against the coalescence model from the detection of gravitational waves (e.g., \\citealt{ct02}), or a mildly relativistic outflow \\citep{lp98,k05}, progress in our understanding of the nature and diversity of the progenitor systems and the underlying physics requires a much larger statistical sample than the current data set of two short GRBs with secure redshifts (050709 and 050724) and two with putative redshifts (050509b and 050813).  In addition, only GRBs 050709 and 050724 have broad-band afterglow coverage which allowed a determination of the energy release and circumburst density, and provided evidence for ejecta collimation \\citep{bpc+05,ffp+05,p05}.  With a larger sample of well-studied short GRBs we can begin to address the properties of the progenitor systems and the GRB mechanism through a census of the energy release, the jet opening angles, the density of the circumburst environment, and the relative fraction of short GRBs in star-forming and elliptical galaxies.  In this vein, we present in this paper detailed optical, X-ray and radio observations of the bright \\grb, which we localize to a star-forming galaxy at $z=0.5464$.  We show from the evolution of the various light curves and from modeling of the broad-band afterglow data that the beaming-corrected energy release of this burst is comparable to those of previous short-hard bursts.  In addition, we find that the circumburst density is low indicating that the progenitor system is unlikely to reside in a star forming region. Finally, we show that the GRB experienced an episode of significant energy injection, which was accompanied by reverse shock emission in the radio band.  This, along with an energy release exceeding $10^{49}$ erg, suggests that the energy extraction mechanism leads to a complex ejecta structure and is possibly driven by MHD processes rather than neutrino anti-neutrino annihilation. ",
        "conclusions": "\\label{sec:disc}  We present broad-band optical, radio, and X-ray data for the afterglow of the short-hard burst \\grb.  Due to our rapid identification of the afterglow we were able to obtain a spectrum at $t\\approx 1.1$ day when the afterglow was in principle sufficiently bright to detect absorption features from the interstellar medium of the host galaxy leading to a direct redshift measurement.  Unfortunately, at the redshift of the host galaxy, $z=0.5464$, no strong metal UV lines are redshifted into the waveband of our spectrum.  Still, rapid identification and spectroscopy of future short bursts may potentially lead to direct redshift measurements.  At $z=0.5464$, \\grb\\ currently represents the most distant short GRB for which we have a secure redshift.  More importantly, we emphasize that \\grb\\ is only the third short burst for which a detailed afterglow study has be performed and the physical parameters have been constrained.  The isotropic-equivalent $\\gamma$-ray and kinetic energy releases are $2.4\\times 10^{51}$ and $1.4\\times 10^{51}$ erg respectively.  These values are a factor of 10 and $10^2$ times larger than the values for the short GRBs 050724 and 050709, respectively.  Adopting the results of our standard synchrotron model and using $t_j\\approx 5$, we constrain the collimation of the ejecta to $\\theta_j\\approx 7^{\\circ}$, comparable to the opening angles of GRBs 050709 and 050724 (\\citealt{ffp+05,bpc+05,p05}; but see \\citealt{gbp+06}). A comparison to the compilation of jet opening angles of long GRBs, shown in Figure~\\ref{fig:angles_hist} demonstrates that while overlap exists between the two distributions, the opening angles measured for the three short GRBs are about a factor of two larger than the median value of $5^{\\circ}$ for long GRBs.  This may not be surprising given that in the case of long GRBs the passage of the jet through the stellar envelope is thought to be responsible for its collimation \\citep{zwm03}. On the other hand, in models of compact object mergers it has been argued that the outflow may be collimated by a neutrino driven wind \\citep{rrd03} or the accretion disk itself \\citep{ajm05} both of which lead to relatively wide opening angles, $\\gtrsim 10^{\\circ}$.  Most likely the same holds true for scenarios such as accretion induced collapse of a neutron star or a white dwarf \\citep{yb98,fbh+99}.  With the inferred opening angle, the beaming-corrected energy of \\grb\\ are similar to those of previous short bursts and comparable to the low end of the distribution for long duration events (e.g., GRB\\,031203; \\citealt{sls04,skb+04}).  This is shown in Figure~\\ref{fig:egamma_eag} where we compile the beaming-corrected gamma-ray energy release and ejecta energy values for both long and short duration bursts.  It has been argued that the majority of the long GRBs have a total (prompt plus afterglow) true energy release within $0.5$ dex of $1.7\\times 10^{51}$ erg \\citep{bkp+03}, with $3$ events having a total energy of $\\lesssim 10^{50}$ erg \\citep{skb+04,skn+06}.  The similar $\\gamma$-ray efficiencies of long and short bursts indicate that the energy dissipation process of the prompt emission is most likely the same.  The energy release inferred for \\grb\\ can also be used to constrain the energy extraction mechanism particularly in the context of merger models.  Two primary mechanisms for energy extraction from the central engine have been discussed in the literature, namely $\\nu\\overline{\\nu}$ annihilation or magnetohydrodynamic (MHD) processes, for example the Blandford-Znajek process (e.g., \\citealt{bz77,lrp05,rrd03}).  Numerical simulations indicate that $\\nu\\overline{\\nu}$ annihilation generally produces a less energetic outflow than MHD processes \\citep{lrg05}.  \\citet{rr02} show that even under the most favorable conditions, neutrino annihilation is unlikely to power a burst with a beaming-corrected energy in excess of few $\\times 10^{48}$ erg, while magnetic mechanisms can produce luminosities $\\gtrsim 10^{52}~\\rm erg~s^{-1}$ \\citep{rrd03}.  Given that the beaming-corrected total energy for GRB\\,051221a is a factor of $\\sim 10$ larger than the predicted yield for the $\\nu\\overline{\\nu}$ annihilation model, this may indicate that MHD processes are responsible.  We next compare the circumburst densities of long and short GRBs.  In Figure~\\ref{fig:ed_hist} we show that the circumburst densities for short bursts cluster near $n\\sim 10^{-2}~\\rm cm^{-3}$.  On the other hand, the densities inferred for long GRBs range from about $10^{-2}$ to $10~\\rm cm^{-3}$ with a median value of about $2~\\rm cm^{-3}$. The fact that the circumburst densities of short GRBs cluster at the low end of the distribution for long GRBs supports the notion that short bursts do not explode in rich stellar wind environments but are instead consistent with interstellar densities.  Finally, we show that the host galaxy of GRB\\,051221a is actively forming stars at a rate of about 1.6 M$_\\odot$ yr$^{-1}$, but at the same time exhibits evidence for an appreciable population of old stars ($\\sim 1$ Gyr).  The inferred star-formation rate is at least an order of magnitude larger than the values inferred for previous short GRB hosts and yet is at the low end of the distribution for long GRB hosts. Similarly, the inferred metallicity for the \\grb\\ host galaxy appears to be higher than those of long GRB hosts indicating that it is more evolved.  It is conceivable that the spread in short GRB host properties reflects a dispersion in the progenitor lifetimes such that the progenitor of \\grb\\ may be younger than those of GRBs 050509b and 050724 which occurred in elliptical galaxies (see also \\citealt{pbc+05,gno+05}).  The potential spread in progenitor lifetimes may also be supported by the smaller offset of \\grb\\ compared to previous short bursts, but we caution that offsets are affected not only by progenitor lifetimes but also by the potential dispersion in kick velocities, host masses, and projection effects.  While the basic properties of short GRBs (redshifts, hosts, isotropic $\\gamma$-ray energies) are now available for several events, it is clear that detailed afterglow observations are required for a complete understanding of the burst properties.  With the addition of \\grb\\ to the existing sample of only two well-studied events, we have embarked on a quantitative study of short bursts that over the next few years will shed light on the diversity of progenitors and energy extraction mechanisms."
    },
    "0601/astro-ph0601039_arXiv.txt": {
        "abstract": "{ According to standard models supernovae produce radioactive $^{44}$Ti, which should be visible in gamma-rays following decay to $^{44}$Ca for a few centuries. \\Ti production is believed to be the source of cosmic $^{44}$Ca, whose abundance is well established. Yet, gamma-ray telescopes have not seen the expected young remnants of core collapse events. The \\Ti mean life of $\\tau$\\about~89~y and the Galactic supernova rate of \\about~3/100~y imply \\about~several detectable \\Ti gamma-ray sources, but only one is clearly seen, the 340-year-old Cas A SNR. Furthermore, supernovae which produce much \\Ti are expected to occur primarily in the inner part of the Galaxy, where young massive stars are most abundant. Because the Galaxy is transparent to gamma-rays, this should be the dominant location of expected gamma-ray sources. Yet the Cas A SNR as the only one source is located far from the inner Galaxy (at longitude 112\\degree). We evaluate the surprising absence of detectable supernovae from the past three centuries. We discuss whether our understanding of SN explosions, their \\Ti yields, their spatial distributions, and statistical arguments can be stretched so that this apparent disagreement may be accommodated within reasonable expectations, or if we have to revise some or all of the above aspects to bring expectations in agreement with the observations. We conclude that either core collapse supernovae have been improbably rare in the Galaxy during the past few centuries, or \\Ti-producing supernovae are atypical supernovae. We also present a new argument based on $^{44}$Ca/$^{40}$Ca ratios in mainstream SiC stardust grains that may cast doubt on massive-He-cap Type I supernovae as the source of most galactic $^{44}$Ca. ",
        "introduction": "\\label{sec:intro}  Supernovae are the agents that drive the evolution of gaseous regions of galaxies. As end points of the evolution of massive stars that have formed out of the interstellar gas, their explosions eject matter enriched with freshly formed isotopes and stir interstellar gas. However, the explosions themselves are still not understood \\citep{2000Natur.403..727B, 2003fthp.conf...39J}. Parametric descriptions are used to describe the core collapses (``cc-SN'', supernovae of types II and Ib/c) \\citep{1995ApJS..101..181W, 1996ApJ...460..408T} as well as thermonuclear explosions of white dwarfs (supernovae of type Ia; \\citealt{1997NuPhA.621..467N}). A prominent issue in astrophysics is whether the supernova explosion itself is a well-regulated, robust physical process, or if intrinsic variability over a wider range of physical conditions are rather common.  Supernova homogeneity by type can be studied in different ways. One approach is to analyze the rate of supernovae of a specific type in different environments and over different time scales. In this work we do this by asking if the current rate of supernovae in our Galaxy which produce radioactive \\Ti is in line with expectations from other observables and from supernova theory.  \\Ti decay offers a unique window to the study of supernova rates. Specific aspects of this window are: \\begin{enumerate} \\item Gamma rays penetrate the entire galactic disk with little extinction \\item \\Ti gamma-rays reflect the current rate of supernovae, with the \\Ti mean decay time scale of $\\tau$ = 89 years; this present-day snapshot which has not yet fed back into chemical evolution can be directly related to the observable current population of massive stars \\item Most \\Ti is co-produced with $^{56}$Ni in relatively frequent core collapse supernovae (cc-SN). The radioactive energy of $^{56}$Ni is responsible for well-observed supernova light \\item Nucleosynthesis of \\Ti is primarily from $\\alpha$-rich freeze-out of nuclear statistical equilibrium and secondarily from silicon burning \\item The origin of abundant cosmic $^{44}$Ca occurs mainly through \\Ti nucleosynthesis \\item \\Ti traces have been found in pre-solar grains which have been attributed to condensation within core-collapse supernovae \\end{enumerate}    In this paper we estimate what the gamma-ray sky of \\Ti sources would be expected to look like by adopting an {\\it average \\Ti source model} having a characteristic source event {\\it recurrence rate, \\Ti yield per event}, and {\\it spatial distribution}. We compare this to the present-day gamma-ray survey and find apparent and serious conflicts. Then we analyze whether deviations from these {\\it average} expectations can occur from the known or expected {\\it variability of models and parameters involved}. We use a Monte Carlo simulation of the expected sky image within reasonable distributions of parameters for that purpose. This leads us to discuss each of the relevant parameters, which may explain an anomalous  \\Ti sky; these are, specifically: \\begin{itemize} \\item the gamma-ray survey quality \\item statistical effects of small samples \\item the adopted supernova rates \\item the supernova explosion and nucleosynthesis models \\item the spatial distribution of supernova events \\item the deviations from smooth chemical evolution \\item the supernova origin of Galactic and solar $^{44}$Ca \\end{itemize} ",
        "conclusions": "\\label{sec:conclusions}  The observed distribution of pointlike sources of \\Ti in our Galaxy is inconsistent (at a probability near 10$^{-3}$) with the combined current understanding of \\Ti production, specifically \\begin{itemize} \\item core collapse supernovae eject $10^{-4}$ \\Msol  of \\Ti, within a factor of two, and \\item the galactic core collapse  SN rate is about 3 per century, \\end{itemize} assuming the last few centuries are representative of the steady-state \\Ti production. The inconsistency is exacerbated if we further demand, as currently understood to be so, that the standard solar abundance of $^{44}$Ca originates from \\Ti-producing supernovae. The larger discrepancy is because the product of the above two quantities falls short by a factor of three of the requisite current $^{44}$Ca production rate. The disagreement persists for any combination of rate and yield whose product is the required value, unless the yield of \\Ti is very much higher and the rates very much lower, i.e., $^{44}$Ca is not made primarily by typical supernovae events, but by very rare ones. These might include He-triggered detonations of sub-Chandrasekhar SNe Ia, or rare variants of core collapses, perhaps those most departing from spherical symmetry. Evidence of He-cap SNIa as source of the $^{44}$Ca abundance might be identifiable in $^{44}$Ca/$^{40}$Ca ratios greater than solar in some mainstream SiC grains.  A future survey with gamma-ray line sensitivity of 10$^{-6}\\ cm^{-2}\\ s^{-1}$ would be expected to detect $\\geq$10 sources (Fig. 2), and so could rule out $^{44}$Ca production by frequent supernovae at confidence 5~10$^{-5}$. Regardless of specific assumptions, the Cas A supernova remnant as the brightest \\Ti remnant in the galaxy is a priori very unlikely. That the brightest SNR should be found in the outer galaxy or that it is over 300 years old are each improbable. It suggests that yields higher than suggested by many current calculations are possible, which, of course, makes the lack of other detectable remnants even more puzzling."
    },
    "0601/astro-ph0601275_arXiv.txt": {
        "abstract": "Observations of the gravitational radiation from well-localized, inspiraling compact object binaries can measure absolute source distances with high accuracy.  When coupled with an independent determination of redshift through an electromagnetic counterpart, these standard sirens can provide an excellent probe of the expansion history of the Universe and the dark energy.  Short $\\gamma$-ray bursts, if produced by merging neutron star binaries, would be standard sirens \\textit{with known redshifts} detectable by ground-based GW networks such as \\ligo-II, \\virgo, and \\aigo. Depending upon the collimation of these GRBs, a single year of observation of their gravitational waves can measure the Hubble constant $h$ to $\\sim2\\%$.  When combined with measurement of the absolute distance to the last scattering surface of the cosmic microwave background, this determines the dark energy equation of state parameter $w$ to $\\sim9\\%$.  Similarly, supermassive binary black hole inspirals will be standard sirens detectable by \\lisa. Depending upon the precise redshift distribution, $\\sim 100$ sources could measure $w$ at the $\\sim4\\%$ level. ",
        "introduction": "With the advent of the Laser Interferometer Gravitational-Wave Observatory (\\ligo), we are on the verge of an era of gravitational-wave astronomy~\\cite{barish99,ligo05}. Among the most interesting expected sources for GW observatories are the inspirals and mergers of compact-object binaries.  LIGO-II, a planned upgrade with tenfold increase in sensitivity, can detect the inspirals and coalescence of stellar-mass binaries within several hundred megaparsecs, while the Laser Interferometer Space Antenna (\\lisa) can study supermassive binary BHs ($M\\sim10^4-10^7 M_\\odot$) throughout the universe ($z\\alt10$).  The idea of using GW measurements of coalescing binaries to make cosmologically interesting measurements has a long history.  As originally pointed out by \\citet{schutz}, observation of the gravitational radiation from an inspiraling binary provides a self-calibrated absolute distance determination to the source. \\citet{chernoff93} and \\citet{finn96} took advantage of this property to show how, by observing many inspiral sources, one can construct the distribution of observed binary mass and GW signal strength, and thereby statistically constrain the values of cosmological parameters. More recently, \\citet{holz05} have shown that \\lisa\\ observations of well-localized supermassive binary black hole (SMBBH) inspirals allow cosmological distance determination with unprecedented accuracy, with typical errors $<1\\%$.  These GW ``standard sirens'' can precisely map out the expansion history of the Universe, offering a powerful probe of the dark energy.  The utility of standard sirens for constraining dark energy is quite similar to that of standard candles, like Type-Ia supernovae.  One advantage of GW standard sirens is that the underlying physics is well-understood.  The radiation emitted during the inspiral phase (as opposed to the merger phase) is well described using the post-Newtonian expansion of general relativity for the BH binary \\cite{blanchet}.  Hence some unknown systematic evolution of the standard sirens over time, mimicking a different cosmology, should not be of concern.  Another advantage is that GW observatories directly measure absolutely calibrated source distances, whereas Type-Ia supernova standard candles provide only relatively calibrated distances.  A major drawback of GW standard sirens is that, although the gravitational waveforms measure distance directly, they contain no redshift information. To be useful as a standard candle, an independent measure of the redshift to the source is crucial. This can be determined through observation of an electromagnetic counterpart, such as the host galaxy of the source. Unfortunately, as GW observatories are essentially all-sky, they generally provide poor source localization, and the host galaxy is not always unambiguously identifiable \\cite{zoltan05}. In cases where source redshifts cannot be determined, the distribution of unlocalized events can be used to place statistical bounds on cosmological parameters \\cite{finn96}.  However, in this paper we will focus upon GW sirens whose redshifts may be measured, as they can provide very tight constraints on cosmology.  Because standard siren distances are absolutely calibrated, even sources at low redshift (e.g.\\ $z\\alt0.2$) can constrain dark energy. This may seem surprising, since at low redshifts the distance-redshift relation is well-described by a linear Hubble relation $D = cz/H_0$, independent of dark energy parameters.  As emphasized by \\citet{hujain04} and \\citet{hu04}, however, absolute distances to sources at low redshift tightly constrain dark energy, when combined with a determination of the absolute distance to the last-scattering surface of the cosmic microwave background.  To understand this, note that cosmological distances are given by a redshift integral of the Hubble parameter, which in turn depends on the sum of energy densities at each redshift: \\begin{eqnarray} D(z_s)&=&\\frac{c}{H_0\\sqrt{\\Omega_K}} \\sinh\\left[\\sqrt{\\Omega_K} \\int_0^{z_s}\\frac{H_0}{H(z)}dz\\right] \\\\ \\frac{H(z)}{H_0}&=& \\sqrt{\\Omega_m(1+z)^3+\\Omega_{\\rm de}(1+z)^{3(1+w)}+\\Omega_K(1+z)^2}. \\nonumber \\end{eqnarray} Here $\\Omega_m+\\Omega_{de}+\\Omega_K=1$, $H_0 = 100\\,h$ km/s/Mpc is the Hubble constant today, and we have assumed a constant equation of state parameter $w$.  If we assume a flat universe ($\\Omega_K=0$), then $\\Omega_{\\rm de}=1-\\Omega_m$, and the only parameters describing the global expansion are $h$, $\\Omega_m$, and $w$. Observations of the primary anisotropies in the cosmic microwave background (CMB) provide two constraints on these three parameters. First, the heights of the acoustic peaks determine the matter density (in g/cm$^3$), which fixes $\\Omega_m h^2$.  Second, the angular scale of the peaks (their location in $l$-space) precisely measures the angular diameter distance to the CMB last-scattering surface, in Mpc.  Absolute distances to low-redshift sources measure the Hubble constant $h$, which then allows all three parameters to be determined \\cite{hujain04,hu04}. Clearly the constraints we present would be substantially degraded if the curvature were not fixed to zero; see \\cite{bernstein05,knox05} for the prospects for precise constraints on curvature.  In addition to the low redshift standard sirens, those at higher redshifts also help constrain dark energy, in the same manner as high-redshift standard candles. \\citet{holz05} discuss how \\lisa\\ observations of SMBBH inspirals can help constrain cosmology. For a dark energy model which is not dramatically different from a cosmological constant $\\Lambda$, the interesting redshift range is when the dark energy density is significant ($z\\alt1$), although note that gravitational lensing degrades the constraints from the highest redshift standard sirens (or candles) ~\\cite{hl05}.  As mentioned above, the gravitational waves from standard sirens measure source distances, but do not measure source redshifts. Some sort of electromagnetic counterpart associated with the merger event will generally be required to use GW sources to determine cosmology. One potential class of GW sources guaranteed to have electromagnetic counterparts are short $\\gamma$-ray bursts (GRBs).  These sources are thought to arise in the mergers of neutron star (NS) binaries, and hence should be strong GW emitters in the frequency band accessible to ground-based observatories. The GRB counterpart to these GW source provides a precise sky localization, which is useful both for determining the redshift to the source, and for significantly improving the GW determination of absolute distance.  As we discuss below, short GRBs occur at a rate large enough for them to provide interesting cosmological constraints. ",
        "conclusions": "We have shown that observations of the gravitational waves emitted by binary compact object inspirals can be a powerful probe of cosmology. In particular, short $\\gamma$-ray bursts appear quite promising as potential GW standard sirens.  The presently observed rate of short GRBs is sufficiently high that within a few years of observation by the next generation of ground-based GW observatories (e.g. \\ligo-II, \\virgo\\ and \\aigo), strong constraints on dark energy parameters may be derived ($\\sigma_w<0.1$).  These inspiraling NS-binaries should be clean sources of gravitational waves; possible sources of contamination, such as tidal effects, magnetic torques, or gasdynamical torques from circumbinary gas, should all be negligible during the crucial inspiral phase (where $v/c \\alt 0.3$).  We emphasize again that the best information about distance measurements comes from combining multiple GW data from instruments that are widely separated.  Good information about the collimation of the gamma rays and thus on the likely inclination of the binary progenitor will also improve the utility of these standard sirens.  Given the great cosmological potential of GW observations of short GRBs, there is strong incentive to extend the lifetime of GRB satellites such as \\textsl{Swift} or \\textsl{HETE}-2 to overlap with next-generation gravitational wave observatories.  The inspirals of SMBBH binaries observed by \\lisa\\ can also provide interesting constraints on dark energy, if the rate of such mergers is high enough to average away noise caused by gravitational lensing.  At present, the total rate and redshift distribution of SMBBH mergers are not well understood, with estimates ranging from a few (or zero) per year, up to hundreds per year, depending upon assumptions \\cite{sesana04,menou01,haehnelt03,koushiappas06}. If the rates are at the high end of these estimates, with a significant fraction at redshifts $z<2$, then $w$ may be constrained at the few percent level."
    },
    "0601/astro-ph0601043_arXiv.txt": {
        "abstract": "It has been proposed that the enrichment in noble gases found by Galileo in Jupiter's atmosphere can be explained by their delivery inside cold planetesimals. We propose instead that this is a sign that the planet formed in a chemically evolved disk and that noble gases were acquired mostly in gaseous form during the planet's envelope capture phase.  We show that the combined settling of grains to the disk midplane in the cold outer layers, the condensation of noble gases onto these grains at temperatures below 20-30K, and the evaporation from high disk altitudes effectively lead to a progressive, moderate enrichment of the disk. The fact that noble gases are vaporized from the grains in the hot inner disk regions (e.g. Jupiter formation region) is not a concern because a negative temperature gradient prevents convection from carrying the species into the evaporating region. We show that the $\\sim 2$ times solar enrichment of Ar, Kr, Xe in Jupiter is hence naturally explained by a continuous growth of the planet governed by viscous diffusion in the protosolar disk in conjunction with an evaporation of the disk and its progressive enrichment on a million years timescale. ",
        "introduction": "Ten years have elapsed since the Galileo probe entered Jupiter's atmosphere and allowed {\\it in situ} measurements of its composition \\citep[e.g.,][]{1996Sci...272..837Y, 1999Natur.402..269O}. One of the key results of the mission is that, when compared to hydrogen, most measurable species appear to be enriched by a factor two to four compared to a solar composition.   Most surprisingly, the list of enriched species includes not only ``ices'' (i.e. species that condense at temperature $\\sim 100\\,$K in the protosolar disk like methane, ammonia), but also the noble gases Ar, Kr, Xe \\citep{2000JGR...10515061M}, which are enriched compared to the new solar abundances \\citep[e.g.][for a review]{2003ApJ...591.1220L} by factors $2.2\\pm 0.8$, $2.0\\pm 0.7$, $2.0\\pm 0.5$, respectively. This led \\citet{1999Natur.402..269O} to propose that Jupiter formed from very low-temperature planetesimals, at a temperature $\\sim 30\\,$K allowing a direct condensation of noble gases onto amorphous ice. This explanation was challenged by \\citet{2001ApJ...550L.227G} who proposed that noble gases would be incorporated into crystalline ice by clathration at slightly higher temperatures, and that they would thus be delivered into Jupiter by icy planetesimals. The large amount of cages available to trap noble gases implies that the ratio of oxygen to other species in Jupiter should be highly non-solar, but the precise factor depends on the condensation sequence \\citep{2001ApJ...550L.227G, 2004P&SS...52..623H, 2005ApJ...622L.145A}.    In this letter, we propose instead that the present abundances of Jupiter indicate that the planet (and other giant planets as well) formed in an chemically evolved protoplanetary disk, in which part of the light gases (hydrogen, helium and maybe neon) had been lost by photoevaporation.  Except as otherwise noted, we will use the disk evolution model described by \\citet{HuGu05} with the following values of the parameters: $M_{\\rm cd}=1\\,\\msol$ (mass in the molecular cloud core), $M_0=0.4\\,\\msol$ (seed mass for the protosun), $T_{\\rm cd}=10\\,$K (ambient interstellar temperature), $\\omega_{\\rm cd}=10^{-14}\\,\\rm s^{-1}$ (assumed uniform rotation rate of the cloud core), $\\alpha=0.01$ (viscosity). ",
        "conclusions": "We have shown that the observed enrichment of Jupiter in noble gases can be explained by a progressive capture of the planet's envelope in sync with the evaporation of the protosolar disk. We infer that the envelope capture phase started relatively late, at about half of the disk's lifetime, probably because of a slow or delayed growth of the protoplanetary core. In our scenario, the final masses of Jupiter and Saturn result naturally from a competition between limited viscous accretion and disk evaporation. We note that a welcome consequence of the late start scenario is a significant suppression of the inward migration when compared to models in which giant planets form early, in locally massive disks (such as the minimum mass solar nebula or more). Provided disk atmospheres can be heated to temperatures of 100\\,K or more, the disk model that we propose further explains the disappearance of disks in Ma timescales, as observed.  We predict that the enrichments in Ar, Kr, Xe in Saturn should be similar to that in Jupiter. It should be larger in Uranus and Neptune, due to a late capture of more enriched gas from the evaporating disk, but also to a pollution of unknown magnitude from their ice cores, significantly more massive than their gaseous envelopes. No isotopic fractionation is expected to arise from the disk evaporation process. However, clues on the ambiant mean temperature of the protosolar molecular cloud core may be obtained from the primordial Ne/Ar, Ne/Kr, and Ne/Xe ratios: if this mean temperature was high enough (e.g. $20\\sim 30$\\,K) not to allow neon to condense onto grains, then neon should have escaped progressively from the disk with hydrogen and helium, yielding small values of these ratios. In all cases these ratios should be equal to or smaller than unity.  These results show the importance of noble gases for tracing back events that occurred in the early Solar System and stress the need for an accurate determination of the compositions of all giant planets."
    },
    "0601/astro-ph0601569_arXiv.txt": {
        "abstract": "We report here on the first pointed X-ray observation of the supernova remnant (SNR) G330.2+1.0. The X-ray morphology is characterized by an extended shell. Its X-ray spectrum is well represented by a single power-law function with a photon index of $\\gamma\\simeq 2.8$ and interstellar absorption of $n_{\\rm H}\\simeq2.6\\times 10^{22}$[cm$^{-2}$]. We interpret this emission as synchrotron radiation from accelerated electrons at the SNR shock, as seen in SN 1006. The surface brightness of the X-ray emission is anti-correlated with the radio emission, and the power-law spectrum is dominated at the western shell where the radio emission is weak.  The co-existence of two distinct (radio bright/X-ray faint and radio faint/X-ray bright) shells in a single supernova remnant challenges our understanding of the particle acceleration and radiation mechanisms in different interstellar environments. The object may be a good target for searching TeV gamma-rays and molecular gas surrounding the blast shock. We also report on the nature of a bright point-like source (AX~J1601$-$5143) to the south of the SNR. ",
        "introduction": "Our understanding of the origin and the acceleration mechanism of cosmic-ray particles has been a great challenge in astrophysics since the last century. X-ray observations have revealed that acceleration is taking place in supernova remnant shells by detecting featureless power-law type spectra (Koyama et al. 1995; Koyama et al. 1997; Slane et al. 2001). These spectra are interpreted as synchrotron radiation from accelerated electrons. The detections of TeV photons in RX~J1713.7$-$3946 (Muraishi et al. 2000; Aharonian et al. 2004) and RX J0852.0$-$4622 (Aharonian et al. 2005) have assured the presence of high-energy ($>$TeV) particles, either leptons or hadrons, in these SNR shells.  The radio source of G330.2+1.0 is characterized by a clumpy distorted shell of \\timeform{11'} diameter (Caswell et al. 1983; Whiteoak, Green 1996). Caswell et al. (1983) attributed its morphology as being due to the interaction of a shock wave with a dense medium, particularly toward the lower galactic latitude (southeast). A clumpiness in the radio map led Whiteoak and Green (1996) to classify G330.2+1.0 as a possible composite type. As we show in the following sections, there is no signature of a Crab-like, compact X-ray source at the position of the clumpy radio shell. We therefore classify G330.2+1.0 as a shell type.  The H\\emissiontype{I} absorption spectrum to the SNR gives a minimum kinematic distance of $4.9\\pm 0.3$ kpc, while the greater distance for velocities interior to the solar circle, 9.9~kpc, can not be excluded (McClure-Griffiths et al. 2001; figure 4 and table 2). A possible association of the SNR with the surrounding H\\emissiontype{I} shell (McClure-Griffiths et al. 2001; subsection 5.3.2 and figure 14) favors that the SNR is close to the absorbing H\\emissiontype{I} clouds at $4.9\\pm0.3$~kpc (or at $9.9$~kpc). It is worth noting that the line of sight of $l=$\\timeform{330D.2} intersects the Norma arm at distances of $\\simeq6.3$~kpc and $\\simeq11.3$~kpc (McClure-Griffiths et al. 2001; figure 14). Therefore, a reasonable distance upper limit to the SNR may be $11.3$~kpc. Hereafter, we denote the SNR distance as $d_{4.9}$ normalized to 4.9~kpc.  For this distance and the radio shell's radius, the Sedov solution's age--radius relation gives a minimum age of $t=3.1\\times 10^{3}\\, (E_{51}/n_{0})^{-1/2}\\, d_{4.9}^{5/2}$ yr, with an explosion energy normalized to $10^{51}$erg and the interstellar medium density normalized to 1\\, ${\\rm cm^{-3}}$. ",
        "conclusions": "Since the absorption columns to G330.2+1.0 and AX~J1601$-$5143 (table \\ref{tab:spectral_parameter}) do not exclude a possibility that they are at a similar distance, we explore if the two objects are physically associated or not.  AX~J1601$-$5143 is probably the same object as 1RXS~J160045.0$-$514307 in the ROSAT all-sky survey faint source catalog (Voges et al. 2000).  The cataloged position is $65''$ from the GIS position. The PSPC countrate, $0.032\\pm0.012$~cps, is as expected from the GIS spectral parameters with either the mekal model or the power-law model, taking into account the statistical uncertainties of the PSPC and GIS data.  If the power-law model is appropriate, AX~J1601$-$5143 could be a Crab-like object or a background active galactic nucleus. The absence of any significant long-term flux variability favors the former interpretation, but the required transverse velocity, $v_{\\rm t} = 3.9\\times 10^{3}(t_{\\rm age}[{\\rm yr}]/3100)^{-1}\\,d_{4.9}\\, {\\rm km\\, s^{-1}}$, for the projected angular distance from the shell center, $\\simeq 8'.8$ or 13\\, $d_{4.9}$~pc, is unreasonably large.  Also, the observed photon index, 2.63, is steeper than any object powered by an energetic pulsar (e.g., Gotthelf 2003).  Thus, it is unlikely that AX~J1601$-$5143 is a Crab-like pulsar/nebula.  If the thermal interpretation is correct, AX~J1601$-$5143 could be a nearby ($d<4.9$~kpc) white dwarf binary with $L_{\\rm X}<5.5\\times 10^{34}\\,d_{4.9}^2\\,{\\rm erg\\,s^{-1}}$. We have not found any clear evidence of line emissions in the GIS spectra, and derive an equivalent width upper limit of $EW<400$~eV for a narrow line at 6.7~keV from helium-like iron. This equivalent width and the elemental abundance (table 1) are not inconsistent with the characteristics found in white-dwarf binaries (e.g., Ezuka, Ishida 1999). The remnant is in the crowded galactic plane, and it is not surprising that a point-like source has been found within the same field of view. We therefore conclude that G330.2+1.0 and AX~J1601$-$5143 are not physically associated, and that G330.2+1.0 is a shell-type SNR without an energetic pulsar.  From a statistical point of view alone, we can not distinguish the better spectral model for the SNR shell (table \\ref{tab:spectral_parameter}). If the thermal model is appropriate, the relatively high temperature ($kT=2.99$~keV) suggests that the SNR is young and emission lines from metal-rich ejecta should be expected.  The absence of significant Mg, Si, S, and Fe K shell lines is incompatible with our knowledge for young, high-temperature SNRs, such as Cas-A, Tycho's, or Kepler's (e.g., Tsunemi et al. 1986; Kinugasa et al. 1999).  The suppression of emission lines could occur if the interstellar medium density is very low, resulting in the low ionization state. This is unlikely, since the clumpy radio morphology suggests shock interaction with the dense medium. Therefore, it is difficult to explain the spectrum in the thermal model. We thus conclude that the most reasonable interpretation for the featureless spectrum in G330.2+1.0 is synchrotron radiation from relativistic electrons accelerated in the SNR shell.  To evaluate multi-wavelength emission, we fit the SNR spectrum with the {\\tt srcut} model, which describes synchrotron radiation from a power-law distribution of electrons with an exponential cut-off (Reynolds, Keohane 1999). For a fixed radio spectral index ($\\alpha=0.3$) and flux density (5 Jy at 1 GHz) (Green 2004), the fit is significantly worse ($\\chi^2/{\\rm dof}=470.6/392$) than that of the simple power-law model. This means that the radio and X-ray emission can not have originated from a simple, single population of electrons in a uniform magnetic field.  The best-fit parameters are the break frequency of $\\nu_{\\rm b} = 4.3\\times 10^{15}$~Hz and the absorption column of $n_{\\rm H}=5.1\\times 10^{22}\\, {\\rm cm^{-2}}$. If the radio flux density is considered as a free parameter, the fit becomes better ($\\chi^2/dof=470.6/392$) with $\\nu_b = 7.0\\times 10^{16}$~Hz, a flux density of $S=7.3\\times 10^{-3}$~Jy, and $n_{\\rm H}=2.2\\times 10^{22}\\, {\\rm cm^{-2}}$. The break frequency corresponds to an electron energy of $E_{\\rm b} = 30 \\cdot (B_5)^{-1/2}$~TeV. Here, $B_5$, is the magnetic field of the emitting region normalized to 5\\, ${\\rm \\mu G}$.  The small best-fit flux density ($7.3\\times 10^{-3}$~Jy at 1\\, GHz) compared to that observed (5 Jy) suggests that there are at least two populations of particles, and only a tiny fraction of the radio-emitting particles has an extended tail toward the TeV region.  As can be seen in figure \\ref{fig1}, the surface brightness of the X-ray emission is anti-correlated with the radio emission.  The X-ray emission is strongest at the western shell, where the lower radio intensity toward the higher galactic latitude suggests a lower ISM density. It is likely that moderately accelerated (GeV) electrons are efficiently decelerated in the eastern shell interacting with the dense ISM. In contrast, the lower ISM density in the western shell results in efficient acceleration to the TeV range.  The co-existence of two distinct (radio bright/X-ray faint and radio faint/X-ray bright) shells in a single supernova remnant gives us an interesting opportunity to study the acceleration mechanisms in different ISM environments.  Interestingly, the X-ray luminosity of the shell, $L_{\\rm SNR}=4.6\\times 10^{34}\\,d_{4.9}^2\\,{\\rm erg\\,s^{-1}}$ (0.7--10~keV) is comparable to that for RX~J1713.7$-$3946 ($1\\times 10^{35}\\,d_{1}^2$), SN~1006 ($\\simeq 2\\times 10^{35}\\,d_{2.2}^2\\, {\\rm erg\\,s^{-1}}$), RX~J0852.0$-$4622 ($2\\times10^{34}\\,d_{1}^2\\, {\\rm erg\\,s^{-1}}$) (Slane et al. 1999; Slane et al. 2001), within a factor of $\\simeq 5$.  There may be common physics at work for tuning the acceleration efficiency in these SNRs.  In summary, we reported on the discovery of a featureless X-ray spectrum in the SNR G330.2+1.0. We interpret this emission as coming from synchrotron radiation of electrons accelerated at the SNR shock as seen in SN~1006 and RX~J1713.7$-$3946. Our preliminary analysis of the archived XMM-Newton data with short exposure ($\\sim 9$~ks) qualitatively confirmed our findings presented herein. Further analysis of the XMM-Newton data with an emphasis on the nature of AX~J1601$-$5143 will be presented elsewhere. Since the surface brightness of the shell is low, X-ray observations with Suzaku, which has a large effective area and low background, will be useful for an accurate determination of the SNR spectrum. Although neither G330.2+1.0 nor the compact source was detected in the recent H.E.S.S. survey (Aharonian et al. 2006), additional deep searches for TeV emission may be useful. For constraining different radiation mechanisms (inverse Compton effect, pion decay, non-thermal bremsstrahlung) in the TeV region, radio observations of the environmental molecular gas (e.g., Fukui et al. 2001) will be of great help."
    },
    "0601/astro-ph0601333_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\setcounter{equation}{0}   Analysis of the Hubble diagram of high redshift Type Ia supernovae (SNe Ia) has discovered that the expansion of the Universe is currently accelerating \\cite{SCP}. In addition, combining measurements of the acoustic peaks in the angular power spectrum of the cosmic microwave background (CMB) anisotropy which indicate the flatness of the Universe \\cite{CMB} and the matter power spectrum of large scale structure (LSS) which is inferred from galaxy redshift surveys like the Sloan Digital Sky Survey (SDSS) \\cite{SDSS} and the $2$-degree Field Galaxy Redshift Survey ($2$dFGRS) \\cite{2dFGRS} has confirmed that a component with negative pressure (dark energy) should be added to the matter component to make up the critical density today.  The cosmological constant and/or a quintessence field are the most commonly accepted candidates for dark energy. The latter is a dynamical scalar field leading to a time dependent equation of state, $\\omega_{\\phi}$. Also, this scalar field has a fluctuating, inhomogeneous component in order to conserve the equivalence principle corresponding to the response of the new component to the inhomogeneities in the surrounding cosmological fluid \\cite{CDS}.  Several new observational effects produced by the existence of quintessence are imprinted in the CMB anisotropies and the matter power spectrum when compared to models with the cosmological constant. The locations of the acoustic peaks in the CMB angular power spectrum are shifted due to their dependence on the amount of dark energy today and at last scattering as well as $\\omega_{\\phi}$ \\cite{Doran, SLee}. Usually the Universe is dominated by the quintessence at late times ($z \\sim {\\cal O}(1)$), when the gravitational potential associated with the density perturbations is changed due to a time dependent $\\omega_{\\phi}$. This enhances the CMB anisotropies at large angular scales by the integrated Sachs-Wolfe effect (ISW) \\cite{ISW}. Thus the amplitudes of both CMB angular and matter power spectra decrease at large scale for the fixed Cosmic Background Explorer Satellite (COBE) normalization compared to those in the cold dark matter model with a cosmological constant ($\\Lambda$CDM)~\\cite{Dodelson1}.  The possibility that a scalar field at early cosmological times follows an attractor-type solution and tracks the evolution of the visible matter-energy density has been explored \\cite{Ratra}. This may help alleviate the severe fine-tuning associated with the cosmological constant problem. However, this still cannot explain the reason why the dark energy and the dark matter have comparable energy densities at present. Recently models considering the coupling of quintessence to dark matter have been investigated as a possible solution for this late time coincidence problem \\cite{coupQ}. However, several of these models using a simple coupling can be ruled out by observational constraints \\cite{coupQ1}.  These non-minimally coupled quintessence models have several different observational effects compared to the minimally coupled models. One of the most important effects is a different scaling of the cold dark matter (CDM) compared to that of CDM of the minimally coupled case. Since CDM scales as $\\rho_{c} = \\rho_{c}^{(0)} a^{-3 + \\xi}$ where $\\xi < 0$, there will be more CDM energy density at early epoch compared to the case with $\\xi = 0$ ({\\it i.e.} $n _c =0$). As the coupling is increased, the locations of the acoustic peaks are also shifted to smaller scales. However, the amplitudes of the odd-number peaks decrease due to the decrease of the baryon density at early time and the increasing ISW effect as the coupling scaled with the COBE normalization. The amplitudes of the even-number peaks increase due to the increase of the CDM energy density. Even though the ratios of the height of the first peak to those of higher peaks between different couplings are quite similar to one another, these are quite different from the $\\Lambda$CDM model \\cite{KMT}. The location of the turnover in the matter power spectrum corresponds to the scale that entered the Hubble radius when the universe became matter-dominated ($a_{eq}$). Thus the coupling might shift the turnover scale.  This paper is organized as follows. In the next section we show the basic equations of linear perturbations of the coupled quintessence model. In Sec. 3, we derive the formal entropy perturbation due to multiple fluids. We also consider the possibility of isocurvature perturbations of the quintessence. We check the coupling effects on CMB and matter power spectrum in Sec. 4. The effect of coupling on the metric perturbation is considered in Sec. 5. Our conclusion is in the last section. ",
        "conclusions": "\\setcounter{equation}{0}  We have analyzed the linear perturbations of the cosmological models for the scalar field with its self-interaction potential, $V(\\phi) = V_{0} \\exp(\\lambda \\phi^2 /2)$, and its coupling to the CDM, $\\exp[B_{c}(\\phi)] = \\exp(n_{c} \\lambda \\phi^2 /2)$. The evolution of the non-perturbed background scalar field $\\bar{\\phi}$ occurred in the tracking regime throughout the radiation dominated epoch. The full analysis of the non-perturbed and perturbed equation of each species in the conformal Newtonian gauge has been done including the proper energy-momentum transfer vector due to the coupling, $Q_{(d) \\nu}$. The Boltzmann equations have been modified as to account for the coupling between a scalar field and the CDM.  We have seen that the energy-momentum of each species may not be conserved as a result of the coupling of the scalar field to the CDM. Thus we have redefined the concepts of entropy perturbations. We have shown that the isocurvature perturbation of the quintessence has been damped to zero with time in the tracking regime during the radiation epoch. Thus we have constrained our considerations on the adiabatic perturbation.  We have considered the CMB anisotropy spectrum and the matter power spectrum for the non-minimally coupled models. Additional mass and source terms in the Boltzmann equations induced by the coupling give the rapid changes of the Newtonian potential $\\Phi$ and enhance the ISW effect in the CMB power spectrum. The modification of the evolution of the CDM, $\\rho_{c} = \\rho_{c}^{(0)} a^{-3 + \\xi}$, changes the energy density contrast of the CDM at early epoch. We have adopted the current cosmological parameters measured by WMAP within $1 \\sigma$ level. With the COBE normalization and the WMAP data we have found the constraint of the coupling $n_{c} \\leq 0.01$. The locations and the heights of the CMB anisotropy peaks have been changed due to the coupling. Especially, there is a significant difference for the heights of  the second and the third peaks among the models. Thus upcoming observations continuing to focus on resolving the higher peaks may constrain the strength of the coupling. The suppression of the amplitudes of the matter power spectra could be lifted by a bias factor. However, a detailed fitting is beyond the scope of this paper. The turnover scale of the matter power spectrum may be used to constrain the strength of the coupling $n_{c}$.  Finally, we have investigated the metric perturbations including the coupling between the scalar field and the CDM. There is no difference to the curvature perturbation $\\zeta$ for the different couplings in the adiabatic case. However, the effects of the coupling have been absorbed in the Boltzmann equations already."
    },
    "0601/astro-ph0601619_arXiv.txt": {
        "abstract": "Bright optical and X-ray flares have been observed in many Gamma-ray Burst (GRB) afterglows. These flares have been attributed to  late activity of the central engine. In most cases the peak energy is not known and it is possible and even likely that there is a significant far-ultraviolet component. These far-ultraviolet photons escape our detection because they are absorbed by the neutral hydrogen before reaching Earth. However, these photons cross the blast wave produced by the ejecta that have powered the initial GRB. They can be inverse Compton upscattered by hot electrons within this blast wave. This process will produce a strong sub-GeV flare that follows the high energy (soft X-ray) tail of the far-UV flare but lasts much longer and can be detected by the upcoming {\\em Gamma-Ray Large Area Telescope} (GLAST) satellite. This signature can be used to probe the spectrum of the underlying far-ultraviolet flare. The extra cooling produced by this inverse Compton process can lower the X-ray emissivity of the forward shock and explain the unexpected low early X-ray flux seen in many GRBs. ",
        "introduction": "\\label{sec:MevGevInt}  The XRT on board of {\\it Swift} detected, during the last year, numerous X-ray flares in GRB afterglows (Burrows et al. 2005; Nousek et al. 2006; Goad et al. 2006; Romano et al. 2006; Falcone et al. 2006; O'Brien et al. 2006). These observations confirmed earlier findings of BeppoSAX (Piro et al. 1998, 2005; Galli \\& Piro 2006; in't Zand et al. 2004) and ASCA (Yoshida et al., 1999). These flares have been interpreted as arising from late time activity of the central engine (King et al. 2005; Perna et al. 2006 and Proga \\& Zhang 2006) producing either internal shocks (Fan \\& Wei 2005; Zhang et al. 2006; Zou, Xu \\& Dai 2006; Wu et al. 2006) or internal magnetic dissipation (Fan, Zhang \\& Proga 2005a).  The flare detected in the afterglow of GRB 050502b peaks in the soft X-rays (Falcone et al. 2006). However, the peak energy of most flares is unknown (B. Zhang, 2006, private communication). It is possible, and even likely, that a significant fraction of the energy or even most of it is emitted in the far-ultraviolet (FUV) band. For example in the internal energy dissipation model the typical synchrotron radiation frequency depends sensitively on the physical parameters (Fan \\& Wei 2005; Zhang et al. 2006; Fan et al. 2005a) and the synchrotron self-absorption frequency is $\\sim{\\rm 10^{15}~Hz}$ (Fan \\& Wei 2005). It is possible, therefore, that the observed X-ray flares are the high energy tails of FUV flares. It is also possible that there are FUV flares that have not been detected at all. Further support for this idea arises from the possible interpretation of the optical flare seen in GRB 050904 (B\\\"oer et al. 2006) as a late time activity of the inner engine (Wei, Yan \\& Fan 2006).  Even if FUV flares exist they won't be observed as the FUV photons are absorbed by the neutral hydrogen in the GRB host galaxy as well as in our Galaxy.  We show here that an underlying FUV flare will be upscattered and produce (after inverse Compton) a sub-GeV flare that may be detected by the upcoming {\\em Gamma-Ray Large Area Telescope} (GLAST; see http://glast.gsfc.nasa.gov/) satellite. Though (sub-)GeV flashes in GRB afterglows could arise in other scenarios (e.g., M\\'esz\\'aros \\& Rees 1994; Plaga 1995; Granot \\& Guetta 2003; Dermer \\& Atoyan 2004; Beleborodov 2005; Fan, Zhang \\& Wei 2005b), the high energy photon flashes predicted in this Letter can be distinguished easily since they follows the high energy (soft X-ray) tail of the FUV flares (see also Wang, Li \\& M\\'esz\\'aros 2006), which is unexpected in any alternative model. Our prediction could be tested by the cooperation of {\\it Swift} and GLAST. This is possible as {\\it Swift} XRT usually slews to the GRB source in $\\sim 100$ seconds and, the field of view of the GLAST burst monitor (GBM) is all sky not occulted by the earth and the Large Area Telescope (LAT) will slew to the GRB direction automatically in $\\sim 5$ minutes. ",
        "conclusions": "\\label{sec:MG_Dis} Bright X-ray flares have been detected a large group of {\\it Swift} GRB afterglows. These flares have been attributed to late activity of the central engine. In most cases the peak energy is not known and it is possible that there is a significant far-ultraviolet component. These far-ultraviolet photons escape our detection because they are absorbed by the neutral hydrogen both in host galaxy and in our Galaxy  before reaching Earth. We suggest here that these far-ultraviolet photons are IC upscattered by hot electrons within the shock that is ahead of them. This shock is driven by the blast wave produced by the ejecta that powered the initial GRB. This IC process will produce a strong sub-GeV burst what can be detected by upcoming {\\em Gamma-Ray Large Area Telescope} (GLAST) satellite if the far-ultraviolet/X-ray flare is bright enough (the energy fluence ${\\cal F}\\sim 10^{-6}~{\\rm ~erg~cm^{-2}}$). Alternatively, if most flare photons are in keV band rather than in far-ultraviolet band, the total number of photons being scattered (and the detected number of photons) is much smaller though the typical energy is much higher (Generally, they are in GeV-TeV energy range, see Wang et al. 2006). It is not easy to collect enough photons for significant detection.  We have also analyzed the sub-GeV detections of EGRET in view of this model. In some events (for example, GRB 930131, GRB 940217 and GRB 941017), the sub-GeV photons are delayed. Out of these bursts, the spectrum of GRB 941017 is rather hard and it cannot be explained by this model. On the other hand the other two events seem to be consistent with the model. However, these sub-GeV photons could be generated in other scenarios (e.g., M\\'esz\\'aros \\& Rees 1994 Plaga 1995; Dermer \\& Atoyan 2004; Fan et al. 2005b). Without simultaneous soft X-ray observations we cannot confirm our model.  Finally, we point out that the extra cooling induced by continuous flux of FUV photons that pass through the forward shock reduces the X-ray flux produces by this shock front. Thus if the central engine does not turn off and gives out a  significant emission of FUV photons several thousand seconds after the GRB, this emission will reduce the X-ray flux emitted by the forward shock at the beginning of the afterglow phase. This possibility provides an alternative explanation to the weak and slowly declining early X-ray light curve observed by {\\it Swift} in many GRBs (Nousek et al. 2006). As this process involves emission of sub-GeV photons it could be tested by simultaneous observations of {\\it Swift} and GLAST in the near future."
    },
    "0601/astro-ph0601105_arXiv.txt": {
        "abstract": "We present an X-ray tour of diffuse emission in the 30~Doradus star-forming complex in the Large Magellanic Cloud using high-spatial-resolution X-ray images and spatially-resolved spectra obtained with the Advanced CCD Imaging Spectrometer aboard the {\\em Chandra X-ray Observatory}.  The dominant X-ray feature of the 30~Doradus nebula is the intricate network of diffuse emission generated by interacting stellar winds and supernovae working together to create vast superbubbles filled with hot plasma.  We construct maps of the region showing variations in plasma temperature (T = 3--9 million degrees), absorption ($N_H$ = 1--6$\\times 10^{21}$~cm$^{-2}$), and absorption-corrected X-ray surface brightness ($S_X$ = 3--126$\\times 10^{31}$~ergs~s$^{-1}$~pc$^{-2}$).  Enhanced images reveal the pulsar wind nebula in the composite supernova remnant N157B and the {\\em Chandra} data show spectral evolution from non-thermal synchrotron emission in the N157B core to a thermal plasma in its outer regions.  In a companion paper we show that R136, the central massive star cluster, is resolved at the arcsecond level into almost 100 X-ray sources.  Through X-ray studies of 30~Doradus the complete life cycle of such a massive stellar cluster can be revealed. ",
        "introduction": "}  Images of spiral and irregular galaxies have been defined historically by regions of massive star formation, signposts demarcating important features such as spiral arms, bars, and starbursts.  They remind us that galaxies really are evolving, with continuous injection of energy and processed material into the galaxy.  Massive star-forming regions (MSFRs) present us with a microcosm of starburst astrophysics, where stellar winds from O and Wolf-Rayet (WR) stars compete with supernovae (SNe) to carve up the neutral medium from which they formed, igniting a new generation of stars.  With X-ray observations, we see the sources that shape the larger view of a galaxy, injecting energy and processed material into the disk and halo ISM through winds, ionization fronts, SNe, superbubbles, and chimneys.  About 80\\% of all SNe occur in superbubbles, so cosmic rays generated by SNe also may be accelerated in superbubbles \\citep[e.g.][]{Higdon98, Parizot04}.  X-ray observations probe different energetic components than traditional visual and infrared (IR) studies, penetrating the obscuring material of the natal cloud while minimizing confusion from foreground and background objects.  X-ray studies also detect the presence of past SNe through the shocks in their extended remnants.  Some superbubbles are well-known bright X-ray sources, where multiple SNe from past OB stars produce soft X-rays filling bubbles 50--100~pc in size with $L_x \\sim 10^{35-36}$~ergs~s$^{-1}$ \\citep{Chu90}, while others (not recently brightened by such ``cavity'' SNe) are X-ray faint \\citep{Chu95}. MSFRs thus exhibit a complicated mixture of point and extended structures that are easily confused by low-resolution X-ray telescopes.  30~Doradus (30~Dor) is a classic example of this.  The {\\em Chandra X-ray Observatory} \\citep[Figure~\\ref{fig:csmoothimage} and][hereafter Paper~II]{Townsley05b} and {\\em XMM-Newton} \\citep{Dennerl01} share the highest-quality views yet achieved of this extraordinary and complicated X-ray field.  30~Dor is a field of superlatives.  Located in the Large Magellanic Cloud (LMC), it is the most luminous Giant Extragalactic H{\\sc II} Region (GEHR) and ``starburst cluster'' in the Local Group.  It hosts the massive compact star cluster R136 \\citep[][called ``RMC~136'' in SIMBAD]{Feast60}, a testbed for understanding recent and ongoing star formation in the 30~Dor complex \\citep{Walborn02}.  Although R136, with a mass of $\\sim 6 \\times 10^4$~M$_{\\odot}$ \\citep{Brandl05}, does not quite qualify as a ``super star cluster'' \\citep{Meurer95}, it has more than 50 times the ionizing radiation of the Orion Nebula.  Nearby, one finds the large supernova remnant (SNR) N157B (30~Dor~B), embedded in an H{\\sc II} region generated by the OB association LH~99 \\citep{Chu92} and containing an X-ray pulsar \\citep{Marshall98}.  {\\em ROSAT} images showed two fainter sources associated with WR stars near the central star cluster, at least five plasma-filled superbubbles, and the N157B SNR \\citep{Wang95, Wang99}.  30~Dor's proximity and complexity provide us with a unique and important microscope into the starburst phenomenon in galaxies.  Figure~\\ref{fig:csmoothimage} illustrates the spatial and spectral complexity of the X-ray emission from 30~Dor with a smoothed image of our {\\em Chandra} observation obtained with the Advanced CCD Imaging Spectrometer (ACIS).  The aimpoint of the $17\\arcmin \\times 17\\arcmin$ ACIS Imaging Array (ACIS-I) was the R136 cluster.  At $D=50$~kpc, the distance we assume throughout this paper and Paper~II, $1^{\\prime} \\sim 14.5$~pc, so ACIS-I covers $\\sim 240$~pc $\\times$ 240~pc.  Two off-axis CCDs from the ACIS Spectroscopy Array (ACIS-S) were operated for this observation, adding a $\\sim 120$~pc $\\times$ 240~pc image of nearby interesting structures (see Figure~\\ref{fig:introimage}), albeit with poor spatial resolution.  A more finely binned image of the ACIS-I data, highlighting the X-ray point sources and showing a J2000 coordinate grid, can be found in Figure~1 of Paper~II.  30~Dor lies at the confluence of two LMC supergiant shells and just the main nebula is $\\sim 250$~pc in diameter.  Since the LMC is viewed nearly face-on, confusion between 30~Dor and other disk structures is minimized.  SNe pervade the region but most go undetected due to age and environment \\citep{Chu90}.  Nearby are two pulsars (B0540$-$69.3 in the SNR N158A and J0537$-$6910 in the SNR N157B) and SN1987A (Figure~\\ref{fig:introimage}).  Wide-field ground-based H$\\alpha$ images \\citep[e.g.][]{Meaburn84,Chu90,Wang99}, enhanced by recent {\\em HST} \\citep{Walborn02} and {\\em Spitzer} data\\footnote{\\url{http://www.spitzer.caltech.edu/Media/releases/ssc2004-01/release.shtml}}, show that the combined actions of stellar winds and SNe from several generations of high-mass stars in 30~Dor have carved its ISM into an amazing display of arcs, shells, pillars, voids, and bubbles, ranging over spatial scales of 1--100~pc \\citep{Brandl05}.  This field provides a unique bridge between Galactic giant H{\\sc II} regions (e.g.\\ W49A, W51A, NGC~3603) and GEHRs \\citep[e.g.\\ NGC~5471 and other regions in M101,][]{Chen05}.  30~Dor is the result of multiple epochs of high-mass star formation in a vast molecular cloud complex chemically enriched by many SNe; it has no Galactic counterpart in terms of mass or complexity of its star formation history and is not even equalled by other GEHRs in the Local Group.  It provides us with a unique view of the most fundamental building block of the starburst phenomenon in galaxies; currently we are not able to resolve comparable star-forming complexes in starburst galaxies, only those many times larger and more powerful \\citep[e.g.\\ NGC~253, NGC~7714/7715;][]{Strickland02,Smith05}.  In \\S\\ref{sec:background} we provide a brief summary of X-ray observations of the 30~Dor complex.  Our {\\em Chandra} observations and data analysis are outlined in \\S\\ref{sec:observation}.  An overview of the X-ray morphology and global spectral properties of the complex is presented in \\S\\ref{sec:global}.  The rest of the paper provides more in-depth studies of specific diffuse X-ray components of the 30~Dor complex:  the superbubbles centered on R136 (\\S\\ref{sec:superbubbles}), the SNR N157B (\\S\\ref{sec:N157B}), and other diffuse structures, with a comparison to the H$\\alpha$ kinematic study of 30~Dor by \\citet{Chu94} (\\S\\ref{sec:otherdiffuse}).  We conclude with a comparison of the X-ray data to recent H$\\alpha$ and IR data and a summary of our results.  The X-ray point source population of 30~Dor revealed by this dataset is described in Paper~II. ",
        "conclusions": "}  \\subsection{The Relationship between Hot, Warm, and Cool Interstellar Material in 30~Dor \\label{sec:hotcold}}  The morphology of the hot X-ray emitting plasma is extremely complicated and bears little resemblance to theoretical calculations of individual SNRs or superbubbles, which generally predict quasi-spherical structures.  Insight into the origins of these structures emerges from comparison with the distribution of interstellar material traced by H$\\alpha$ emission from ionized gas and IR emission from dust.  The comparison with H$\\alpha$ has been made since the earliest X-ray images of 30~Dor were obtained by {\\em Einstein} \\citep[e.g.][]{Chu90,Walborn91,Wang91a}.  Figure~\\ref{fig:Halpha} shows a high-resolution H$\\alpha$ image from the MCELS project \\citep{Smith00} in red, with adaptively-smoothed soft-band ACIS images (point sources removed) in green and blue.  This image is reminiscent of earlier combinations of H$\\alpha$ and X-ray data that showed that diffuse soft X-ray emission was anticorrelated with H$\\alpha$ emission, often filling the cavities outlined by ionized gas \\citep[e.g.][]{Wang99}.  These high-resolution images allow us to refine that picture somewhat, giving more detailed information on the relative locations of harder and softer X-rays and H$\\alpha$ emission.  Comparing Figure~\\ref{fig:Halpha} to a similar image based on {\\em ROSAT}/HRI data in \\citet[][Figure 3]{Wang99}, both images show that the northern Shell 5 and southern Shell 2 are very clearly outlined by large-scale H$\\alpha$ structures.  X-rays from the northern part of SNR N157B coincide with a dark cavity in the H$\\alpha$ emission caused by a dust lane \\citep{Chu92}.  Conversely, the outer loop of Shell 3, labeled as X-ray emitting region 18 in Figure~\\ref{fig:diffregions}, coincides with a similar loop in H$\\alpha$; in Figure~\\ref{fig:Halpha} we can see that the northern part of this loop is softer than the southern part.  Both images show that a prominent H$\\alpha$ dark region in Shell 2 (southeast of R136) contains faint X-ray emission, while the bright H$\\alpha$ filaments that outline it bifurcate our X-ray void labeled region 15.  Figure~\\ref{fig:Halpha} shows that Shell 2 X-rays are relatively hard and Figure~\\ref{fig:nhktmaps} reveals that this is due to a combination of intrinsic spectral hardness and heavy obscuration in this part of 30~Dor.  The West Ring is now more clearly seen to be a clumpy ring of X-ray emission, fainter in its center, with some soft X-ray emission coincident with the Hodge 301 cluster on the northeast edge of the ring.  Small-scale clumps (such as our regions 7 and 8) are now distinct from the smoother underlying X-ray emission. With {\\em Chandra's} high on-axis spatial resolution, we have been able to excise the point source X-ray emission from Figure~\\ref{fig:Halpha}, so it is clearer that bright diffuse regions 5 and 6 are distinct from the R136 cluster and from the nearby WR stars R139, R140, and R145.  The 8$\\mu$m {\\em Spitzer} data (Figure~\\ref{fig:spitzer}) highlighting warm dust also add insight into the diffuse X-ray structures.  Hot X-ray plasma fills the interiors of superbubbles that are outlined by warm dust and emission from PAHs (Brandl et al.\\ 2006, in preparation).  The X-ray morphology and possible confinement are more fully appreciated when anchored by these IR data.  Heated dust provides an envelope to the base of the cylindrical X-ray Shell 5.  X-ray emission could be suppressed in this region or it could be present but absorbed.  An IR-bright V-shaped ridge of emission (reminiscent of the Carina Nebula) separating Shell 3 from Shell 5 is devoid of observable soft X-ray plasma, while the bright X-ray spot dominating Shell 1 is nearly devoid of IR emission.  The thermal plasma that constitutes the outer regions of the N157B SNR fills a large cavity in the warm dust. One of the brightest clumps of diffuse X-rays (our region 5) fills a distinct hole in both the H$\\alpha$ and IR emission.  To give a more complete view of the wide range of emission from 30~Dor, Figure~\\ref{fig:3bands} combines the 8$\\mu$m data tracing PAH emission and warm dust (red) from Figure~\\ref{fig:spitzer}, the H$\\alpha$ data tracing ionized $10^4$~K gas (green) from Figure~\\ref{fig:Halpha}, and the 1120--2320~eV X-ray data tracing $10^7$~K plasma (blue) from Figure~\\ref{fig:csmoothimage}.  Note that here the X-ray image comes from the ACIS data processed with {\\em csmooth}, with the point sources left in place, not the adaptively-smoothed images used in Figures~\\ref{fig:Halpha} and \\ref{fig:spitzer} where the point sources were removed.  Figure~\\ref{fig:3bands} shows the full ACIS-I field of view and is scaled to show the full extent of faint diffuse X-ray emission across the field.  The zoomed image in Figure~\\ref{fig:3bandszoom} is scaled to show just the brighter patches of X-ray emission in the center of the field and removes the color saturation in the core.  Although the {\\em Chandra} data are not deep enough to match the spatial resolution seen in the H$\\alpha$ and {\\em Spitzer} data, it is clear that the 30~Dor complex cannot be understood by visual and IR studies alone, no matter how high the quality of those datasets.  The high-energy emission is a near-perfect complement to the longer-wavelength emission, filling cavities in the complex that are outlined by H{\\sc II} regions.  These are in turn outlined by warm dust, because ultraviolet radiation in the H{\\sc II} regions destroys PAHs.  X-rays from the northern parts of SNR N157B fill a prominent hole in the 8$\\mu$m emission.  The eastern side of the Shell 5 X-ray emission, which appeared unconfined by warm dust in Figure~\\ref{fig:spitzer}, is clearly defined by a large region of H$\\alpha$ emission, while its western side shows a narrower H{\\sc II} region bordered by dust.  The southwestern side of the 30~Dor complex shows more warm dust than the northern or eastern sides.  There is a notable absence of X-ray emission in the southeastern corner of the image, where the stellar cluster SL~639 \\citep{Shapley63,Melnick87,Bica99} is clearly seen in the H$\\alpha$ and 8$\\mu$m data.  Some ionized gas is present around the cluster along with substantial amounts of heated dust.  This cluster may be too young to have produced a substantial wind-blown bubble or any supernovae, or its ionizing stars may not generate enough wind power to blow a substantial bubble.  Conversely, if it is older than $\\sim$10~Myr it may have dispersed its hot gas.  It will be an interesting site to search for faint diffuse X-rays in longer observations.  The famous central arcs north and west of R136 are bright at 8$\\mu$m as well as in H$\\alpha$, illustrating the transition layers from cold molecular material to heated dust to ionized gas that characterize much of the 30~Dor complex.  Figure~\\ref{fig:3bandszoom} shows that the bright X-ray emission lies in the cavities interior to these ionization fronts, but the saturated center of Figure~\\ref{fig:3bands} suggests that some diffuse X-rays are seen superposed on the H$\\alpha$ and IR emission.  This is true in other parts of the image as well and implies that the soft X-rays come from regions that lie in front of the denser IR-emitting material since they avoided being absorbed.  This is a reminder that the complex ISM in 30~Dor may be absorbing similar soft X-rays along other lines of sight.  Thus the regions exhibiting hot plasma may be connected via tunnels or fissures that are not visible to us in these images.  The regions where we do not see soft X-ray emission are not necessarily lacking in hot plasma, especially if it is clear from the {\\em Spitzer} data that substantial absorbing material lies along the line of sight.  These data thus augment and support the ideas developed over the last twenty years for the evolution of the 30~Dor complex.  Powerful stellar winds from extremely massive stars are carving holes in an extensive GMC and filling those holes with hot plasma.  Additional hot plasma is added as those stars explode inside the cavities they created.  In this process, cold gas is pushed aside and confines the hot plasma into shapes otherwise difficult to understand.  Since this plasma appears to emit only soft X-rays for most of its lifetime, some of it may be absorbed by intervening molecular material.  Tunnels and fissures could exist in the GMC that allow hot plasma to flow between apparently unconnected regions, although none of the hot plasma seen by {\\em Chandra} appears completely unconfined (there are no empty bubbles that appear to have vented their X-ray plasma into the surrounding ISM). The H$\\alpha$ emission traces the ionization fronts, H{\\sc II} region ``sheets'' that mark the transition between the million-degree plasma and the neutral material.  This transition layer is heated, evaporated, and is accelerating away from the cloud interface.  Evolved stars are distributed across the entire 30~Dor nebula, providing the fuel for the SNRs that occasionally brighten the superbubbles in X-rays and the enhanced density that allows the plasma to cool.  Neutral clumps may also enhance the X-ray emission, especially in regions close to the remains of the GMC.   \\subsection{Summary of Findings}  Diffuse emission dominates the morphology of 30~Dor in soft X-rays, caused by hot plasma filling the large superbubbles created by generations of star formation and subsequent SNe in this region.  This is dramatically illustrated by combining the {\\em Chandra} data with H$\\alpha$ and {\\em Spitzer} images (Figures~\\ref{fig:Halpha}--\\ref{fig:3bandszoom}); the $10^6$--$10^7$~K X-ray plasmas are enveloped and probably confined in most cases by the cooler gas and warm dust that define the classic picture of 30~Dor. This example illustrates the need for high-quality X-ray observations of star-forming regions; new facets of both stellar and diffuse components are revealed by high-energy data.  The combined spatial and spectral resolution afforded by ACIS shows that these superbubbles are not uniformly filled with a single-temperature gas; great variety is seen, on a range of spatial scales, in absorption, plasma temperature, and intrinsic surface brightness (Figure~\\ref{fig:nhktmaps}).  Some bubbles are center-filled, perhaps revealing interactions with cold gas left in shell interiors \\citep{Arthur96,Wang99}, while others are edge-brightened with distinct central voids not caused by obscuring foreground material.  The faintest regions (voids and the periphery of the main 30~Dor nebula) are a factor of 20 fainter in intrinsic surface brightness than the brightest regions, which have surface brightnesses $>1 \\times 10^{33}$~ergs~s$^{-1}$~pc$^{-2}$.  Our diffuse regions range in size from $1\\arcmin$ ($\\sim 14.5$~pc) for small surface brightness enhancements to $>7\\arcmin$ for the large shells.  For comparison with other GEHRs, we have provided integrated spectra and fits for the main 30~Dor nebula and for the N157B SNR (Table~\\ref{tbl:globalspectable}).  Spectral fits to the diffuse emission show moderate absorption ($N_H = 1$--$6 \\times 10^{21}$~cm$^{-2}$) and soft thermal plasmas with $kT = 0.3$--0.8~keV (3--9~MK).  Many diffuse regions exhibit elevated abundances (above the nominal value of $0.3Z_{\\odot}$), perhaps indicative of line emission, but none show prominent emission lines. Although not a definitive demonstration of a supernova origin for the X-ray emission, these spectral fit results are consistent with that interpretation.  SNR N157B, also possibly the result of an explosion within a pre-existing cavity, also lacks strong emission lines, as does the Galactic example of a possible cavity SNR discussed here, HB3.  The brightest point source in the field is PSR~J0537$-$6910 in the SNR N157B, imaged almost $7\\arcmin$ off-axis.  Although the large PSF there prevents us from performing accurate spectral analysis of the pulsar and its pulsar wind nebula, these data clearly show spectral changes as a function of distance from the pulsar; the pure non-thermal spectrum steepens and combines with a thermal component, finally giving over to a pure thermal spectrum in the outer regions of the SNR (\\S\\ref{sec:N157B}).  Several interesting smaller-scale structures emerge with the high-resolution {\\em Chandra} observations.  One feature (the ``West Ring'') may be the relic signature of a past cavity supernova that exploded inside a pre-existing wind-blown bubble generated by the Hodge~301 cluster.  As first shown by CK94, bright patches of X-ray emission often coincide with high-velocity features known from H$\\alpha$ echelle studies, indicating that X-rays are produced in shocks in 30~Dor's ISM (\\S\\ref{sec:otherdiffuse}).  The MCELS H$\\alpha$ image, the {\\em Spitzer} image, and the recent mosaic of {\\em HST} data \\citep{Walborn02} are all dominated by highly-structured arcs and shells on 1--10~pc scales; in the central region, these reveal the interfaces between the central cavity created by R136 and the surrounding molecular clouds \\citep{Scowen98} and demonstrate how R136 is shredding its natal environment (\\S\\ref{sec:hotcold}).  Here and throughout the 30~Dor nebula, an appreciation of the X-ray emission is necessary to further our understanding of the processes working to shape this complex.  As the images in \\S\\ref{sec:hotcold} show, it is imperative that the diffuse X-ray emission become part of the ``classic'' picture of 30~Dor.  New stellar clusters are now forming in the dense knots that remain in 30~Dor's GMC; their collapse was probably triggered by R136 \\citep{Walborn02}.  These clusters are not yet resolved in X-rays, demonstrating that there is much more work to be done to achieve a complete picture of the X-ray emission from 30~Dor.    \\subsection{30~Dor and Other Massive Star-forming Regions \\label{sec:othermsfrs}}  30~Dor's diffuse X-ray emission is, overall, substantially different than that seen in Galactic MSFRs that are too young to have produced SNe.  M~17 shows strong diffuse soft X-rays probably due to wind-wind and/or wind-cloud collisions \\citep{Townsley03}, but the primary component of this emission is hotter (kT = 0.6~keV) than many 30~Dor regions and has low surface brightness ($10^{31.9}$~ergs~s$^{-1}$~pc$^{-2}$).  These quantities are most similar to the regions we call ``voids'' in 30~Dor (regions 15, 19, 20, 21, and 22 in Figure~\\ref{fig:diffregions}).  The diffuse X-ray emission reported for the massive Galactic clusters NGC~3603 \\citep{Moffat02} and Arches \\citep{Yusef-Zadeh02} has higher surface brightness more comparable to the brighter regions in 30~Dor \\citep[$10^{32.6}$ and $10^{33.1}$~ergs~s$^{-1}$~pc$^{-2}$ respectively, calculated from Table 4 of][]{Townsley03} but the plasmas are much hotter (kT = 3.1~keV and 5.7~keV respectively).  Either a very different mechanism is generating hard diffuse emission in these clusters \\citep[e.g. wind collisions,][]{Canto00}, young SNe dominate the diffuse emission, or substantial unresolved point source emission is corrupting the measurements.  \\citet{Chu93} proposed that the soft diffuse X-ray emission seen in the Carina star-forming complex \\citep{Seward82} could be due to a cavity SNR inside a superbubble blown by Carina's many massive stellar clusters; thus the Carina complex should be a Galactic analog to one of the superbubbles seen in 30~Dor.  The entire Carina complex shows integrated X-ray emission with $kT \\sim 0.8$~keV and surface brightness $\\geq 10^{32.2}$~ergs~s$^{-1}$~pc$^{-2}$ \\citep{Seward82,Townsley03}. {\\em Chandra} and {\\em XMM} resolve this X-ray emission into thousands of harder but comparatively faint stellar sources \\citep[e.g.][]{Albacete03,Evans03} and bright diffuse emission pervading the complex yet not centered on the O or WR stars or the stellar clusters.  No obvious SNR is present, but the complex is old enough to have produced SNe, as evidenced by the presence of evolved stars.  This bright, soft diffuse emission is quite comparable in surface brightness and extent to the smaller regions sampled in 30~Dor.  Using new {\\em Chandra} data centered on the Trumpler~14 OB cluster in Carina, we show that the brightest regions of diffuse emission are well-separated from the massive stars in the Carina complex and show filamentary morphology, both consistent with a cavity supernova origin \\citep{TownsleyIAUS227}.  Thus we suspect that Carina serves as a good microscope for understanding the processes powering the X-ray emission from 30~Dor.  Given the morphological complexity that 30~Dor displays at all wavelengths, it is surprisingly similar to other GEHRs, most notably NGC~604 in M33, the second-largest GEHR in the Local Group.  Although not dominated by a massive central star cluster, the large-scale loops and voids of NGC~604 are also characteristic of superbubbles and are filled with hot gas emitting soft X-rays \\citep[][Brandl et al.\\ 2006, in preparation]{Maiz04}.  The second-largest GEHR in the LMC, N11, also exhibits a central massive stellar cluster surrounded by a superbubble and containing bright, soft diffuse X-rays \\citep{MacLow98,Naze04}. Although more distant GEHRs are not resolvable with current technology, we should expect all GEHRs to be made up of a complex mix of pointlike and diffuse X-ray components with the hard-band X-ray luminosity dominated by massive stars and the soft-band X-ray luminosity dominated by the effects of recent cavity SNe expoding near the edges of superbubbles, as we see in 30~Dor.   \\subsection{Concluding Comments}  We have analyzed an early {\\em Chandra}/ACIS observation of 30~Doradus, concentrating here on the diffuse X-ray structures and in Paper II on the point sources.  This study documents a wide variety of diffuse X-ray-emitting sources:  a complex hierarchy of diffuse structures, from small-scale knots and wisps to huge superbubbles; a composite SNR including its pulsar, pulsar wind nebula, and cometary tail; an X-ray ring possibly due to a cavity SNR; bright X-ray patches associated with high-velocity H$\\alpha$ structures.  A coherent understanding of the structures is beginning to emerge from a multiwavelength comparison of the X-ray, H$\\alpha$, and mid-IR maps.  We demonstrate a variety of data analysis tools for study of ACIS fields with mixtures of pointlike sources and diffuse structures.  All software used in this study is publicly available.  From this work, we conclude the following: \\begin{itemize} \\item  30~Dor's integrated diffuse emission is well-fit by a single-temperature plasma with $T \\simeq 7$~MK, roughly solar line strengths, and a single absorbing column density of $N_H \\simeq 3 \\times 10^{21}$~cm$^{-2}$.  The soft-band luminosity is dominated by the diffuse structures, but these disappear above 2~keV and the integrated emission is dominated by just a few bright point sources (\\S\\ref{sec:globalspectra} and Paper~II). \\item  ACIS spectra of 30~Dor's diffuse emission regions often suggest chemical enrichment, another argument in favor of a SNR origin for the X-ray emission (Table~\\ref{tbl:diffspectable}). \\item  Annular spectra around the SNR N157B, centered on PSR~J0537$-$6910, indicate electron cooling through a steepening power law slope, a transition from non-thermal to thermal spectra with distance from the pulsar, and chemical enrichment in the SNR (\\S\\ref{sec:N157B}). \\item  A powerful method for understanding the complex spectral and spatial information contained in high-resolution X-ray studies of massive star-forming complexes like 30~Dor is to study maps of column density, plasma temperature, and intrinsic surface brightness derived from X-ray spectral fitting (Figure~\\ref{fig:nhktmaps}). \\item  The main 30~Dor nebula exhibits several X-ray ``voids'' not caused by obscuration, most notably the hot region 15 (\\S\\ref{sec:superspectra}).  These resemble the wind collision plasma seen in nearby H{\\sc II} regions like M~17 (\\S\\ref{sec:othermsfrs}). \\item  Comparing the diffuse X-ray emission with kinematic studies of the warm gas in H{\\sc II} region complexes is essential to disentangle the confusing morphological information.  While the fast shocks found in 30~Dor by CK94 often correlate well with regions of high X-ray surface brightness, the correlation is not complete; some fast shells lack bright X-rays while some X-ray features are not known to exhibit high velocities (\\S\\ref{sec:highvel}). \\item Some shells do not currently contain massive clusters to supply wind-generated X-rays.  The source of diffuse X-rays may be centered SNe that heated the interior of these shells without producing fast shocks against the shell walls \\citep{Chu90}. \\item  The X-ray structure that we call the West Ring (our region 17) is adjacent to CK94's NW Loop and has X-ray properties very similar to the Galactic cavity SNR HB3.  We propose that it is a cavity SNR that exploded inside 30~Dor's Shell 3 and that its projenitor most likely came from the massive cluster Hodge~301 (\\S\\ref{sec:westring}). \\item  Bright X-ray patches often have cooler plasma temperatures than their surrounding shells, possibly indicating that the increased densities associated with fast shocks allow the associated X-ray plasma to cool more efficiently than the plasma associated with the large, low-density shells. \\end{itemize}  Our upcoming 100~ks ACIS-I observation of 30~Dor will provide the high-quality X-ray dataset that is needed to place the X-ray emission in context with state-of-the-art observations in other wavebands.  We expect this observation to reveal more of the high-mass stellar population in 30~Dor and to give even more detailed information on the complex morphology of the wind-blown bubbles and superbubbles, both with higher spatial resolution imaging and with spectroscopy on finer spatial scales.  True progress in understanding, though, will require cooperation.  By combining the power of today's Great Observatories and high-quality ground-based data, we are confident that unique insight awaits."
    },
    "0601/cond-mat0601148_arXiv.txt": {
        "abstract": "We derive the energy of the surface between the normal and superfluid components of a mixed phase of a system composed of two particle species with different densities. The surface energy is obtained by the integration of the free energy density in the interface between the two phases. We show that the mixed phase remains as the favored ground state over the gapless phase in weak coupling. We find that the surface energy effects emerge only at strong coupling. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/hep-ph0601095_arXiv.txt": {
        "abstract": "We study neutrino spin oscillations in gravitational fields. The quasi-classical approach is used to describe the neutrino spin evolution. First we examine the case of a weak gravitational field. We obtain the effective Hamiltonian for the description of neutrino spin oscillations. We also receive the neutrino transition probability when a particle propagates in the gravitational field of a rotating massive object. Then we apply the general technique to the description of neutrino spin oscillations in the Schwarzschild metric. The neutrino spin evolution equation for the case of the neutrino motion in the vicinity of a black hole is obtained. The effective Hamiltonian and the transition probability are also derived. We examine the neutrino oscillations process on different circular orbits and analyze the frequencies of spin transitions. The validity of the quasi-classical approach is also considered. ",
        "introduction": "Neutrino physics becomes one the most interesting fields of elementary particle physics, especially in the wake of the recent experimental achievements in studying of solar and reactor neutrinos (see, e.g., Refs.~\\cite{Aha04,Ali05}). One of the most intriguing puzzles in neutrino physics in the neutrino oscillations problem. Neutrinos of one type can be converted into another type if there is a mixing between different neutrino eigenstates. Presently three major opportunities for neutrino oscillations are discussed in literature. The first one is neutrino flavor oscillations. The example of this oscillations type are $\\nu_e\\to\\nu_{\\mu,\\tau}$ or $\\nu_\\mu\\to\\nu_\\tau$ conversions that are the possible explanations of the solar and atmospheric neutrino problems. The second type of neutrino oscillations is the transitions between neutrino helicity states within one flavor, or neutrino spin oscillations. The last opportunity is the combination of the two first neutrino oscillations types, i.e. the transition when both flavor and helicity states can change. This type of neutrino oscillations is called neutrino spin-flavor oscillations. It can be important in solving of the solar neutrino problem. In this paper we restrict ourselves to studying of neutrino spin oscillations.  It was realized more than twenty years ago that external fields drastically change the process of vacuum neutrino oscillations. In Refs.~\\cite{Wol78,MikSmi85eng} it was established that neutrino weak interactions with background matter result in the resonant enhancement of neutrino flavor oscillations. The electromagnetic interactions are of great importance in studying of neutrino spin and spin-flavor oscillations. For example, the neutrino interaction with an external electromagnetic field provides one of the mechanisms for the mixing between different helicity eigenstates. It is also interesting to study the influence of gravitational fields on neutrino oscillations (neutrino does not seem to participate in strong interactions). Despite the gravitational interaction is relatively weak compared to electromagnetic and weak interactions there are rather strong gravitational fields in the Universe. Thus neutrino oscillations in gravitational fields are of interest in astroparticle physics and cosmology.  The influence of the gravitational interaction on neutrino oscillations has been studied in many publications (see, e.g., Refs.~\\cite{CasMon94,AhlBur96,PirRoyWud96,Ahl97,CarFul97,KosMew04,SinMobPap04,DvoGriStu05}). First we mention Refs.~\\cite{PirRoyWud96,CarFul97} where the formalism for the description of neutrino oscillations in gravitational fields in presence of matter and external electromagnetic fields was worked out. One of the approaches to the description of the spin dynamics of a neutrino in a gravitational field consists in decomposing of the Dirac Hamiltonian (in presence of an external gravitational field) and establishing terms that mix different \\emph{chirality} components of the neutrino wave function. This method was used in Ref.~\\cite{CasMon94}. To describe the \\emph{helicity} evolution in a gravitational field one can use the evolution equation in Heisenberg representation with the Hamiltonian in Foldy-Wouthuysen representation. This approach was implemented in Ref.~\\cite{SinMobPap04}. Neutrino spin oscillations and spin light of a neutrino in weak gravitational fields were studied in our work~\\cite{DvoGriStu05}. Gravity induced neutrino flavor oscillations were examined in Refs.~\\cite{AhlBur96,Ahl97}.  There is, however, another approach for the treatment of the spinning particle dynamics in an external gravitational field. The system of equations for the description of the spinning particle motion in gravitational fields was derived by A.~Papapetrou in Ref.~\\cite{Pap51}. Since then plenty of works where Papapetrou equations were analysed have been published. The major difficulties in solving Papapetrou equations are that these equations are non-linear in the particle spin as well as the motion of a spinning particle deviates from the geodesical. Thus to describe the motion of a \\emph{finite size} spinning particle in a gravitational field one has to take into account the contribution of the spin term into the particle motion law. It was established in Ref.~\\cite{KhrPom98} that Papapetrou equations can be significantly simplified when one uses the quasi-classical approach and linear in spin approximation of these equations for a \\emph{point} particle.  In this paper we study neutrino spin oscillations in gravitational fields within the quasi-classical approach. Note that the quasi-classical treatment of neutrino spin oscillations in moving and polarized matter under the influence of electromagnetic fields was considered in Refs.~\\cite{EgoLobStu00,LobStu01,DvoStu02JHEP}. In Sec.~\\ref{General} we formulate the main equations necessary for the description of the neutrino motion and spin evolution. The limit of a weak gravitational field is examined in Sec.~\\ref{WGF}. On the basis of the general equation presented in Sec.~\\ref{General} we obtain the effective Hamiltonian for neutrino spin oscillations. This effective Hamiltonian is valid for arbitrary neutrino velocity. We also compare the neutrino spin evolution equation with the analogous one derived in our previous work \\cite{DvoGriStu05} and find out that the equation obtained in the present work correctly accounts for the neutrino velocity dependence. Then we receive neutrino transition probability when a particle propagates in the gravitational field of a rotating massive object. In Sec.~\\ref{SM} we apply the general technique to the description of neutrino spin oscillations in the Schwarzschild metric. We obtain the neutrino spin evolution equation which is valid for the neutrino motion even in the vicinity of a black hole. The effective Hamiltonian and the transition probability are also derived. We examine neutrino oscillations process on different circular orbits and analyze the frequencies of the spin transitions. In Sec.~\\ref{CONCL} we discuss our results. The calculation of covariant derivatives of the vierbein vectors (Appendix~\\ref{VF}) and the basic elements of $SL(2,C)$ group (Appendix~\\ref{LorGroup}) are also presented. ",
        "conclusions": "In conclusion we note that neutrino spin oscillations in gravitational fields within the quasi-classical approach have been studied. We have started from the particle spin evolution equation in the vierbein frame. On the the basis of this approach we have investigated neutrino spin oscillations in a weak gravitational field created by a massive rotating object. The three dimensional spin evolution equation has been obtained. We have improved analogous calculation performed in our previous work \\cite{DvoGriStu05}. The neutrino spin evolution equation derived in the present work is valid for arbitrary neutrino velocities and correctly takes into account the neutrino velocity dependence. Then we have written down the neutrino oscillations effective Hamiltonian [Eq.~\\eqref{HeffJ}] and transition probability [Eqs.~\\eqref{Probabr} and \\eqref{Phir}] when a neutrino propagates along the radial direction. The neutrino spin conversion rate has been analyzed for different neutrino velocity directions. Then we have examined neutrino spin oscillations in the gravitational field of a non-rotating black hole. We have studied a neutrino moving on circular orbits, with our calculations being valid for the neutrino motion in the vicinity of a black hole. We have derived the effective Hamiltonian [Eq.~\\eqref{HeffBH}] and transition probability [Eq.~\\eqref{probBH}] in this case. It has been demonstrated that neutrino spin oscillations occur in the Schwarzschield metric. We have also analyzed the dependence of the neutrino spin oscillations frequency on the radius of the orbit. The validity of the quasi-classical approach to the description of neutrino spin oscillations in a gravitational field has been analyzed in our work.  It should be noted that the method used in the present paper for the studying of neutrino spin oscillations is valid not only for neutrinos but also for any spin-$1/2$ particles. The basic equations \\eqref{zetaeq} and \\eqref{GEB} were demonstrated in Ref.~\\cite{SilTer05} to be applicable to a massive Dirac particle in a weak static gravitational field. Therefore we can study, e.g., electron spin evolution with help of this approach. It is known that rather frequently massive gravitational objects have inherent magnetic fields and are surrounded by the moving matter (accretion disks). Thus, if we discuss the particle spin oscillations process comprehensively, i.e. take into account all factors, the cases of neutrinos and electrons are different because these particles have diverse types of the interaction with magnetic field and with the background matter. In principle massless particles can be also included into consideration (see Sec.~\\ref{SM}). Despite of the fact that the mass of a particle is absent in Eqs.~\\eqref{zetaeq} and \\eqref{GEB}, the law of motion (i.e. the function $u^a(\\mathbf{r},t)$, which is involved in the spin evolution equations) of massless particles differs from that of massive particles. However in the case of massless particles the helicity and chirality are the same and one can apply the methods used in the recent work~\\cite{Muk05}.  According to Eqs.~\\eqref{Omega1}-\\eqref{Omega3} there is no neutrino spin flip in case of radial neutrino motion in the Schwarzschield metric. This our result agrees with the deduction of Ref.~\\cite{PirRoyWud96} where the radial propagation of neutrinos from active galactic nuclei was studied. However one can hardly agree with the claims of Ref.~\\cite{CarFul97} that a spherically symmetric, static Schwarzschild space-time cannot cause the particle spin flip. That statement confront both our results and the analysis of Ref.~\\cite{SilTer05} in which it was shown by means of direct calculations that in the Schwarzschield metric the spin of a Dirac fermion can precess and change its direction with respect to the particle momentum. The calculations in that paper were based on Dirac equation in curved space. That result also agrees with the statement of our paper that the gravitational field of a non-rotating black hole can cause neutrino spin oscillations.  Note that the results of our paper can be applied to the description of the neutrino spin evolution in various astrophysical media. It is known that relic massive neutrinos can cluster into a halo around a galaxy contributing to cold dark matter. Despite of the fact that the average size of a halo ranges between $10^{2}$ and $10^{3}\\thinspace\\text{kpc}$ for a typical galaxy like the Milky Way (see, e.g., Ref.~\\cite{RinWon04}), there is a possibility for a neutrino to be gravitationally captured on orbits close to the central massive object. In this case strong gravitational field can essentially contribute to the neutrino spin oscillations process."
    },
    "0601/astro-ph0601223_arXiv.txt": {
        "abstract": "We present deep {\\it HST/ACS} observations in \\gp\\rp\\ip\\zp\\ towards the $z=4.1$ radio galaxy TN J1338--1942 and its overdensity of $>$30 spectroscopically confirmed \\lya\\ emitters (LAEs).  We select 66 \\gp-band dropouts to \\zp$_{,5\\sigma}=27$, 6 of which are also a LAE. Although our color-color selection results in a relatively broad redshift range centered on $z=4.1$, the field of TN J1338--1942 is richer than the average field at the $>$5$\\sigma$ significance, based on a comparison with GOODS.  The angular distribution is filamentary with about half of the objects clustered near the radio galaxy, and a small, excess signal ($2\\sigma$) in the projected pair counts at separations of $\\theta<10\\arcsec$ is interpreted as being due to physical pairs.  The LAEs are young (a few $\\times10^7$ yr), small ($\\langle r_{hl}\\rangle=0\\farcs13$) galaxies, and we derive a mean stellar mass of $\\sim10^{8-9}$ M$_\\odot$ based on a stacked \\ks-band image.  We determine star formation rates, sizes, morphologies, and color-magnitude relations of the \\gp-dropouts and find no evidence for a difference between galaxies near TN J1338--1942 and in the field.  We conclude that environmental trends as observed in clusters at much lower redshift are either not yet present, or are washed out by the relatively broad selection in redshift.  The large galaxy overdensity, its corresponding mass overdensity and the sub-clustering at the approximate redshift of TN J1338--1942 suggest the assemblage of a $>10^{14}$ M$_\\odot$ structure, confirming that it is possible to find and study cluster progenitors in the linear regime at $z\\gtrsim4$. ",
        "introduction": "\\label{sec:intro}  \\noindent The formation and evolution of structure in the Universe is a fundamental field of research in cosmology. Clusters of galaxies represent the most extreme deviation from initial conditions in the Universe, and are therefore good evolutionary probes for studying the formation of the large-scale structure. While clusters of galaxies have been studied extensively in the relatively nearby Universe, their evolutionary history becomes obscure beyond roughly half the Hubble time \\citep{blakeslee06,mullis05,stanford06}.  Their progenitors are extremely difficult to identify when the density contrast between the forming cluster and the field becomes small, and mass condensations on the scales of clusters are extremely rare at any epoch \\citep{kaiser84}.  Even the most distant clusters known contain, besides a population of star-forming galaxies, an older population of relatively red and massive galaxies \\citep[e.g.][]{dressler99,vandokkum00,goto05}.  The scatter in the color-magnitude relation for cluster ellipticals at $z\\sim1$ is virtually indistinguishable from that at low redshift, suggesting that some of the galaxy populations in clusters have very old stellar populations \\citep[e.g.][]{stanford98,blakeslee03_clusters,wuyts04,holden05}.  The epoch of cluster and cluster galaxy formation is presumed to be marked by the violent build-up of the stellar mass and morphologies of these early-type galaxies.  Large samples of star-forming Lyman break galaxies (LBGs) that are selected in the rest-frame UV have been used to probe the cosmic star formation rate history as well as the large-scale structure far beyond $z=1$.  The epoch corresponding to the peak in the star formation rate density ($z\\sim2-3$) is preceded by a significant but modest decrease in the star formation rate density from $z\\sim3$ out to $z\\sim6$ \\citep{madau96,steidel99,giavalisco04_results,bouwens07}.  Deep near-IR NICMOS observations in and around the GOODS fields find small samples of $z\\sim7-8$ galaxies \\citep{bouwensillingworth06,bouwens04_z7} and indicate that this decline in the number of $UV$ luminous galaxies continues from $z\\sim7-8$ \\citep{bouwensillingworth06}.  This suggests that cosmic star formation is a rather extended process, with only modest but continual changes over time.  LBGs, as well as the partially overlapping population of \\lya\\ emitters (LAEs), are strongly clustered at $z=3-5$, and are highly biased relative to predictions for the dark matter distribution \\citep{giavalisco98,adelberger98,ouchi04_r0,lee05,kashikawa05}.  The biasing becomes stronger for galaxies with higher rest-frame UV luminosity \\citep{giavalisco01}.  \\citet{ouchi04_r0} found that the bias may also increase with redshift and dust extinction. They suggest that the reddest LBGs could be connected with sub-mm sources or the extremely red objects \\citep[EROs,][]{elston88,mccarthy01,daddi02}. By comparing the number densities of LBGs to that of dark halos predicted by \\citet{sheth99}, they concluded that LBGs at $z\\sim4$ are hosted by halos of $1\\times10^{11}-5\\times10^{12}$ M$_\\odot$, and that the descendants of those halos at $z=0$ have masses that are comparable to the masses of groups and clusters. The derived halo occupation numbers of LBGs increase with luminosity from a few tenths to roughly unity, implying that there is only one-to-one correspondence between halos and LBGs at the highest masses.  Studies of sizes indicate that high-redshift star-forming galaxies are compact in size ($\\sim0.1$\\arcsec--0.3\\arcsec), while large ($\\gtrsim0.4$\\arcsec) low surface brightness galaxies are rare \\citep{bouwens04_sizes}, and there is a clear decrease in size with redshift for objects of fixed luminosity \\citep{ferguson04,bouwens04_sizes}. Morphological analysis of LBGs indicates that they often possess brighter nuclei and more disturbed profiles than local Hubble types degraded to the same image quality \\citep[e.g.][]{lotz04}.  Despite these advances in the study of the evolution of the highest redshift galaxies, {\\it galaxy clusters} have been studied out to only $z=1.5$ while it is clear that they began forming at much earlier epochs.  Finding and studying these cluster progenitors may yield powerful tests for (semi-)analytical models and $N$-body simulations of structure formation, and give new clues to how the most massive structures in the Universe came about.  Several good candidates for galaxy overdensities, possibly `protoclusters\\footnote{The term {\\it protocluster} is commonly used to describe non-virialized galaxy overdensities at high redshift ($z\\gtrsim2$) that, when evolved to the present day, have estimated masses comparable to those of the virialized galaxy clusters. See Sect. 6 for a more physical definition.}', have been discovered at very high redshift \\citep[e.g.][]{pascarelle96,keel99,francis01,moller01,steidel98,steidel05,shimasaku03,ouchi05,kashikawa07}. These structures have been found often as by-products of wide field surveys using broad or narrow band imaging.  Luminous radio galaxies have also been used to search for overdense regions within the large-scale structure at high redshift.  The association of distant, powerful radio galaxies with massive galaxy and cluster formation is mainly based on two observational clues. First, some high redshift radio galaxies are very massive systems\\footnote{We note that the measure of the absolute stellar masses of galaxies is still a matter of debate given the different implementations in the literature of the post-AGB stellar phase in the modeling of the rest frame near-infrared galaxy emission \\citep[e.g., see][]{maraston06,bruzual07,eminian07}.} \\citep[e.g.][]{debreuck02,dey97,pentericci01,villarmartin05,miley06,seymour07}, suggesting that their host galaxies are likely to be the progenitors of massive red sequence galaxies or even the brightest cluster galaxies (BCGs) that dominate the deep potential wells of clusters. Second, in some cases high redshift radio galaxies have been shown to have companion galaxies. At $1.5<z<2$, there is evidence for overdensities of red galaxies associated with radio sources, consistent with moderately rich Abell-type clusters \\citep[e.g.][]{sanchez99,sanchez02,thompson00,hall01,best03,wold03}. At $z>2$, several radio galaxies have been found to be surrounded by overdensities of LAEs of a few, discovered through deep narrow band imaging and spectroscopic follow-up with the Very Large Telescope (VLT) of the European Southern Observatory \\citep[e.g.][]{pentericci00,kurk03,venemans02,venemans04,venemans05,venemans07}.  Focussing on the excesses of LAEs discovered in the vicinity of distant radio sources, we are performing a survey of LBGs in such radio-selected protoclusters with the {\\it Advanced Camera for Surveys} on the {\\it Hubble Space Telescope} \\citep[HST/ACS;][]{ford98}.  In \\citet{miley04} and \\citet{overzier06} we reported on the detection of a significant population of LBGs in the (projected) region around radio galaxies at $z=4.1$ (TN J1338--1942) and $z=5.2$ (TN J0924--2201). Here, we will present a detailed analysis of the ACS observations of protocluster TN J1338--1942 at $z=4.1$, augmented by ground-based observations with the VLT. This structure is amongst the handful of overdense regions so far discovered at $z>4$, as evidenced by 37 LAEs that represent a surface overdensity of $\\sim5$ compared to other fields \\citep[][]{venemans02,venemans07}. The FWHM of the velocity distribution of the LAEs is 625 km s$^{-1}$. The mass overdensity as well as the velocity structure is consistent with the global properties of $z\\sim4$ protoclusters derived from simulations \\citep[e.g.][]{delucia04,suwa06}.  The radio galaxy itself is highly luminous in the rest-frame UV, optical and the sub-mm, indicating a high star formation rate. It has a complex morphology which we have interpreted as arising from AGN feedback on the forming ISM and a massive starburst-driven wind \\citep{zirm05}.  The main issues that we will attempt to address here are the following. What are the star formation histories, physical sizes and morphologies of LAEs and LBGs?  In particular we wish to study these properties in relation to the overdense environment that the TN J1338--1942 field is believed to be associated with, analogous to galaxy environmental dependencies observed at lower redshifts and predicted by models \\citep[e.g.][]{kauffmann04,postman04,delucia05}. How does sub-clustering associated with the TN J1338--1942 structure compare to the `field'? What is the mass overdensity of the structure, and what is the relation to lower redshift galaxy clusters?  In Sect. 2 we will describe the observations, data reduction and methods. We present our sample of LBGs in Sect. 3, and describe the rest-frame UV and optical properties of LBGs and LAEs. In Sect. 4 we present the results of a nonparametric morphological analysis.  In Sect. 5 we will present further evidence for a galaxy overdensity associated with TN J1338--1942 and investigate its clustering properties. We conclude with a summary and discussion of our main results (Sect. 6).  We use a cosmology in which $H_0=72$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M=0.27$, and $\\Omega_\\Lambda=0.73$ \\citep{spergel03}.  The luminosity distance is 37.1 Gpc and the angular scale size is 6.9 kpc arcsec$^{-1}$ at $z=4.1$.  The lookback time is 11.9 Gyr, corresponding to an epoch when the Universe was approximately 11\\% of its current age.  All colors and magnitudes quoted in this paper are expressed in the $AB$ system \\citep{oke71}. ",
        "conclusions": "\\label{sec:discussion}  \\subsection{The physical properties of LBGs and LAEs}  {\\it SFRs} -- We studied the star forming properties of 66 \\gp-dropouts to \\zp$=27$ mag, and 12 LAEs (6 of which are also in the \\gp-dropout sample). The SFRs were in the range 1--100 M$_\\odot$ yr$^{-1}$ (Table \\ref{tab:lbgs}), with LAEs being limited to $<14$ M$_\\odot$ yr$^{-1}$ (Table \\ref{tab:laes}). Applying an average extinction ($E(B-V)=0.1$ mag) gives SFRs of up to $\\sim300$ M$_\\odot$ yr$^{-1}$.  {\\it Ages/dust} -- The LBGs and LAEs have very blue continua ($\\beta_{iz}\\approx-2$) when averaged over the entire sample. We derived a UV-slope magnitude relation of $\\beta_{iz}=(-0.16\\pm0.05)(z_{850}-25~\\mathrm{mag})-(1.84\\pm0.11)$, and found that the slopes of the $L\\gtrsim L^*$ LBGs are consistent with the average slopes determined for mildly reddened LBGs at $z\\sim3$ redshifted to $z=4$ \\citep[Fig. \\ref{fig:lbgsbetas};][]{papovich01,shapley03}.  The $\\beta_{iz}-z_{850}$ relation can be interpreted as a SFR- or mass-extinction relation implying $E(B-V)\\approx0.13$ mag at \\zp$\\approx23$ mag and $E(B-V)\\approx0.0$ mag at \\zp$\\approx27$ mag for a fixed age of $\\sim70$ Myr.  We derived rest-frame UV-optical colors, and found LBG ages in the range 10--100 Myr, with $\\sim50$\\% of the LBGs having ages $<50$ Myr with respect to our base template \\citep[exponentially declining, $\\tau=10$ Myr, $Z=0.2Z_\\odot$ and $E(B-V)=0.16$ mag from][]{papovich01}. We also found evidence for a relation in \\ip--\\ks\\ vs. \\ks, similar as found for \\bp-dropouts in GOODS \\citep[Fig. \\ref{fig:uvopt1};][]{papovich04}. This is likely to be interpreted as a stellar mass-age or mass-dust relation, in the sense that more massive galaxies have higher optical luminosities and redder UV-optical colors due to aging or dust.  {\\it Masses} -- None of the LAEs was detected in the \\ks-band, but we found a $3\\sigma$ detection through stacking. The stacked magnitude of \\ks$=25.8^{+0.44}_{-0.32}$ mag implied \\ip--\\ks$\\approx0.0$, while a stack of LBGs with similar UV magnitudes gave a $5\\sigma$ detection of \\ks$=25.14^{+0.23}_{-0.18}$ mag and \\ip--\\ks$\\approx0.7$. Although the difference in \\ks\\ magnitude is only $\\sim2\\sigma$, other studies based on longer wavelength data suggest that LAEs may indeed be fainter at optical and infrared wavelengths than LBGs while having similar SFRs \\citep{pentericci07}. The difference can be interpreted as the LAEs being younger or less massive than LBGs. We use the \\ks-band magnitudes to derive stellar masses, finding $\\sim3\\times10^8$ M$_\\odot$ for the faintest LAEs and LBGs to $\\sim2\\times10^{10}$ M$_\\odot$ for the brightest LBGs (Fig. \\ref{fig:uvopt2}). The mean stellar mass of the LAEs in the protocluster region is in good agreement with the mean stellar masses of LAEs selected at $z=3-5$ \\citep{gawiser06,pirzkal06,pentericci07,nilsson07}.  The possible difference between the LAEs and LBGs is qualitatively consistent with \\citet{charlot93}, who find that \\lya\\ emission may primarily escape during a relatively short, dustless phase after the onset of star formation. However, observations in \\lya\\ and the UV of some local starbursts having properties similar to that of high redshift LBGs seem to suggest instead that the escape of \\lya\\ photons may have more to do with the properties of the gas kinematics rather than with dust \\citep{kunth03}.  The nature of the origin of \\lya\\ emission therefore remains unclear for the moment.  {\\it Sizes} -- For $L\\ge L^*$ galaxies we found a mean half-light radius of $0\\farcs22$ ($\\sim2$ kpc). Although this is slightly smaller than the average size reported by \\citet{ferguson04} based on a GOODS \\bp-dropout sample, \\citet{ferguson04} measured $r_{hl}$ out to much larger circular annuli. The results are consistent when we apply a small aperture correction from Fig. \\ref{fig:completeness}.  The mean $r_{hl}$ of the \\gp-dropouts is comparable to that of \\bp-dropouts culled from the UDF and GOODS fields by \\citet{bouwens04_sizes}, and is therefore consistent with the $\\propto(1+z)^{-1.05\\pm0.21}$ size scaling law that connects \\up-\\bp-,\\vp-, and \\ip-dropouts at fixed luminosities \\citep{bouwens04_sizes}. The $r_{hl}$ of the $L<L^*$ LAEs and LBGs are comparable ($\\sim1.5$ kpc). The radii of LAEs in the \\rp\\ filter that contains \\lya\\ are similar to those in the \\ip\\ filter which images our $z\\sim4$ sample purely in the UV continuum, indicating that the highest surface brightness \\lya\\ originates from the same region as the stellar continuum.  {\\it AGN} -- Although several of the LAEs appear pointlike, the spectra show no evidence for broad \\lya\\ or high ionization lines. Likewise, X-ray observations of a large field sample at $z\\sim4.5$ show no positive detections of AGN among \\lya\\ emitters \\citep{wang04}, and \\citet{ouchi07} find an AGN fraction of only $\\simeq1$\\% among \\lya\\ emitters at $z=3-4$.  On the other hand, several of the LAEs in the protocluster near radio galaxy MRC 1138--262 at $z=2.16$ have been detected with {\\it Chandra} indicating that the AGN fraction of such protoclusters could be significant \\citep[][]{pentericci02,croft05}.  {\\it Morphologies} -- The HST images show that the \\gp-dropouts span a wide range of morphologies, with some showing clear evidence for small-scale interactions (Fig. \\ref{fig:lbgzstamps}).  The morphological parameters $G$, $M_{20}$ and $C$ determined for a bright subset of the \\gp\\ dropouts suggest that the morphologies resemble those of \\bp\\ dropouts in GOODS \\citep[Fig. \\ref{fig:nick},][]{lotz04,lotz05}.  \\subsubsection{Summary}  \\noindent To summarize, the properties of our UV-selected ``protocluster'' sample of LBGs and LAEs are in good agreement with the physical properties of galaxies in field samples. In other words, we find no evidence for trends that can be linked to the relatively rich environment of the TN1338 structure, in contrast to what has been observed in lower redshift structures \\citep[e.g.][but see \\citet{peter07}]{steidel05,kodama07}.  Even if such trends in stellar mass, age, size or morphology exist at $z>4$, the relatively crude redshift selection applied to select the \\gp-dropouts at the approximate redshift of TN1338 ($\\delta(z)\\sim0.5$, FWHM) may wash out its signal due to galaxies in the immediate fore- and background.  \\subsection{Properties of the protocluster}  \\subsubsection{Clustering properties} \\label{sec:halosizes}  \\noindent We have presented evidence for an overdensity of \\gp-dropouts in TN1338 (Fig. \\ref{fig:cells}), and that this population has significant sub-clustering across the ACS field (Fig. \\ref{fig:surfdens}). The radio galaxy lies in a $\\sim7.5$ (co-moving) Mpc `filament' formed by the majority of the \\gp-dropouts. The discovery of this substructure with ACS ties in closely with the clustering seen at larger scales. \\citet{intema05} present the large-scale distribution of relatively bright $B$-dropouts towards TN1338 in a $25\\arcmin\\times25\\arcmin$\\ field observed with the Subaru Telescope, showing several significant density enhancements amidst large voids.  In contrast, \\citet{venemans02,venemans07}, using data from two $7\\arcmin\\times7\\arcmin$ VLT/FORS fields, found that the LAEs are relatively randomly distributed. Based on the absense of substructure and their small velocity dispersion they concluded that the LAEs might just be breaking away from the local Hubble flow.  Additional evidence for this might be contained in the fact that the LAEs in the much smaller ACS field seem to prefer regions that are generally devoid of the UV-selected LBGs.  We showed evidence that LAEs are generally younger (and possibly less massive) than objects in our UV-selected sample. Also, \\gp-dropouts that were detected in \\ks\\ are both brighter and redder than LAEs, suggesting that objects with both higher ages or dust and larger stellar masses lie in the densest regions of the TN1338 field \\citep[see also][]{kashikawa07}. This may be expected if the oldest and most massive galaxies predominantly end up in the population of red sequence galaxies dominating the inner part of clusters at $z\\lesssim1.5$.  We have compared the clustering of the \\gp-dropouts to that expected based on the angular clustering of \\bp-dropouts at a similar redshift. We found a small excess of clustering at the smallest angular scales ($\\theta<10\\arcsec$) compared to mock samples with a built-in large-scale correlation function (Fig. \\ref{fig:wt}).  The small-scale excess is to be expected when \\wt\\ is dominated by non-linear sub-halo clustering at small scales \\citep{kravtsov04,ouchi05c,lee05}.  A radius of $\\sim0.3-0.6$ (co-moving) Mpc is similar to the virial radius, $r_{200}$, of dark matter halos with masses of $10^{12-13}$ M$_\\odot$ \\citep[see][]{ouchi05c}, where $r_{200}$ is defined as the radius of a sphere in which the mean density is $200\\times$ the mean density of the Universe \\citep{mo02}. The linear bias, $b$, within these radii can reach values of $>10-50$, compared to a bias of $2.5-4$ for the field \\citep{ouchi05c,lee05}. Detection of this small-scale clustering implies that a significant fraction of the \\gp-dropouts may share the same halo, possibly associated with the forming structure near TN1338.  Given their overall strong clustering properties and cosmic variance, $z\\sim4$ LBGs will likely have significant field-to-field variations even on angular scales larger than currently probed by GOODS \\citep{somerville04}. Future deep, wide surveys will demonstrate the uniqueness of protocluster structures such as found in the TN1338 field. Based on shallower samples, albeit of significantly larger areas, it has been found that the surface density of TN1338-like object concentrations are comparable to that expected based on the co-moving volume densities of local clusters \\citep[e.g.][]{steidel98,shimasaku03,ouchi05,intema05}. Observations of these kind of systems could constrain the halo mass function and the halo occupation distribution at high redshift at the very high mass end.  \\subsubsection{Mass of the overdensity}  \\noindent A proper determination of the mass of the TN1338 protocluster requires a good estimate of the total volume, which depends both on the angular and redshift distribution of the LBGs and LAEs. We were not able to confirm any additional LBGs at the exact redshift of the radio galaxy, not only because of the faintness and therefore relatively uncertain photometric redshifts of the targets, but also because all objects with high equivalent width \\lya\\ (the most efficient method of confirming objects at these high redshifts) had already been found previously in this field.  We nevertheless indirectly obtained redshifts of the 6 \\gp-dropouts that are also in the \\lya\\ sample (excluding the radio galaxy).  From the redshift distribution of LBGs in the GOODS simulations we expect 2.33 field LBGs at a redshift of $z=4.1$ with $|z-\\delta z|<0.03$.  The volume overdensity in TN1338 is then $\\delta_g\\equiv(\\rho-\\bar{\\rho})/\\bar{\\rho}=1.6$, which can be considered to be a lower limit since it assumes that no other \\gp-dropouts lie within $\\delta z$ of $z=4.1$. Alternatively, given that only 20\\%--25\\% of field LBGs have large enough equivalent width \\lya\\ emission to be detected as a LAE \\citep{steidel00}, the 10 LAEs (with \\zp$<27$ mag) associated with TN1338 would represent 40--50 LBGs (including LAEs) in a relatively narrow redshift range. Note that if we subtract the average number of \\gp-dropouts expected in a TN1338-sized field ($\\sim26$) from the number of \\gp-dropouts observed, the surplus amounts to $\\sim40$ \\gp-dropouts. This is an independent confirmation of the results of \\citet{venemans02,venemans07} that were based on LAEs alone.  If we assume that the LBGs in TN1338 have the same overdensity as measured for the LAEs by \\citet{venemans02,venemans07}, the overdensity in the TN1338 ACS field would be $\\delta_g\\approx3.5$.  We can relate the true mass overdensity, $\\delta_m$, to the observed \\gp-dropout overdensity, $\\delta_g$, through $1+b\\delta_m=|C|(1+\\delta_g)$, where $C=1+f-f(1+\\delta_m)^{1/3}$ (with $f=\\Omega_m(z)^{4/7}$) corrects for redshift-space distortion due to peculiar velocities assuming that the structure is just breaking away from the Hubble expansion \\citep{steidel98}. Taking $b\\sim3$ for \\bp-dropouts with \\zp$<26-27$ from \\citet{lee05} gives $\\delta_m\\sim0.8$ for $\\delta_g=3.5$. This can be related to a total mass of $M\\simeq(1+\\delta_m)\\bar{\\rho}V\\gtrsim10^{14}$ M$_\\odot$, where $\\bar{\\rho}$ is the present-day mean density of the Universe and $V$ is the volume probed by our ACS observations.  The mass overdensity corresponds to a linear overdensity of $\\delta_L\\sim0.5$, which when evolved to the present epoch corresponds to a linear overdensity of $\\delta_L\\sim2$ \\citep[see][]{steidel05,overzier06b}.  This exceeds the linear collapse threshold of $\\delta_c=1.69$, so this ``protocluster'' will have virialized by $z=0$.  We note that these calculations depend on several critical assumptions, such as the size of the volume \\citep[likely to be larger than the ACS field,][]{venemans02,intema05}, the magnitude of the overdensity (here assumed to be constant over the volume), and the exact bias value corresponding to the objects used to determine the overdensity.  The validity of these assumptions and the typical properties of such structures will be the subject of a detailed comparison with numerical simulations in a future work (Overzier et al., in prep.).  \\subsubsection{Redshift evolution of the overdensity}  \\noindent The progenitors of galaxy clusters must have undergone rapid and intense star-formation (and possibly AGN activity) at $z\\gtrsim2$. The star formation in these `protoclusters' is not only responsible for the buildup of the present-day stellar mass in cluster galaxies, but also for the chemical enrichment of the intracluster medium \\citep[e.g.][]{ettori05,venemans05}.  The ages of the stellar populations in massive red-sequence galaxies at $z\\sim1$ are sufficiently high for them to have begun forming at redshifts $2\\lesssim z\\lesssim5$ \\citep[e.g.][and references therein]{blakeslee03,glazebrook04,mei06,homeier06,rettura06}.  Simulations suggest that there may be significant differences between the redshift of formation and redshift of assembly for the stellar mass in massive early-type galaxies: \\citet{delucia05} found that for elliptical galaxies with stellar mass larger than $10^{11}$ M$_\\odot$ the median redshift at which 50\\% of the stars were formed is $\\sim2.5$, but the median redshift when those stars were actually assembled into a single galaxy lies only at $\\sim0.8$.  In the same paper, they showed that the star formation properties of ellipticals depend strongly on the environment. For elliptical galaxies in clusters, the average ages can be up to 2 Gyr higher than those of similar mass ellipticals in the field. For an elliptical that is $\\sim1$ Gyr older compared to the field, 50\\% of its stellar mass will already have been formed at $z\\sim4$. This stellar mass is likely to be formed in much smaller units, while the number of major mergers is considered to be relatively small (a few).  It is likely that the stellar mass formed by protocluster galaxies will end up in quiescent cluster early-types at lower redshifts, in accordance with the color-magnitude and morphology-density relations. However, there is a significant discrepancy between the masses of the LBGs and LAEs (both in protoclusters and in the field) and the masses of cluster early-type galaxies of $\\gtrsim10$, indicating that a large fraction of the stellar mass still has to accumulate through merging. Detailed observations of protocluster regions on much larger scales ($\\sim$50 co-moving Mpc) are needed to test if the number density of LBGs is indeed consistent with forming the cluster red sequence population through merging.  Alternatively, protocluster fields may also host older galaxies of significant mass, analogous to the population of distant red galaxies found at $2<z<4$ \\citep[e.g.][]{franx03,vandokkum03,webb05}.  These objects are believed to be the aged and reddened descendants of LBGs that were UV luminous only at $z\\gtrsim5-6$, and form a population that is highly clustered \\citep{quadri07}. Attemps are currently being made in finding such objects toward protoclusters \\citep[e.g.][]{kajisawa06,kodama07}, but spectroscopic confirmation is difficult.  \\subsection{Comparison to simulations}  \\noindent Our results are in agreement with studies of large-scale structure and protoclusters using $N$-body simulations. \\citet{suwa06} studied global properties of protoclusters by picking up the dark matter particles belonging to clusters at $z=0$ and tracing them back to high redshift. The simulations showed that clusters with masses of $>10^{14}$ $M_\\odot$ can be traced back to regions at $z=4-5$ of 20--40 $h^{-1}$ Mpc in size, and that these regions are associated with overdensities of typical halos hosting LAEs and LBGs of $\\delta_g\\sim3$ and mass overdensities $\\delta_m$ in the range 0.2--0.6. For randomly selected regions of the same size, the galaxy and mass overdensities were found to be mostly $\\lesssim0$, as expected due to the fact that massive halos are relatively rare. Although some of the overdense regions in the simulations having a similar overdensity as our protocluster candidate do not end up in clusters at $z=0$, the simulations show that most regions with an overdensity on the order of a few at $z\\sim5$ will evolve into clusters more massive than $10^{14}$ $M_\\odot$ ($\\gtrsim50$\\% for $\\delta_g\\gtrsim2$).  Also, recent numerical simulations of CDM growth predict that quasars at $z\\sim6$ may lie in the center of very massive dark matter halos of $\\sim4\\times10^{12}$ $M_\\odot$ \\citep{springel05,li06}. They are surrounded by many fainter galaxies, that will evolve into massive clusters of $\\sim 4\\times10^{15}$ $M_\\odot$ at $z=0$. The discovery of galaxy clustering associated with luminous radio galaxies and quasars at $z>2$ \\citep[e.g.][this paper]{stiavelli05,venemans07,zheng06} is consistent with that scenario."
    },
    "0601/astro-ph0601709_arXiv.txt": {
        "abstract": "We present an interferometric search for large molecules, including methanol (CH$_3$OH), methyl cyanide (CH$_3$CN), ethyl cyanide (CH$_3$CH$_2$CN), ethanol (CH$_3$CH$_2$OH), and methyl formate (CH$_3$OCHO), in comets LINEAR (C/2002 T7) and NEAT (C/2001 Q4) with the Berkeley-Illinois-Maryland Association (BIMA) array. In addition, we also searched for transitions of the simpler molecules CS, SiO, HNC, HN$^{13}$C and $^{13}$CO .  We detected transitions of CH$_3$OH and CS around Comet LINEAR and one transition of CH$_3$OH around Comet NEAT within a synthesized beam of $\\sim$20$''$.  We calculated the total column density and production rate of each molecular species using the variable temperature and outflow velocity (VTOV) model described by Friedel et al.\\ (2005). Considering the molecular production rate ratios with respect to water, Comet T7 LINEAR is more similar to Comet Hale-Bopp while Comet Q4 NEAT is more similar to Comet Hyakutake. It is unclear, however, due to such a small sample size, whether there is a clear distinction between a Hale-Bopp and Hyakutake class of comet or whether comets have a continuous range of molecular production rate ratios. ",
        "introduction": "Comets are believed to contain the most pristine material remaining from the presolar nebula and are providing important insights into the formation mechanisms of complex molecular species. Comets are primarily located in two distinct regions of the solar system. The Oort cloud, the source of long period (P$>$200 years) comets, is a spherically symmetric distribution of comets that encompasses the solar system out to a distance of nearly 100,000 AU. The Kuiper-Edgeworth belt that lies in the ecliptic plane just beyond the orbit of Pluto out to several hundred AU is the source of short period (20$<$P$<$200 years) comets (Jewitt 2004). Recent models have suggested that the Oort cloud comets may have had origins in the entire giant planet region between Jupiter and Neptune.  These models suggest that comets were formed in a much wider range of nebular environments and probably experienced thermal and collisional processing before they were ejected into the Oort cloud (Weissman 1999).  This processing may have ``homogenized'' the cometary nuclei of Oort cloud comets.  Observations comparing long and short period comets, which are believed to have had different origins, will allow us to address the molecular diversity in comets (A'Hearn et al.\\ 1995, Mumma et al.\\ 2005).  Furthermore, a key goal in astrochemistry is to learn whether the molecular diversity seen in high and low mass hot molecular cores (HMCs) is similar to the chemistry of the primordial solar nebula.  The molecular diversity in comets may provide the link between the formation mechanisms of complex molecular species seen in high mass HMCs and low mass young stellar objects to the formation of young solar systems.   Therefore, only through comprehensive surveys of HMCs and comets will a consistent, coherent picture emerge regarding the basic chemistry involved in the evolution of HMCs and comets.  Most observations of the molecular composition of cometary gas and dust have taken place in the visible, ultraviolet (see e.g.\\ Hutsem\\'{e}kers et al.\\ 2005 and references therein) or infrared parts of the spectrum (see e.g.\\ Helbert et al.\\ 2005; Dello Russo et al.\\ 2001).  Also, a vast inventory of interstellar ices and volatiles was discovered by satellites passing through the coma of Comet Halley (see e.g.\\ Mitchell et al.\\ 1992).  However, there have been several successful detections of the rotational transitions of molecular species, including organic compounds, at millimeter wavelengths (see e.g.\\ Crovisier et al.\\ 2004a,b; Biver et al.\\ 2002). Many of these species are important in prebiotic organic chemistry (e.g.\\ H$_2$O, HCN, CH$_3$OH, aldehydes, nitriles) and are believed to be parent molecules (originating from the comet nucleus) rather than the products of photodissociation or gas phase chemistry in the cometary coma.  Two prominent examples of parent molecules of prebiotic importance are hydrogen cyanide (HCN) (see e.g.\\ Friedel et al.\\ 2005) and methanol (CH$_{3}$OH) (see e.g.\\ Ikeda et al.\\ 2002).  In the Spring of 2004, there was a rare opportunity to observe two dynamically new Oort cloud comets passing into the inner solar system and within $\\sim$0.3 AU of the Earth: Comet LINEAR (C/2002 T7) (hereafter, T7 LINEAR) and Comet NEAT (C/2001 Q4) (hereafter, Q4 NEAT). This paper presents the results of an effort to further investigate the astrochemistry of comets by observing several large molecular species in Comets T7 LINEAR and Q4 NEAT with the Berkeley-Illinois-Maryland Association (BIMA) array\\footnote{Operated by the University of California, Berkeley, the University of Illinois, and the University of Maryland with support from the National Science Foundation.}; these species include methanol (CH$_3$OH), methyl cyanide (CH$_3$CN), ethyl cyanide (CH$_3$CH$_2$CN), ethanol (CH$_3$CH$_2$OH), and methyl formate (CH$_3$OCHO). In addition, we also searched for transitions of the simpler molecules CS, SiO, HNC, HN$^{13}$C and $^{13}$CO.  We were successful in obtaining single-field images, cross-correlation spectra, and production rates for cometary methanol (CH$_3$OH) in both comets and CS in Comet T7 LINEAR.  Upper limits were found for the remaining species in both comets. ",
        "conclusions": "\\subsection{Column Densities and Production Rates}\\label{sec:NT}  As described in Friedel et al.\\ (2005), the VTOV model calculates the total column density and production rate of a cometary species assuming optically thin emission in LTE and that the temperature and outflow velocity within the coma vary with cometocentric distance.  The VTOV model calculations to determine the total column density and production rate are discussed in detail in Friedel et al.\\ (2005) Appendix A. Here, we discuss the implication of the production rates and column densities. Using the VTOV model, tables 3 and 4 list the total beam-averaged column densities and production rates of each molecular species observed toward Comets T7 LINEAR and Q4 NEAT. The VTOV model was also used to determine the total beam-averaged column density and production rate of CS even though it is believed to be the daughter species. The parent molecule of CS may be CS$_2$, which has a very short lifetime of $\\sim$10$^3$ s or less (Snyder et al.\\ 2001).  This short lifetime suggests that CS is formed in the inner coma and the measured line width of CS in Comet LINEAR is similar to HCN and CH$_3$OH, which are known parent species. Thus, at the spatial resolution of our observations, CS can be fitted as though it were a parent species.   \\subsection{Relative Production Rates of X to H$_2$O and CN}  \\subsubsection{Comet LINEAR} Based upon the H$_2$O production rates as measured by Schleicher et al.\\ (2005, private communication) at different heliocentric distances, Friedel et al.\\ (2005)  determined the H$_2$O production rate to be $\\sim2\\times10^{29}$ s$^{-1}$ during our observing period.  The relative production rates of each molecular species relative to H$_2$O are listed in column 6 of Table 3. The relative production rate ratios of HCN (Friedel et al.\\ 2005), CH$_3$OH, and CS with respect to H$_2$O for Comet LINEAR are most similar to those observed around Comet Hale-Bopp (C/1995 O1) (($2.1-2.6)\\times10^{-3}$, $4\\times10^{-2}$, and $(3.7-4.3)\\times10^{-3}$, respectively) (Snyder et al.\\ 2001, Remijan et al.\\ 2004). Schleicher et al.\\ (2005) estimated the CN production rate to be $\\sim1.4\\times10^{26}$ s$^{-1}$.  We also attempted to determine the relative production rates of each nitrogen bearing molecular species relative to CN. However in this work, no species observed toward Comet LINEAR with a CN bond was detected above the 3 $\\sigma$ detection limit.  The upper limits to the production rate ratios are listed in column 7 of Table 3.  \\subsubsection{Comet NEAT}  Schleicher et al.\\ (2005) also determined the H$_2$O and CN production rates for comet NEAT during our observing period (as reported by Friedel et al.\\ 2005).  These are very similar to water production rates measured from Comet Hyakutake at similar heliocentric distances ($\\sim1\\times10^{29}$ s$^{-1}$ Lis et al.\\ 1997).  The relative production rates of each molecular species relative to H$_2$O are listed in column 6 of Table 4.  The relative production rate ratios of HCN (Friedel et al.\\ 2005) and CH$_3$OH relative to H$_2$O are most similar to those measured around Comet Hyakutake (C/1996 B2) ($1.2\\times10^{-3}$ and $1.7\\times10^{-2}$, respectively) (Biver et al.\\ 1999). The average CN production rate is Q(CN)$\\sim2.6\\times10^{26}$ s$^{-1}$.  The relative production rates of each nitrogen bearing molecular species relative to CN are listed in column 7 of Table 4, however the only species observed toward Comet NEAT with a CN bond was CH$_3$CN which was not detected above the 3 $\\sigma$ detection limit.  The upper limit to the production rate ratio was $<$4.0$\\times$10$^{-2}$.   \\subsection{Classifying Comets T7 LINEAR and Q4 NEAT}  Comets are diverse objects, as has been noted by many researchers in the past. For example, A'Hearn et al.\\ (1995) reported an optical study of 85 comets, Biver et al.\\ (2002) reached this conclusion based upon a radio survey of 24 comets, while Mumma et al.\\ (2005) reached the same conclusion from the hypervolatile gases CO and CH$_4$ from observations at IR wavelengths. Many comet observers have searched for differences between comets from the Oort cloud and the short-period (Jupiter family) comets to attempt to find clear differences in chemical composition depending on their place of formation. As more observations become available, it becomes harder to differentiate between the various comets classes, in particular after the Deep Impact mission (e.g., Mumma et al, 2005).  We can now include Comets T7 LINEAR and Q4 NEAT in the chemical comparison from the detections of both CH$_3$OH and HCN (Friedel et al.\\ 2005).  Both comets T7 LINEAR and Q4 NEAT are believed to be Oort cloud comets given their measured orbital parameters and our observations show that both comets are chemically very similar.  However, from the Biver et al.\\ (2002) 30 m survey, there appears to be a distinction between comets rich in HCN (HCN/H$_2$O$>$0.2\\%) including comets Hale-Bopp (C/1995 01), 109P/Swift-Tuttle, 1P/Halley and 9P/Tempel 1 (Mumma et al.\\ 2005) compared to comets with an HCN/H$_2$O abundance ratio $\\sim$0.1\\% which include, among several others, comets Hyakutake (C/1996 B2), Austin (C/1989 X1) and Levy (C/1990 K1). It is unclear, however, due to such a small sample size, whether in fact there are two distinct classes of comets (i.e.\\ Hale-Bopp and Hyakutake class) or if comets have a range of different molecular production rate ratios. Given the abundance ratios of Comet T7 LINEAR, it is more similar Comet Hale-Bopp that was both HCN and CH$_3$OH rich.  Comet Q4 NEAT on the other hand, has column densities of HCN and CH$_3$OH more similar to Comet Hyakutake. Furthermore, this distinction appears across both Oort cloud and Kuiper-Edgeworth belt comets as illustrated by recent Deep Impact observations of Comet 9P/Tempel 1 (Mumma et al.\\ 2005). Thus, while our observations still tend to support a chemical distinction between Hale-Bopp and Hyakutake class comets due to the abundance rations of CH$_3$OH/H$_2$O and HCN/H$_2$O, more observations of both long and short period comets are necessary to increase the sample size.  With better statistics, it may be possible to reliably determine the origins of both Oort cloud and Kuiper-Edgeworth belt comets in the pre-solar nebula."
    },
    "0601/astro-ph0601468.txt": {
        "abstract": "We report the results of a spectroscopic search for debris disks surrounding 41 nearby solar type stars, including 8 planet-bearing stars, using the {\\it Spitzer Space Telescope}. With accurate relative photometry using the  Infrared Spectrometer (IRS) between 7-34 $\\micron$ we are able to look for excesses as small as $\\sim$2\\% of photospheric levels with particular sensitivity to weak spectral features. For stars with no excess, the $3\\sigma$ upper limit in a band at 30-34 $\\mu$m corresponds to $\\sim$ 75 times the brightness of our zodiacal dust cloud. Comparable limits at 8.5-13 $\\mu$m correspond  to $\\sim$ 1,400 times the brightness of our zodiacal dust cloud. These limits correspond to material located within  the $<$1 to $\\sim$5 AU region that, in our solar system, originates from debris associated with the asteroid belt.  We find excess emission longward of $\\sim$25 $\\mu$m from five stars of which  four also show excess emission at 70 $\\mu$m. This emitting dust must be located around 5-10 AU. One star has 70 micron emission but no IRS excess.  In this case, the emitting region must begin outside 10 AU; this star has a known radial velocity planet. Only two stars of the five show emission shortward of 25 $\\micron$ where spectral features reveal the presence of a population of small, hot dust grains emitting in the 7-20 $\\mu$m band. The data presented here strengthen the results of previous studies to  show that excesses at 25 $\\micron$ and shorter are rare: only 1 star out of 40 stars older than 1 Gyr or $\\sim 2.5$\\% shows an excess.  Asteroid belts 10-30 times more massive than our own appear are rare among mature, solar-type stars. ",
        "introduction": "The infrared excess associated with main sequence stars \\citep{aumann1984} is produced by dust orbiting at distances ranging from less than 1 to more than 100 AU. In discussing these debris disks we often use the analogy of our Solar System, where dust is generated by collisions between larger bodies in the asteroid and  Kuiper belts, as well as outgassing comets. As in the Solar System, we distinguish between: a) hot (200-500 K) debris located relatively close to the star ($<$1 AU out to $\\sim$ 5 AU) and arising in a region dominated by refractory material, e.g. debris associated with an asteroid belt \\citep{dermott02};  and b) cooler debris ($<30-200$ K) located beyond 5 AU that may be associated with more volatile, cometary material, i.e. the analog of our Kuiper Belt. Observed dust luminosities range from $\\ld \\simeq 10^{-5}$ to greater than $10^{-3}$ for main-sequence stars, compared to $\\ld \\simeq 10^{-7}-10^{-6}$ inferred for our own Kuiper belt  \\citep{stern1996,  backman1995} and $\\sim 10^{-7}$ measured for our asteroid belt \\citep{bp1993, dermott02}.  The FGK survey of solar-type stars is a \\spit\\ GTO program designed to search for excesses around a broad sample of 148 nearby, F5-K5 main-sequence field stars,   sampling wavelengths from 7-34 \\um\\ with the Infrared Spectrograph (IRS; Houck et al. 2004) and 24 and \\70um with  the Multiband Imaging Photometer for Spitzer (MIPS; Rieke et al. 2004). In the context of our  paradigm,  IRS is well suited to searching for and characterizing hot dust in an asteroid belt analogue while MIPS is better suited to studying cooler material in a Kuiper Belt analogue. The sample consists of two components --- stars with planets and those without. The MIPS observations of planet-bearing stars are discussed in Beichman et al. (2005a) and of non planet-bearing stars in Bryden et al. (2006).  This paper presents the results of the IRS survey of a sub-sample of 39 stars, 7 with planets, and the remainder without planets based on current knowledge from radial velocity programs. In addition, the paper includes IRS follow-up observations of 2 stars (one of which is planet-bearing) that were not originally scheduled for spectroscopy,  but which were found to have a 70 $\\mu$m excess. Results for one star with a notable IRS spectrum, HD 69830, have been presented elsewhere \\citep{beichman2005b} but will be summarized below in the context of the complete survey.  The IRS portion of the program was designed with four goals in mind: 1) to extend the wavelength coverage beyond that possible with MIPS to look for dust  emitting either shortward or longward of MIPS's $24\\, \\mu$m band; 2) to take advantage of IRS's broad wavelength coverage to reduce systematic errors and thus, possibly, identify smaller excesses than previously  possible; 3) to constrain the spatial distribution of dust grains by assessing the range of temperatures present in a given debris disk; and 4) to search for mineralogical features in the spectra of any excesses. In this paper we describe the sample of stars ($\\S2$);  present our calibration procedure and give the reduced spectra for each star ($\\S3$ and Appendix-1); discuss detections of or limits to IR excesses ($\\S4$); present models of the excesses in the 4 stars with IRS and MIPS 70 um excesses and discuss the nature of their debris disks; discuss overall results and conclusions ($\\S5$ and $\\S6$). ",
        "conclusions": "We have used the IRS spectrometer on the Spitzer Space Telescope  to look for weak excesses around main sequence, solar type stars. After careful calibration, the IRS spectra allow us to achieve relatively low values of $\\ld$ compared to previous studies. At 8-13 $\\micron$ the IRS data provide an average 3$\\sigma$ limit of $\\ld =14 \\pm 6 \\times 10^{-5}$ or roughly  1,400 times the level of our solar system zodiacal cloud for material in the 1 AU region of the target stars.  Two stars, HD 72905 and HD 69830, have detectable emission at these wavelengths and in both cases there is evidence for emission by small dust grains. At 30-34 $\\micron$, we reach a more  sensitive, limit on $\\ld$ of $0.74 \\pm 0.31 \\times 10^{-5}$, or roughly  74 times the level of our solar system zodiacal cloud for material in the $5-10$ AU  region which may represent the inner edge of the Kuiper Belt. Four stars (HD 7570, HD 72905, HD 76151, HD 206860) showed evidence for long wavelength IRS emission that appears to link smoothly to a 70 $\\micron$ MIPS excess. HD 69830 shows no evidence for 70 $\\micron$ excess and thus for colder, more distant material.  The lack of IRS excess emission toward planet-bearing stars may be due to the small size of the sample observed to date, to the effects of planets removing dust from the orbital locations probed at IRS wavelengths, or to more general processes that remove dust from regions interior to a few AU \\citep{dominik2003}. We suggest that asteroid belts located between 1-10 AU and as massive as 8-35 times our own asteroid belt are rare around mature, solar type stars."
    },
    "0601/astro-ph0601479_arXiv.txt": {
        "abstract": "\\begin{center} {\\large\\bf Abstract} \\end{center} We present a perturbative treatment of the evolution under their mutual self-gravity of particles displaced off an infinite perfect lattice, both for a static space and for a homogeneously expanding space as in cosmological N-body simulations. The treatment, analogous to that of perturbations to a crystal in solid state physics, can be seen as a discrete (i.e. particle) generalization of the perturbative solution in the Lagrangian formalism of a self-gravitating fluid. Working to linear order, we show explicitly that this fluid evolution is recovered in the limit that the initial perturbations are restricted to modes of wavelength much larger than the lattice spacing. The full spectrum of eigenvalues of the simple cubic lattice contains both oscillatory modes and unstable modes which grow slightly faster than in the fluid limit. A detailed comparison of our perturbative treatment, at linear order, with full numerical simulations is presented, for two very  different classes of initial  perturbation spectra. We find that the range of validity is similar to  that of the perturbative fluid approximation (i.e. up to close to ``shell-crossing''), but that the accuracy in tracing the evolution is superior. The formalism provides a powerful tool to systematically calculate discreteness effects at early times in cosmological N-body simulations. ",
        "introduction": "The standard paradigm for formation of large scale structure in the universe is based on the growth of small initial density fluctuations in a homogeneous and isotropic medium (see e.g. \\cite{peebles_80}). In the currently most popular cosmological models, a dominant fraction (more than 80 \\%) of the clustering matter in the universe is assumed to be in the form of microscopic particles which interact essentially only by their self-gravity.  At the macroscopic scales of interest in cosmology the evolution of the distribution of this matter is then very well described by the Vlasov or ``collisionless Boltzmann'' equations coupled with the Poisson equation (see e.g. \\cite{binney_87}).  A full solution, either analytical or numerical, of these equations starting from appropriate initial conditions (IC) is not feasible. There are, on the one hand, various perturbative approaches to their solution (for reviews see e.g. \\cite{sahni_95,bernardeau_02}), which allow one to understand the evolution in some limited range (essentially of small to moderate amplitude fluctuations). On the other hand, there are cosmological N-body simulations (for reviews see \\cite{bagla_97,bertschinger_98,hockney_99}), which solve numerically for the evolution of a system of $N$ particles interacting purely through gravity, with a softening at very small scales. The number of particles $N$ in the very largest current simulations \\cite{springel_05} is $\\sim 10^{10}$, many more than two decades ago, but still many orders of magnitude fewer than the number of real dark matter particles ($\\sim 10^{80}$ in a comparable volume for a typical candidate).  The question inevitably arises of the accuracy with which these ``macro-particles''  trace the desired correlation properties of the theoretical models. This is the problem of discreteness in cosmological N-body simulations. It is an issue which is of considerable importance as cosmology requires ever more precise predictions for its models for comparison with observations.  Up to now the primary approach to the study of discreteness in N-body simulations  has been through numerical studies of convergence (see e.g. \\cite{power_03,diemand_04}), i.e., one changes the number of particles in a simulation and studies the stability of the measured quantities. Where results seem fairly stable, they are assumed to have converged to the continuum limit. While this is a coherent approach, it is far from conclusive as, beyond the range of perturbation theory, we have no theoretical ``benchmarks'' to compare with. Nor is there any systematic theory of discreteness effects, e.g., we have no theoretical knowledge of the $N$ dependence of the convergence. Given that typically $N$ is varied over a very modest range (typically one or two orders of magnitude) compared to that separating the simulation from the model (typically 70 orders of magnitude) there is much room for error.  Different mechanisms by which discreteness effects may make the evolution of N-body simulations different to that of the fluid limit have been discussed in the literature. A very basic consideration is that of the discreteness effects introduced already in the IC, before any dynamical evolution. Indeed there are necessarily discrepancies between the correlation properties of the discretised IC and those of the input theoretical model, as there are intrinsic fluctuations associated with the particles themselves.  For analysis and discussion of these effects see, e.g., \\cite{baertschiger_02,knebe_02,knebe_03,baertschiger_03,joyce_04,joyce_04b}. What is probably the most obvious effect of discreteness, and certainly the one most emphasised in the literature, is two body collisionality: pairs of particles can have strong interactions with one another, which is an effect absent in the collisionless limit. For analysis and discussion of these effects see, e.g., \\cite{melott_93,melott_97,splinter_98,baertschiger_02t,binney_02}.    In this paper we present an approach which allows in principle a systematic understanding of discreteness effects {\\it between these two regimes}, in the evolution from the IC up to the time when two body collisions start to occur.  We do so by developing a perturbative solution to the fully discrete cosmological N-body problem, which is valid in this regime.  This essentially analytic solution can be compared to the analogous fluid ($N \\rightarrow \\infty$) solution, and one can understand exhaustively the modifications introduced, at a given time and length scale, by the finiteness of $N$.  While the usefulness of the approach is restricted to the regime of validity of this perturbative approach, we can gain considerable insights into the effects of discreteness and how they introduce error.  Some of the essential results have already been briefly reported in \\cite{jmgbsl_05}. In this paper we describe in much greater detail the perturbative method used to describe the evolution, and evaluate its regime of validity by extensive comparison with numerical simulations. In a forthcoming paper \\cite{joyce_06} we will discuss the application of this method to the study of discreteness effects in N-body simulations, providing precise quantifications of these effects in the regime in which our treatment is valid.  The perturbative scheme we employ is one which is well known and standard in solid state physics, as it is that used in the treatment of perturbations of a crystal about the local or global minimum of its internal energy. Indeed the class of cosmological N-body simulations we consider are those which start by making very small perturbations to particles initially placed on a perfect simple lattice \\cite{efstathiou_85,navarro_96,bertschinger_98,power_03,smith_03,diemand_04}. Up to an overall change in sign, our perturbative scheme is precisely that one would use for the analysis of the extensively studied Coulomb lattice or Wigner crystal (see e.g. \\cite{clark_57,pines_63}), of $N$ particles on a lattice interacting by a pure unscreened Coulomb force. At linear order (harmonic analysis) one simply solves a $3N \\times 3N$ eigenvalue problem to determine the eigenmodes and eigenvalues of the displacements off the crystal.  This can be done at very low cost in computational resources because of the symmetries of the lattice. Stable (i.e.  dynamically oscillating) modes in one problem become unstable (growing and decaying) modes in the other problem, and vice versa.  One consequence of this which we will discuss briefly here, and more extensively in \\cite{joyce_06}, is that, for what concerns discreteness in N-body simulations, there are qualitatively different features on the simple cubic (sc) lattice compared to the body centered cubic (bcc) or face centered cubic (fcc) lattice.  A crucial step in our analysis is evidently the comparison of our solutions for the evolution with those obtained from the treatment of the continuous self-gravitating system.  This latter problem can also be solved in a limited range with a perturbative treatment of the equations of a self-gravitating fluid, which are obtained by truncation of the collisionless Boltzmann equation. Given that our perturbation scheme works with the displacements of the particles, one might anticipate that the appropriate perturbation scheme to compare with is that given in the Lagrangian formulation of these fluid equations, in which the evolution is described in terms of the trajectories of fluid elements \\cite{buchert_92}.  We show explicitly, at linear order, that this is the case: taking the limit in which the perturbations to the lattice are of wavelengths much larger than the lattice spacing $\\ell$ the evolution described by our scheme maps precisely on to that at the same order in the Lagrangian description of the fluid. The Zeldovich approximation, which is simply the asymptotic form of this solution, can then be understood in very simple analogy with the long-wavelength coherent ``plasma oscillations'' in a unscreened charged plasma.    The paper is organized as follows. In the next section we describe the perturbative treatment for perturbations off a perfect lattice, for the specific case of gravity. We work firstly, for simplicity, in a static Euclidean universe, giving the explicit expressions for the evolution from general IC (i.e. any perturbation from the lattice). In the following section we explicitly solve the evolution for the case of a simple cubic lattice, and discuss in detail the structure of the spectrum of eigenvalues. In Sect.  IV we generalize our results to the case of an expanding universe, and then show explicitly the recovery of the fluid limit given by the solution at the same order of the Lagrangian formulation of the equations of a self-gravitating fluid. In the next section we present a comparison of the evolution described by our approximation with that of full numerical N-body simulations.  We consider both uncorrelated initial perturbations (a ``shuffled lattice'') and a set of highly correlated perturbations (with a power spectrum of density fluctuations $\\sim k^{-2}$), like that in current cosmological models. We verify that the agreement is very good, and that the evolution is traced with considerably better accuracy than by the fluid limit at the same order. Both approximations break down when particles start to approach one another (i.e. ``shell-crossing'' in fluid language), which we parametrise through an appropriate statistical measure. In the final section we summarise our results and discuss various further developments of the method presented here which could be pursued, notably the extension of the perturbation scheme to higher than linear order. This may throw further light on how insights about discreteness gained using this formalism, which will be discussed at length in \\cite{joyce_06}, extend into the regime of highly non-linear evolution. We also discuss briefly the possible interest of solving the cosmological N-body problem on the bcc lattice, as well as the possible extension of our method to IC generated on ``glassy'' configurations. ",
        "conclusions": "In this paper we have described in detail a new perturbative treatment which describes the evolution of $N$ self-gravitating particles of equal mass initially perturbed off a perfect lattice, subject to periodic boundary conditions, both in a static spacetime and in a cosmological (expanding) background.  We have reported specifically the spectrum of eigenvectors and eigenvalues for the modes of the displacement field on a {\\it simple cubic} lattice, which is the case of relevance in cosmology. While the fluid limit ($N \\rightarrow \\infty$) is recovered for long wavelength modes, the full spectrum for the finite $N$ system contains both modes with negative eigenvalues, corresponding to oscillations, and modes with exponents greater than in the fluid limit. Further the eigenvalues depend explicitly not only on the modulus of the wave-vector $\\bk$, but also on its orientation with respect to the axes of the lattice. The breaking of rotational invariance in the lattice is thus  imprinted in the evolution of the system. We have shown, by detailed comparison with numerical simulations, that the linear order of the scheme has approximately the same range of validity as the corresponding (linear) order of the fluid theory, up to when particles come very close to one another (i.e. up to ``shell-crossing'' in fluid language).  However it traces the real evolution with greater accuracy than its fluid counterpart.  In the context of cosmological simulations this means that our method provides a precise tool for quantifying {\\it fully}, up to shell crossing, the effects of discreteness in these simulations: these effects are nothing other than the difference between the finite $N$ evolution and the fluid limit. In a forthcoming paper \\cite{joyce_06} we will explore more extensively this application, giving precise quantitative measures of these effects adapted for use in ``correcting'' such N-body simulations.  We conclude this paper by commenting on a few possible developments of the perturbative theory we have described here.  Firstly the method can evidently also be extended to higher order, just as has been done in the analogous treatment of condensed matter system (see e.g. \\cite{carr_61,carr_61b}).  It is straighforward to generalize Eq.~\\eqref{def-dyn} to an expansion to any order. The $\\mu$ component of the force reads: \\bea \\label{expansion_force} F_{\\mu}(\\br)&=&\\sum_{n=0}^{\\infty}\\sum_{\\bR'\\neq \\bR}\\frac{1}{n!}\\mathcal{G}_{\\mu,\\nu_1\\dots\\nu_n}^{(n)}(\\bR-\\bR')\\\\\\nonumber &\\times& [u_{\\nu_1}(\\bR)-u_{\\nu_1}(\\bR')]\\dots [u_{\\nu_n}(\\bR)-u_{\\nu_n}(\\bR')] . \\eea The tensor $\\mathcal{G}_{\\mu,\\nu_1\\dots\\nu_n}^{(n)}$ is a function only of the interaction potential $v(\\br)$ and is given by: \\be \\label{dynamical_general} \\mathcal{G}_{\\mu,\\nu_1\\dots\\nu_n}^{(n)}(\\bR)= - \\left[ \\frac{\\partial^{(n+1)}w(\\ve r)}{\\partial r_{\\mu}\\partial r_{\\nu_1}\\dots\\partial r_{\\nu_n}} \\right]_{\\ve r= \\ve R}. \\ee The analysis can be followed through at any order in analogy to linear order. Transforming to reciprocal space the problem simplifies, but there is now the added complexity of a coupling of modes. A first calculation of interest would be to map this description, in the fluid limit, onto the corresponding one at the same order in the fluid Lagrangian theory. This latter treatment has been explored extensively in the cosmological literature and compared in detail with numerical simulations (see e.g. \\cite{buchert_93} and references therein). The study of the corrections to this limit should allow one to get some insight on the interplay of non-linearity (which is of course also a feature of the fluid model) and the effects of discreteness. Non-linearity in gravity involves the transfer of power from larger to smaller scales, an effect which is often qualitatively argued to make discreteness effects of less consequence. One might hope to see such a mechanism at play, if indeed it is there.  The approach presented here may also prove useful in providing insight about the nature of existing approximations to self-gravitating systems which go beyond the simple fluid limit, as it provides an ``exact'' evolution of a self-gravitating finite N-body system in a certain range. For example approximations have been developed to self-gravitating systems involving pressure terms associated to velocity dispersion (see e.g. \\cite{buchert_93,buchert_98,tatekawa_02,tatekawa_04b,dominguez_04,buchert_05}). In principle these terms can be calculated exactly using our perturbative scheme and the improvement (if this is the case) they allow to the approximation of the full evolution better understood.  Another direction in which this treatment can be generalised is to the consideration of other initial configurations. We have analysed here almost exclusively the case of perturbations from a sc lattice, as this is the kind of lattice used in cosmological simulations. Our treatment can easily be generalised to other lattices, and a comparative study of the discreteness effects should be straightforward. Without doing any calculation however one simple and interesting result can be given. It is known that both the bcc and fcc lattice are stable (or at least meta-stable) configurations for the Coulomb lattice.  This means that there are no unstable modes. For gravity this implies that there are \\emph{only} unstable modes. There are therefore no oscillating modes, and thus, by the Kohn sum rule, no modes with exponent greater than in the fluid limit. Consequently either of these lattices would appear to be better lattices to use in N-body simulations, in which one wishes evidently to approach as closely as possible the fluid evolution. The bcc lattice would appear to be the more interesting of the two, as it is known to be \\cite{torquato_03} the most densely packed lattice.  It is also a more isotropic configuration than either the sc or fcc lattice.  Another case of interest to analyse would be that of  ``glassy'' configurations, which are often used as an alternative to the sc lattice in cosmological simulations \\cite{white_94,jenkins_98,springel_05}. These configurations are generated by simulating a set of point particles evolving under negative gravity (i.e. Coulomb forces), with an appropriate damping term. By doing so one arrives at a configuration in which the forces are very small, but which is more isotropic than the lattice (in which the forces are exactly zero). In the approximation that the initial forces are negligible one could, in principle, carry out the same kind of analysis as given here. The only difference is that the $3N$ eigenmodes of the displacement field will not be plane waves, which greatly complicates the analysis compared to the case of the lattice. Numerically however such a solution should be feasible (for any specified glassy configuration), and it would be necessary if this method is to be used to give a precise quantification of discreteness effects from these IC as we can now give, using the analysis presented here, for the case of a lattice. We do not expect the results, however, to be very different (either qualitatively or quantitively): the effects described here are essentially sampling effects which depend on the sampling scale (the lattice spacing $\\ell$ above) and not on the precise nature of the sampling.  The particular manifestation of anisotropy which we have observed in the sc lattice will necessarily be quite different, and likely less pronounced, in the glassy case, but the average slowing down of growth at smaller scales would be expected to be very similar in magnitude."
    },
    "0601/astro-ph0601129_arXiv.txt": {
        "abstract": "We calculate NLTE line-driven wind models of selected O stars \\zm{in the spectral range of O4 to O9} in the Small Magellanic Cloud (SMC). We compare predicted basic wind properties, i.e. the terminal velocity and the mass-loss rate with values derived from observation. We found relatively good agreement between theoretical and observed terminal velocities. {On the other hand,} predicted mass-loss rates and mass-loss rates derived from observation {are in} a good agreement only for higher mass-loss rates. {Theoretical} mass-loss rates lower than approximately $10^{-7}\\,\\text{M}_\\odot\\,\\text{year}^{-1}$ are significantly higher than those derived from observation. {These results confirm the previously reported problem of weak winds, since our calculated mass-loss rates are in a fair agreement with predictions of Vink et al.~(2001).} We study multicomponent models for these winds. For this purpose we develop a more detailed {description of} wind decoupling. {We show that the instability connected with} the decoupling of individual wind elements {may occur for low-density winds}. In the case of winds with {very} low observed mass-loss rate the multicomponent effects are important for the wind structure, however this is not able to consistently explain the difference between predicted mass-loss rate and mass-loss rate derived from observation {for these stars}. {Similar to previous studies, we found the dependence of wind parameters on the metallicity.} We conclude that the wind mass-loss rate significantly {increases with} metallicity {as $\\dmdt\\sim Z^{0.67}$}, whereas the wind terminal velocity on average depends on metallicity only slightly, {namely $v_\\infty\\sim Z^{0.06}$} (for studied {stars}). ",
        "introduction": "In recent few years 8m class telescopes became routinely available for {the} stellar research. This enabled {detailed} study of many stars in the Local Group. Clearly, many astrophysically important types of stars can be studied from another perspective. This is especially true for stars from the Small Magellanic Cloud (SMC). Their lower metallicity compared to the solar value ({the average metallicity of individual elements is} $Z/Z_\\odot=0.2$, e.g. Venn \\citeyear{ven}, Bouret at al. \\citeyear{bourak}) enables {to} study {the} stellar properties and evolution with respect to metallicity.  One of the most important properties of hot stars, that can influence both {the} observed spectrum and {the} stellar evolution, is the presence of the stellar wind. The existence of significant dependence of the basic wind properties (i.e. {the} {amount of mass expelled from the star per unit of time, the} mass-loss rate and the wind velocity \\zm{at large distances from the star}, {the} terminal velocity) on the metallicity has been anticipated already at the very beginning of the theoretical study of hot star winds {(Abbott \\citeyear{abpar}, Kudritzki et al. \\citeyear{kupapu})}. Although the most abundant elements in the observed universe, hydrogen and helium mostly contribute to the wind density, their contribution to the radiative acceleration is very small (e.g. Abbott \\citeyear{abpar}). Heavier elements like carbon, nitrogen, oxygen or iron {are much more important for the wind acceleration \\zm{of present hot stars}. This is because heavier elements} have {effectively} much more lines that are available to absorb the stellar radiation and to accelerate the stellar wind. Clearly, in most cases the radiative force shall be higher for higher metallicity. Consequently, the basic wind properties shall depend on {the} metallicity.  Although relatively simple expressions for the dependence of the mass-loss rate and {the} terminal velocity on the metallicity can be obtained even on the basis of the standard CAK (Castor, Abbott \\& Klein \\citeyear{cak}) theory (e.g. Puls et al. \\citeyear{puspo}), the detailed dependence is probably more complicated (Puls et al. \\citeyear{puspo}, Vink et al. \\citeyear{vikolamet}). It is clear that more precise (NLTE) wind models are necessary to study this problem in detail.  {There are several independent NLTE wind models that can be used for the study of the metallicity dependence of the wind properties. These codes differ in the level of sophistication of the treatment of the wind problem. One of the most important aspects that influences the reliability of individual  NLTE wind models is the inclusion of proper bound-free and bound-bound blanketing in the UV domain. While models of Vink et al. (\\citeyear{vikolamet},  hereafter VKL) use Monte Carlo method for the solution of the radiative transfer equation, models of  Pauldrach et al. (\\citeyear{pahole}) solve this equations in detail. Wind models of Gr\\\"afener \\& Hamann (\\citeyear{graham}), that were successfully used for the calculation of wind models of WR stars, use comoving-frame formulation of the radiative transfer equation. All these models to some extent properly take into account blanketing effects in the UV domain that are important for the correct calculation of the ionization balance and of the radiative force. Hence, these models are able to reliably predict the most important wind parameter -- the  mass-loss rate and its metallicity dependence.  However, also the wind velocity field and especially the terminal velocity can be influenced by the metallicity. To study this effect, the solution of hydrodynamic equations is necessary. Since models of VKL, that are widely used in hot star evolutionary calculations, do not solve these equations, these models assume a prespecified velocity law and cannot be used for the predictions of the velocity structure. This is an important difference between models of VKL and e.g.~Pauldrach et al. (\\citeyear{pahole}) or Gr\\\"afener \\& Hamann (\\citeyear{graham}). Hydrodynamical wind models for different metallicities were calculated by Kudritzki (\\citeyear{kudmet}). According to these models the wind terminal velocity decreases with decreasing metallicity, in agreement with SMC observations (see an extensive compilation of Kudritzki \\& Puls \\citeyear{kupul}).  However, the wind metallicity does not vary the radiative force only. The momentum acquired by the heavier elements from the radiative field is transferred to the bulk wind component (i.e.~hydrogen and helium) via the Coulomb collisions. Since the frictional force caused by these collisions depends on number densities of heavier elements and hydrogen-helium component, the effects induced by the collisions are more important in the low-metallicity environment (Kudritzki \\citeyear{kudmet}, Krti\\v cka et al. \\citeyear{gla}). NLTE models available to study this problem were presented by  Krti\\v{c}ka \\& Kub\\' at (\\citeyear{nltei}, hereafter KK1).  Some hot stars with low luminosities (and consequently low mass loss rates) show signatures of very weak winds, much weaker than that deduced from standard wind theory (Bouret et al. \\citeyear{bourak}; Martins et al. \\citeyear{martin,okali}). Since this discrepancy is based mostly on wind models of Vink et al. (\\citeyear{vikolamet}), it would be interesting to \\zm{test} this result also with independent NLTE wind models and to test whether this discrepancy can be caused by multicomponent effects.}  {T}he study of the metallicity dependence of stellar wind properties is not important only for our understanding of stellar winds or {massive} stars itself. Since stellar winds play a role in the chemical enrichment of galaxies and are responsible for input of a large amount of momentum and energy in the interstellar medium, the detailed knowledge of variations of basic wind parameters with metallicity is important also for other fields of astrophysics.  As we have discussed, wind properties of hot stars shall depend on the stellar metallicity. {L}ow-metallicity environment of SMC offers a unique possibility to test these available predictions. To do so, we present here wind models for selected SMC hot stars. ",
        "conclusions": "We have presented NLTE wind models {of cooler} O stars in SMC. These models enable prediction of basic wind parameters, i.e. {the} mass-loss rates and {the} terminal velocities. We have compared these predicted wind parameters with those derived from observation. We have concluded that there is a relatively good agreement between observed and predicted terminal velocities with some scatter. This scatter can be partly attributed to the poorly known chemical composition of SMC stars. Moreover, the existence of unidentified binaries in our sample can also influence the comparison. Last but not least, the simplification involved in our code (e.g.~{simplified treatment of radiative transfer equation}\\zm{, like neglect of line overlaps}) can also influence the results. {However, t}he situation with mass-loss rates is different. For relatively high mass-loss rates (i.e.~$\\dmdt\\gtrsim10^{-7}\\,\\smrok$) there is again a good agreement between predicted quantities and those derived from observations. However, for smaller mass-loss rates the agreement is rather poor, predicted mass-loss rates are significantly higher than those derived from observation. Note that this is not an effect of our models only, since there is {a} similar disagreement between observations and models of {VKL}  (see B03 and {Mr04}).  We have tested whether the mentioned \\zm{systematic} disagreement {between theoretical and observ{ed} values of mass-loss rates} can be caused by the decoupling of some metals from the mean flow. We have shown that for stars with relatively good agreement between observed and predicted mass-loss rates the velocity differences are small (and thus there is no decoupling of wind components). On the other hand for stars for which the predicted mass-loss rates are systematically too high {either} the {instability connected with the} decoupling of wind components may occur {or} the frictional heating {increases wind temperature}. Note however that {only for one star} we have obtained the possibility of decoupling close to the stellar surface for velocities lower than the escape velocity. {(Note that in order to obtain fall back of the wind material and consequently lower wind mass-loss rate it is necessary to the decoupling occur below the point where wind velocity is equal to the escape velocity.)} {With this exception,} {presented} models are, to our knowledge, close to the star stable for perturbations of multicomponent flow. This means that we are not able without any further assumptions to explain \\zm{very} low {observed} mass-loss rate{s} of {most} studied stars {with $\\dmdt\\lesssim10^{-7}\\,\\smrok$} {(although in the outer wind regions of these stars the multicomponent effects are important)}. However, there are several possibilities how to (from the theoretical point of view) obtain decoupling of wind components for  the velocities lower than the escape velocity also for other stars (like slightly different metallicity, overestimation of theoretical mass-loss rates\\zm{, neglect of some heavier elements} or improved treatment of the multicomponent flow). {Another way how to explain low wind mass-loss rates may be connected with} the thin-wind {effect} (Puls et al. \\citeyear{puspo}, Owocki \\& Puls \\citeyear{owpu}).   The decoupling discussed here slightly differs from that studied so far (cf. Springmann \\& Pauldrach \\citeyear{treni}, KKII). {I}n the previous studies it was usually assumed that metals and passive component \\zm{may} decouple completely. {W}e present here a more detailed and more realistic picture of the decoupling of wind components. {W}e showed that some elements which have very low abundance, however which significantly contribute to the radiative force, may decouple \\zm{(or cause the instability connected with the decoupling)} {individually} from the mean flow, {that}  is basically not decoupled. {Moreover,} the decoupling may occur for higher mass-loss rates than it was previously assumed. {We present an approximate formula for the test of importance of discussed decoupling of individual heavier elements.}  {For our comparison we used mass-loss rates derived from the observation assuming smooth winds. If the O star winds are clumped, as is indicated by observations of e.g.~B03 or Martins et al. (\\citeyear{okali}), then the  mass-loss rates of O stars are likely lower than predicted by current hot star wind theory and we obtain a disagreement between observations and theory for all stars. On the other hand a better treatment of the opacity sources in the UV domain with account of the line overlaps \\zm{and comoving-frame radiative transport} may help to overcome this potential disagreement.}   We have studied the dependence of wind parameters on metallicity. It is well-known that the mass-loss rate increases with {increasing} metallicity. We have derived that the mass-loss  rate scales with metallicity as $\\dmdt\\sim Z^{0.67}$. The terminal velocity of individual stars also varies with metallicity mainly due to the {sensitivity of the radiative force on detailed wind state in outer wind regions}. {However, o}n the average the terminal velocity varies with metallicity only slightly (for studied values of metallicity) {as} $v_\\infty\\sim Z^{0.06}$. As a consequence, the ratio of wind terminal velocity to the surface escape velocity $v_\\infty/v_\\text{esc}\\sim2.3$ is only slightly lower than for Galactic stars."
    },
    "0601/astro-ph0601403_arXiv.txt": {
        "abstract": "We present the Fundamental Plane (FP) for 38 early-type galaxies in the two rich galaxy clusters RXJ0152.7--1357 ($z=0.83$) and RXJ1226.9+3332 ($z=0.89$), reaching a limiting magnitude of $M_B =-19.8\\,{\\rm mag}$ in the rest frame of the clusters. While the zero point offset of the FP for these high redshift clusters relative to our low redshift sample is consistent with passive evolution with a formation redshift of $z_{\\rm form}\\approx 3.2$, the FP for the high redshift clusters is not only shifted as expected for a mass-independent $z_{\\rm form}$, but rotated relative to the low redshift sample. Expressed as a relation between the galaxy masses and the mass-to-light ratios the FP is significantly steeper for the high redshift clusters than found at low redshift. We interpret this as a mass dependency of the star formation history, as has been suggested by other recent studies. The low mass galaxies ($10^{10.3}\\,M_{\\sun}$) have experienced star formation as recently as $z\\approx 1.35$ (1.5 Gyr prior to their look back time), while galaxies with masses larger than $10^{11.3}\\,M_{\\sun}$ had their last major star formation episode at $z > 4.5$. ",
        "introduction": "The Fundamental Plane (FP) for elliptical (E) and lenticular (S0) galaxies is a key scaling relation, which relates the effective radii, the mean surface brightnesses and the velocity dispersions in a relation linear in log-space (e.g., Dressler et al.\\ 1987; Djorgovski \\& Davis 1987; J\\o rgensen et al.\\ 1996, hereafter JFK1996). The FP can be interpreted as a relation between the galaxy masses and their mass-to-light (M/L) ratios. For low redshift cluster galaxies the FP has very low internal scatter, e.g.\\ JFK1996. It is therefore a powerful tool for studying the evolution of the M/L ratio as a function of redshift (e.g., J\\o rgensen et al.\\ 1999; Kelson et al.\\ 2000; van de Ven et al.\\ 2003; Gebhardt et al.\\ 2003; Wuyts et al.\\ 2004; Treu et al.\\ 2005; Ziegler et al.\\ 2005). These authors all find that the FP at $z$=0.2--1.0 is consistent with passive evolution of the stellar populations of the galaxies, generally with a formation redshift $z_{\\rm form} > 2$.  Most previous studies of the FP at $z$=0.2--1.0 cover fairly small samples of galaxies in each cluster and are limited to a narrow range in luminosities and therefore in masses, making it very difficult to detect possible differences in the FP slope. A few recent studies indicated a steepening of the FP slope for $z \\sim 1$ galaxies (di Serego Alighieri et al.\\ 2005; van der Wel et al.\\ 2005; Holden et al.\\ 2005). These studies and studies of the K-band luminosity function (Toft et al.\\ 2004) and the red sequence (de Lucia et al.\\ 2004) at $z\\approx 0.8-1.2$ suggest a mass dependency of the formation epoch.  We present the FP for two galaxy clusters RXJ0152.7--1357 at $z=0.83$ and RXJ1226.9+3332 at $z=0.89$. Our samples reach apparent $i'$-band magnitudes of 22.5--22.8 mag, equivalent to an absolute magnitude of $M_B = -19.8\\,{\\rm mag}$ in the rest frame of the clusters.  No other published samples suitable for studies of the cluster galaxy FP at $z>0.8$ go this deep. Our study of these two clusters is part of the Gemini/HST Galaxy Cluster Project, which is described in detail in J\\o rgensen et al.\\ (2005). We adopt a $\\Lambda$CDM cosmology with $H_0 = 70\\,{\\rm km\\,s^{-1}\\,Mpc^{-1}}$, $\\Omega_M=0.3$, and $\\Omega_{\\rm \\Lambda}=0.7$. ",
        "conclusions": "We find that the FP for E/S0 galaxies in the clusters RXJ0152.7--1357 ($z=0.83$) and RXJ1226.9+3332 ($z=0.89$) is offset and rotated relative to the FP of our low redshift comparison sample of Coma cluster galaxies. Expressed as a relation between the M/L ratios and the masses of the galaxies, the high redshift galaxies follow a significantly steeper relation than found for the Coma cluster. We interpret this as due to a mass dependency of the epoch of the last major star formation episode. The lowest mass galaxies in the sample ($10^{10.3}\\,M_{\\sun}$) have experienced significant star formation as recent as $z_{\\rm form} \\approx 1.35$, while high mass galaxies (${\\it Mass} > 10^{11.3}\\,M_{\\sun}$) have $z_{\\rm form} > 4.5$. This is in general agreement with the predictions for the star formation histories of E/S0 galaxies from Thomas et al.\\ (2005) based on their analysis of line index data for nearby galaxies. The scatter of FP for these two $z=0.8-0.9$ clusters is as low as found for the Coma cluster, and we find no significant difference in the scatter for low and high mass galaxies. This indicates that at a given galaxy mass the star formation history for the E/S0 galaxies is quite similar. In a future paper we will discuss these results in connection with our absorption line index data for the galaxies in both high redshift clusters.  \\vspace{0.5cm} Based on observations obtained at the Gemini Observatory (GN-2002B-Q-29, GN-2004A-Q-45), which is operated by AURA, Inc., under a cooperative agreement with NSF on behalf of the Gemini partnership: NSF (US), PPARC (UK), NRC (Canada), CONICYT (Chile), ARC (Australia), CNPq (Brazil) and CONICET (Argentina). Based on observations made with the NASA/ESA Hubble Space Telescope. IJ, KC, and KF acknowledge support from grant HST-GO-09770.01 from STScI. STScI is operated by AURA, Inc. under NASA contract NAS 5-26555."
    },
    "0601/astro-ph0601259_arXiv.txt": {
        "abstract": "{We investigated the role of the cyclotron emission (CE) associated to cosmic magnetic fields (MF) on the evolution of cosmic microwave background (CMB) spectral distortions. We computed the photon and energy injection rates by including spontaneous and stimulated emission and absorption. These CE rates have been compared with those of bremsstrahlung (BR) and double Compton scattering (DC), for realistic CMB distorted spectra at various cosmic epochs. For reasonable MF strengths we found that the CE contribution to the evolution of the CMB spectrum is much smaller than the BR and DC contributions. The constraints on the energy exchanges at various redshifts can be then derived, under quite general assumptions, by considering only Compton scattering (CS), BR, and DC, other than the considered dissipation process. Upper limits to the CMB polarization degree induced by CE have been estimated. ",
        "introduction": "\\label{sec:intro}  The CMB spectrum emerged from the thermalization redshift, $z_{therm}$, with a shape very close to a blackbody (BB) one, owing to the tight coupling between radiation and matter through CS and photon production/absorption processes, DC and BR. In the presence of a cosmic MF, another photon production/absorption process, CE, operates in the cosmic plasma (Puy \\& Peter 1998). Since the very large electron conductivity the magnetic flux through any loop moving with fluid is a conserved quantity. On scale where diffusion can be neglected the field is said to be % {\\it frozen-in}, in the sense that lines of force move together with % the fluid. Assuming that the Universe expands isotropically, magnetic flux conservation implies ${\\bf B}(t)={\\bf B}(t_0)\\left(\\frac{a(t_0)}{a(t)}\\right)^2={\\bf B}_0(1+z)^2 \\, $, ${\\bf B}_0$ being the MF at the present time. Considering the cyclotron spontaneous emission, absorption and stimulated emission terms (Afshordi 2002), we derived the CE contribution to the evolution of the CMB photon occupation number, $\\eta$, as a further term in the Kompaneets equation (Zizzo \\& Burigana 2005):  \\vskip -0.4cm \\begin{eqnarray}\\label{eq:kompaneets+ciclotrone} \\frac{\\partial\\eta}{\\partial t}&=&\\frac{1}{\\phi}\\frac{1}{t_C} \\frac{1}{x^2}\\frac{\\partial}{\\partial x}\\left[x^4\\left[\\phi \\frac{\\partial\\eta}{\\partial x}+\\eta(1+\\eta)\\right]\\right] \\nonumber\\\\ &+&\\left[K_{BR}\\frac{g_{BR}}{x_e^3}e^{-x_e}+K_{DC}\\frac{g_{DC}}{x_e^3}+K_{CE} \\delta(x_e-x_{e,CE})\\right] \\nonumber\\\\ &\\cdot&\\left[1-\\eta(e^{x_e}-1)\\right] \\, ; \\nonumber \\end{eqnarray} \\vskip -0.3cm  \\noindent here the BR and DC rates are defined by the coefficients $K(z)$ and the Gaunt factors, $g_{BR}$ and $g_{DC}$, $t_C=mc^2/[kT_e(n_e \\sigma_T c)]$ is the timescale for the achievement of kinetic equilibrium between radiation and matter, $\\phi=T_e/T_r$ where $T_r=T_0(1+z)$ is the CMB temperature and $T_e$ the electron one, $x=\\phi x_e$ (see, e.g., Burigana, De Zotti \\& Danese 1995 and references therein), $x_{e,CE}={h\\nu_c \\over kT_e} \\simeq 4.97\\times 10^{-5}\\cdot \\phi^{-1}\\tsu27 ^{-1}B_0(1+z) \\, $ is the dimensionless frequency of CE, and $K_{CE}(z)=\\frac{4\\pi^2\\,e\\,c}{3 B(z)}\\,n_e=4.64\\times10^{-4}\\, \\Ohat\\cdot B^{-1}(z)(1+z)^3\\mbox{ s}^{-1} \\, $ where $\\Ohat =\\Omega_b [H_{0}/(50 {\\rm Km} {\\rm s}^{-1} {\\rm Mpc}^{-1})]^2 \\,$, $\\Omega_b$ being the baryon density parameter and $H_0$ the Hubble constant.  In Fig.~1 we compare $x_{e,CE}$ with the frequency $x_{e,c}$, at which the combined DC and BR absorption time $t_{abs}$ is equal to $t_C$ and the maximum frequency $x_{e,abs}$, at which DC and BR could re-establish a BB spectrum after a distortion.  \\begin{figure} \\vskip -0.3cm \\resizebox{\\hsize}{!} {\\includegraphics[]{buriganaf1.ps}} \\vskip -0.3cm \\caption{Comparison between the characteristic frequencies $x_{e,c}$ (thick solid line) $x_{e,abs}$ (thin solid line), and $x_{e,CE}$ for different values of the MF $B_0$: $10^{-9}$G (thick dashed line), $10^{-8}$G (thick dot-dashed line), $10^{-7}$G (thick three dots and dash line). Different values of the fractional energy $\\Delta\\varepsilon/\\varepsilon_i$ ($\\simeq \\mu_0/1.4$ for $\\mu_0 \\ll 1$), injected in the CMB radiation field in the case of early (BE-like) distorted spectra, have been assumed in the panels: $10^{-4}$ (panel {\\it a}), $10^{-2}$ (panel {\\it b}), $1$ (panel {\\it c}). The cosmological parameters have been assumed in agreement with WMAP (see Tab.~3 of Bennett et al. 2003); nevertheless the detailed choice of them is not critical here.} \\label{fig1} \\end{figure} ",
        "conclusions": ""
    },
    "0601/astro-ph0601290_arXiv.txt": {
        "abstract": "We study the effect of the stochastic character of supernova explosions on the anisotropy of galactic cosmic rays below the knee. We conclude that if the bulk of cosmic rays are produced in supernova explosions the observed small and nearly energy independent amplitude of the anisotropy and its phase are to the large extent determined by the history of these explosions in the vicinity of the solar system, namely by the location and the age of the supernova remnants, within a few kpc, which give the highest contibution to the total intensity at the present epoch. Among the most important factors which result in the small magnitude and the energy independence of the anisotropy amplitude are the mixed primary mass composition, the effect of the Single Source and the Galactic Halo. Special attention is given to the phase of the anisotropy. It is shown that the excessive cosmic ray flux from the Outer Galaxy can be due to the location of the Solar System at the inner edge of the Orion Arm which has the enhanced density and rate of supernova explosions. ",
        "introduction": "Since the discovery of cosmic rays (~CR~) the study of their arrival directions is one of the most popular studies aiming eventually to discover their origin. The difficulty of this study is an extremely uniform distribution of CR arrival directions. The anisotropy amplitude in the energy range from tens of GeV to sub-PeV is less than 10$^{-3}$ and does not show any significant energy dependence (~see the recent summary in \\cite{Guil} and references therein~). It is surprising since the amplitude $A$ is usually connected with the diffusion coefficient $D$ as $A = \\frac{3DgradI}{cI}$, where $I$ is the CR intensity and $c$ is the speed of light. Since the diffusion coefficient is admitted to rise with the energy \\cite{Gais} as $E^s$ with $s \\approx 0.3-0.6$, then one would expect a corresponding rise of the anisotropy, which in fact is not observed \\cite{Hill}.  The view of the majority of CR workers about CR origin is that at energies below a PeV CR are produced mainly by the shocks from supernova (~SN~) explosions. These explosions are a stochastic process and the importance of the fluctuations in the cosmic ray production and propagation has been underlined in many works \\cite{Jones,Lee,Berez,Pohl}. The dominant role in these fluctuations is played by explosions of nearby and recent SN \\cite{Ahar,Nish,EW1,EW2}. The anomalous diffusion in the non-uniform interstellar medium (~ISM~) magnifies these fluctuations compared with normal 'gaussian' diffusion \\cite{Lag1,EW3}. Attempts to explain the observed anisotropy by the random character of SN explosions have been made in \\cite{Ptus,Svesh} and by the anomalous diffusion - in \\cite{Lag2}. In the present paper we analyse further the anisotropy of CR in the framework of scenario with the stochastic nature of SN explosions and the subsequent anomalous diffusion through the ISM of our Galaxy. Special attention is devoted to the observed phase of the anisotropy. ",
        "conclusions": "We have shown that if the bulk of the Cosmic Radiation is produced by shocks from SN explosions, the observed characteristics of their anisotropy are the product of local phenomena. These explosions are a stochastic process and the observed characteristics appear to some extent as accidental. The small magnitude and the nearly energy independent amplitude of the anisotropy can be the consequence of the mixed mass composition, the influence of the Single Source - the nearby and recent SN explosion, the smoothing effect of the Galactic Halo and the position of the Sun on the inner edge of the Orion arm. In order to unveil the specific spatial and temporal distribution of nearby SNR and other local objects which determine the characteristics of the observed CR one should study the energy dependence of the anisotropy amplitude and phase for separate primary nuclear charges.  {\\bf Acknowledgments}  The Royal Society and the University of Durham are thanked for financial support."
    },
    "0601/astro-ph0601545_arXiv.txt": {
        "abstract": "{ The high-frequency peaked BL Lac PG\\,1553+113 was observed in 2005 with the H.E.S.S. stereoscopic array of imaging atmospheric-Cherenkov telescopes in Namibia. Using the H.E.S.S. standard analysis, an excess was measured at the 4.0$\\sigma$ level in these observations (7.6 hours live time).  Three alternative, lower-threshold analyses yield $>$5$\\sigma$ excesses. The observed integral flux above 200 GeV is (4.8$\\pm$$1.3_{stat}$$\\pm$$1.0_{syst}$)$\\times$10$^{-12}$ cm$^{-2}$\\,s$^{-1}$, and shows no evidence for variability. The measured energy spectrum is characterized by a very soft power law (photon index of $\\Gamma$=4.0$\\pm$0.6).  Although the redshift of PG\\,1553+113 is unknown, there are strong indications that it is greater than $z$=0.25 and possibly larger than $z$=0.78.  The observed spectrum is interpreted in the context of VHE $\\gamma$-ray absorption by the Extragalactic Background Light, and is used to place an upper limit on the redshift of $z$$<$0.74. ",
        "introduction": "PG\\,1553+113 was discovered in the Palomar-Green survey of UV-excess stellar objects \\cite{discovery_paper}.  It is classified as a BL Lac object based on its featureless spectrum (\\cite{redshift_initial, redshift_question}) and its significant ($m_p$ = 13.2-15.0) optical variability \\cite{optical_variability}. PG\\,1553+113 is well studied from the radio to the X-ray regime, and has been the subject of several multi-wavelength observation campaigns.  In X-rays it has been detected by multiple observatories, with energy spectra measured by both {\\it BeppoSAX} \\cite{BeppoSAX_det} and {\\it XMM Newton} \\cite{XMM_det}. Based on its broad-band spectral energy distribution (SED), PG\\,1553+113 is now classified as a high-frequency peaked BL Lac \\cite{classification}, similar to essentially all of the AGN detected at VHE (very high energy; $>$100 GeV) energies.  The redshift of PG\\,1553+113 was initially determined to be $z$=$0.36$ (\\cite{redshift_initial}).  However, it was later claimed that this value is flawed as it was based on a spurious emission line, caused by a bright spot in the {\\it International Ultraviolet Explorer} image, that was misidentified as Lyman-$\\alpha$ (\\cite{redshift_question}). To date no emission or absorption lines have been measured despite more than ten observation campaigns with optical instruments, including two with the 8-meter VLT telescopes \\cite{no_lines}.  Estimates of the redshift can also be made using photometric observations. No host galaxy was resolved in {\\it Hubble Space Telescope} (HST) images of PG\\,1553+113 taken during the HST survey of 110 BL Lac objects \\cite{Hubble_image}. Approximately 80\\% of these BL Lacs (88) have known redshifts, of which all 39 with $z$$<$0.25 and 21 of the 28 with 0.25$<$$z$$<$0.6 have their hosts resolved \\cite{z_extreme}.  The HST results were recently used to set a lower limit of $z$$>$0.78 \\cite{z_extreme} for PG\\,1553+113. This redshift limit is based on the assumption of the nucleus being present in a typical host galaxy, and general correlation between photometrically determined redshifts and those determined spectroscopically. However, a departure from this correlation, or the possibility of an atypical host galaxy is not unexpected.  The possibility of such a large redshift is of critical importance to VHE  observations due to the absorption  of VHE photons (\\cite{EBL_effect, EBL_effect2}) by pair-production on the Extragalactic Background Light (EBL). This absorption, which is energy dependent and increases strongly with redshift, distorts the VHE energy spectra observed from distant objects, and may even render the VHE detection of very distant objects impossible.  Although the redshift of PG\\,1553+113 is essentially unknown, it is viewed as a promising candidate for detection as a VHE emitter \\cite{luigi_AGN} based on its SED. However, it has not been previously detected in the VHE regime. Only flux level upper limits have been reported by the Whipple collaboration \\cite{whipple_UL} and the Milagro group \\cite{Milagro_AGN}. Here, evidence for $>$200 GeV $\\gamma$-ray emission from PG\\,1553+113 is presented, based on observations made with the H.E.S.S. (High Energy Stereoscopic System) experiment. The data are used to set an upper limit on the source redshift. ",
        "conclusions": "\\begin{figure} \\centering \\includegraphics[width=8.7cm]{figures/Hl143_fig3.eps} \\\\ [-0.3cm] \\caption{The intrinsic photon index of PG\\,1553+113, with 1$\\sigma$ statistical error contours, versus redshift for the cases of {\\it minimal} and {\\it maximal} EBL absorption. An intrinsic photon index below the horizontal line ($\\Gamma_{\\mathrm{int}}$=1.5) is not considered realistic.  The uppermost (dashed-dotted) curve is the sum of the intrinsic photon index and 2$\\sigma$ statistical error for the case of minimal EBL absorption.} \\label{index_vs_z_plot} \\end{figure}  PG\\,1553+113 possibly represents the most distant BL Lac object detected at VHE energies. Therefore, the observed VHE spectrum is expected to be affected by EBL absorption. Unfortunately the unknown redshift of this object does not allow its VHE spectrum to be used in constraining the poorly-measured EBL flux density (see e.g. \\cite{HESS_EBL_paper} 2005d). However, assuming a redshift allows the determination of the spectrum intrinsic to the BL Lac (i.e. with the effect of EBL absorption removed). The intrinsic flux is related to the observed flux by F$_{\\mathrm{int}}$\\,=\\,e$^{\\tau(z,E)}$\\,F$_{\\mathrm{obs}}$, where $\\tau(z,E)$ is the optical depth due to pair production ($\\gamma_{_{\\mathrm{VHE}}}$\\,$\\gamma_{_{\\mathrm{EBL}}}$\\,$\\rightarrow$\\,$e^{+}$\\,$e^{-}$). Assuming that the main uncertainty arises from the overall normalization of the EBL density, this optical depth can be expressed as $\\tau(z,E)$\\,=\\,$f$\\,$\\tau_{0}(z,E)$, where $\\tau_{0}(z,E)$ is calculated using a reference EBL-density model, and $f$ is a scaling factor for this density. At energies relevant to the H.E.S.S. observations, $\\tau(z,E)$ increases with both redshift and energy.  Therefore, a larger redshift, or higher EBL flux density, causes a stronger absorption of higher-energy $\\gamma$-rays, resulting in a harder intrinsic spectrum (F$_{\\mathrm{int}}$\\,$\\propto$\\,$E^{-\\Gamma_{\\mathrm{int}}}$). The EBL density model of \\cite{joel_EBL} is chosen here since it includes the effects of galaxy evolution which become important at redshifts greater than $\\sim$0.2. This model's fluxes are above recently published upper limits (\\cite{HESS_EBL_paper} 2005d). However, scaling its values by $f$=0.85 approximately traces these limits\\footnote{A different model and hence scaling factor is used in \\cite{HESS_EBL_paper} (2005d).} and is therefore considered a {\\it maximal} EBL flux density. Given that the EBL flux density is poorly-known, a {\\it minimal} EBL scenario is also considered. Here, the original model is scaled by $f$=0.6, approximately tracing the lower limits on the EBL flux density provided by galaxy counts \\cite{pozzetti}.  Using these parameterizations, the intrinsic spectrum of PG\\,1553+113 is calculated for different redshift values.  Figure~\\ref{index_vs_z_plot} shows the intrinsic photon index ($\\Gamma_{\\mathrm{int}}$) versus redshift. Assuming the intrinsic photon index of the BL Lac is not harder than $\\Gamma_{\\mathrm{int}}$=1.5 (see \\cite{HESS_EBL_paper} 2005d for arguments supporting this assumption), an upper limit on the redshift of PG\\,1553+113 of $z$$<$0.74 is determined using the {\\it minimal} EBL. Above this redshift, where $\\Gamma_{\\mathrm{int}}$=0.3$\\pm$0.6, the sum of $\\Gamma_{\\mathrm{int}}$ and twice its statistical error is harder than 1.5. Clearly, the redshift upper limit is smaller if a higher EBL flux density is assumed. This is the first time that a VHE observation is used to usefully constrain the redshift of an astrophysical object.  Future observations of PG\\,1553+113 should allow for a more precise determination of the photon index, possibly resulting in a stronger redshift constraint."
    },
    "0601/astro-ph0601598_arXiv.txt": {
        "abstract": "We present a holographic dark-energy model in which the Newton constant $G_{N}$ scales in such a way as to render the vacuum energy density a true constant. Nevertheless, the model acts as a dynamical dark-energy model since the scaling of $G_{N}$ goes at the expense of deviation  of concentration of dark-matter particles from its canonical form and/or of promotion of their mass to a time-dependent quantity, thereby making the effective equation of state (EOS) variable and different from $-1$ at the present epoch. Thus the model has a potential to naturally underpin Dirac's suggestion for explaining the large-number hypothesis, which demands a dynamical $G_{N}$ along with the creation of matter in the universe. We show that with the aid of observational bounds on the variation of the gravitational coupling, the effective-field theory IR cutoff can be strongly restricted, being always closer to the future event horizon than to the Hubble distance. As for the observational side, the effective EOS restricted by observation can be made arbitrary close to $-1$, and therefore the present model can be considered as a ``minimal'' dynamical dark-energy scenario. In addition, for nonzero but small curvature $(|\\Omega_{k0}| \\lsim 0.003 )$, the model easily accommodates a transition across the phantom line for redshifts $z \\lsim 0.2 $, as mildly favored by the data. A thermodynamic aspect of the scenario is also discussed. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601551_arXiv.txt": {
        "abstract": "We have used {\\em Chandra} to resolve the nearby 70~Oph (K0~V+K5~V) and 36~Oph (K1~V+K1~V) binary systems for the first time in X-rays. The LETG/HRC-S spectra of all four of these stars are presented and compared with an archival LETG spectrum of another moderately active K dwarf, $\\epsilon$~Eri.  Coronal densities are estimated from O~VII line ratios and emission measure distributions are computed for all five of these stars.  We see no substantial differences in coronal density or temperature among these stars, which is not surprising considering that they are all early K dwarfs with similar activity levels. However, we do see significant differences in coronal abundance patterns. Coronal abundance anomalies are generally associated with the first ionization potential (FIP) of the elements.  On the Sun, low-FIP elements are enhanced in the corona relative to high-FIP elements, the so-called ``FIP effect.''  Different levels of FIP effect are seen for our stellar sample, ranging from 70~Oph~A, which shows a prominent solar-like FIP effect, to 70~Oph~B, which has no FIP bias at all or possibly even a weak inverse FIP effect. The strong abundance difference exhibited by the two 70~Oph stars is unexpected considering how similar these stars are in all other respects (spectral type, age, rotation period, X-ray flux).  It will be difficult for any theoretical explanation for the FIP effect to explain how two stars so similar in all other respects can have coronae with different degrees of FIP bias. Finally, for the stars in our sample exhibiting a FIP effect, a curious difference from the solar version of the phenomenon is that the data seem to be more consistent with the high-FIP elements being depleted in the corona rather than a with a low-FIP enhancement. ",
        "introduction": "Extensive surveys by the {\\em Einstein} and {\\em ROSAT} X-ray satellites have shown that all late-type main sequence stars are surrounded by $T\\geq 10^{6}$~K coronae analogous to the more easily observed corona surrounding our own Sun \\citep[e.g.,][]{js04}. Although the mechanism or mechanisms responsible for heating these outermost stellar atmospheric regions to such high temperatures are not entirely understood, it is clear that coronal heating and coronal structure are controlled by magnetic activity.  On the Sun, magnetic fields are associated with a number of fascinating phenomena:  sunspots, flares, the solar wind, coronal mass ejections, prominences, etc.  Lack of spatial resolution makes most of these phenomena difficult to study on stars, so most of what we know about stellar magnetic activity comes solely by studying their global radiation.  Since X-rays are produced by stellar coronae, stellar X-ray luminosities have become one of the primary measures of coronal activity, and much work has been done to see how activity as measured by X-ray emission correlates with other stellar properties such as age and rotation.  A thorough review of X-ray studies of stellar coronae is provided by \\citet{mg04}.  X-ray spectroscopy provides more diagnostic information on the characteristics of stellar coronae than the X-ray luminosity. The numerous pulse height spectra obtained by satellites such as {\\em Einstein}, {\\em ROSAT}, and {\\em ASCA} provide some information about the X-ray spectra of many stars, but their very low spectral resolution ($R\\equiv \\lambda/\\Delta\\lambda \\lesssim 50$) severely limits the usefulness of such data.  In contrast, the grating spectra produced by the now defunct {\\em Extreme Ultraviolet Explorer} (EUVE), and by the current {\\em Chandra} and {\\em XMM-Newton} missions can resolve individual lines, providing a more complete arsenal of spectroscopic techniques for inferring coronal characteristics.  For example, such spectra allow more precise determinations of coronal emission measure distributions, which indicate the temperature distribution of the emitting plasma.  These spectra often contain density-sensitive lines that allow the measurement of coronal densities \\citep[e.g.,][]{jun02}.  Finally, elemental abundances in stellar coronae can be estimated from X-ray spectra.  Much attention has been focused on these abundances, since coronae often exhibit curious abundance patterns that depend on first ionization potential (FIP).  For the Sun and other stars of modest or low activity, elements with low FIP are often observed to have enhanced abundances relative to high FIP elements \\citep{uf00}.  This is the so-called ``FIP effect.''  However, coronae of very active stars are often found to behave differently, either having no FIP effect or even an inverse FIP effect where low FIP elements are depleted rather than enhanced in the corona relative to high FIP elements \\citep[e.g.,][]{ma03}.  The origin of these coronal abundance effects is an active area of theoretical study \\citep[e.g.,][]{jml04}.  We report here on X-ray spectra obtained by {\\em Chandra}'s Low Energy Transmission Grating (LETG), using the HRC-S detector.  Our new observations are of two moderately active K dwarf binaries, 70~Oph (K0~V+K5~V) and 36~Oph (K1~V+K1~V).  {\\em Chandra}'s superb spatial resolution allows these binaries to be resolved for the first time in X-rays, meaning that these two observations provide spectra of all four stars in the two binary systems.  In addition to these data, we also analyze archival LETG/HRC-S spectra of another moderately active K dwarf, $\\epsilon$~Eri (K1~V).  The properties of the five stars in our sample are summarized in Table~1.  Our goal is to compare the X-ray spectra of these five very similar K dwarfs with very similar coronal activity levels, in order to see whether these stars have identical coronal properties.  In the case of the 70~Oph and 36~Oph binaries, we can actually compare spectra of stars that are coevol, and therefore have identical ages, presumably identical photospheric abundances, and rotation rates that are observed to be nearly identical in both cases (see Table~1).  Another reason behind our choice of targets is that these are among the few solar-like stars with measured mass loss rates, which provides us the opportunity to compare coronal properties with the strengths of the stellar winds that emanate from these coronae.  The wind measurements, which are based on H~I Lyman-$\\alpha$ absorption from the interaction regions between the winds and the interstellar medium \\citep{bew05a}, suggest that despite their similarities the three stellar systems have significantly different mass loss rates.  With respect to the solar mass loss rate of $\\dot{M}_{\\odot}=2\\times 10^{-14}$ M$_{\\odot}$~yr$^{-1}$, the stellar mass loss rates range from $15~\\dot{M}_{\\odot}$ for 36~Oph up to $100~\\dot{M}_{\\odot}$ for 70~Oph (see Table~1).  (Unfortunately, the wind-ISM interaction regions that are used to measure the stellar winds encompass both stars in the 36~Oph and 70~Oph systems, so wind measurements can only be made for the combined winds of these binaries.)  We will see whether there are differences in the coronal X-ray spectra of these stars that may be connected with these differences in wind strength. ",
        "conclusions": "\\subsection{The 70 Oph Abundance Discrepancy}  Perhaps the most interesting result of our analysis is the coronal abundance difference between the two members of the 70~Oph binary, which is unexpected since these two stars are so similar in all other respects. Previous stellar observations have shown that the traditional solar-like FIP effect is generally observed for stars of low to moderate activity \\citep{jml96,jjd97,jml99}. % In contrast, for very active stars an inverse FIP effect is generally observed, where low-FIP elements have coronal abundances that are {\\em lower} relative to photospheric values than is the case for high-FIP elements \\citep{acb01,jjd01,mg01,dph01,jsf03} % Coronae of stars of intermediate activity are generally found to have abundances with little or no dependence on FIP \\citep{jjd95,ma01,bb05}. % Thus, the general picture is of the FIP effect being dependent on activity, with the low-FIP element abundances decreasing in the corona with increasing activity \\citep{ma03,at05}. %  {\\em However, if two stars as similar as 70~Oph~A and B can have drastically different levels of FIP bias, this makes it very difficult for any simple paradigm to explain why different stars have different coronal abundance patterns.}  One potential way out of this dilemma is to appeal to time dependent effects to explain the curious 70~Oph results.  Time variations in FIP bias are observed in solar active regions, which emerge with nearly photospheric abundances and then acquire a FIP bias that increases on a timescale of days \\citep{kgw01}.  Perhaps at the time of our {\\em Chandra} observation of 70~Oph, the visible part of 70~Oph~B's corona was dominated by young active regions while the visible part of 70~Oph~A's corona was dominated by old active regions, in which case the observed difference in FIP behavior is temporary.  The only way to test this hypothesis would be to reobserve 70~Oph, preferably several times.  \\citet{jsf03} looked for abundance variations in multiple observations of the very active K dwarf AB~Dor, with the observations encompassing a range of activity levels for the star.  However, no abundance variability was found.  \\subsection{Depletions of High-FIP Elements in Stellar Coronae}  Another puzzling result to arise from our coronal abundance analysis is the apparent depletion of high-FIP elements in the corona with respect to the photosphere, at least for a few of our stars.  Although the origin of the FIP effect and inverse FIP effect is not well understood, the cause is thought to lie in the chromosphere, where the low-FIP elements are ionized and therefore subject to forces that the neutral high-FIP elements are not subject to.  The general assumption is that these forces, whatever their nature, preferentially act on low-FIP elements, which leads to these elements having nonphotospheric abundances in the corona.  For the Sun and relatively inactive stars, the low-FIP coronal abundance trend is towards high abundances, while for very active stars the trend is towards low abundances.  In this simple scenario, the elements with the highest FIP should always have coronal abundances closest to the photospheric values since they are less subject to whatever forces are causing the FIP effect. This, however, is not consistent with what we see in our data.  For example, for every single one of our stars the coronal O abundance is farther from the photospheric abundance than the Fe and Si abundances, with the data suggesting a depletion of O in the corona in every case.   Note that measured photospheric O, Si, and Fe abundances exist for our stars and are taken into account in the analysis (see Table~1), so a lack of photospheric reference abundances is not an issue here, though it is often a problem in interpreting the coronal abundance measurements for very active stars \\citep[see, e.g.,][]{jsf04}. \\citet{at05} find similar evidence for high-FIP element depletions for many of the active solar-like stars in their sample.  Other examples of coronal abundance patterns that are also difficult for any simple FIP-based fractionation model to explain are reported by \\citet{acb01}, \\citet{jjd01}, \\citet{dph03}, and \\citet{jsf04}.  Another observational result that the simple FIP-dependent scenarios cannot easily explain is that the Ne/O ratios of many different stars are the same, regardless of whether the stars generally show a solar-like FIP effect, no FIP effect, or an inverse FIP effect.  All of our stars and all of those in the larger sample of \\citet{jjd05} suggest ${\\rm Ne/O}\\approx 0.4$.  Either the processes that cause the observed coronal abundance anomales are not dependent on FIP in any simple, monotonic fashion, or there are other fractionation forces at work that have nothing to do with FIP whatsoever.  Some solar flare observations have also suggested abundance patterns similar to that seen for 70~Oph~A in Figure~8, with high-FIP elements being depleted in the corona \\citep{njv81,af95}. The Ne/O ratios observed from solar active regions also show substantial variation about the mean solar value of 0.18, demonstrating a deviation from the average FIP effect even in relatively quiescent plasma \\citep{jts96}.  It has been proposed that photoionization by coronal X-rays can lead to nonequilibrium ionization conditions in the atmospheric regions where fractionation is taking place \\citep{as91}.  Since Ne, for example, has a higher photoionization cross section than O, Ne could be photionized while O remains neutral despite Ne having a higher FIP.  Analogous effects have been proposed to explain a factor of 2 depletion of He in the solar wind \\citep{rvs89,rvs00}.  For active stars, the irradiation of stellar atmospheres by coronal X-rays would be more intense, potentially making such photoionization effects more important.  This could perhaps explain the high Ne/O ratios seen for active stars, but it does not obviously explain other aspects of the observed abundance patterns, such as why high-FIP elements can be depleted in stellar coronae.  Clearly, we are still far from understanding the fractionation forces operating in stellar atmospheres, and far from understanding why they appear to operate so differently on different stars.  \\subsection{The Ne/O Controversy}  The whole issue of Ne/O ratios deserves additional comment since solar/stellar Ne and O abundances have recently become controversial topics.  The discord began when complex 3D hydrodynamic models of the solar atmosphere were used to reassess solar abundances, and it was then claimed that the abundances of many light elements such as C, N, and O needed to be revised downward by $25-35$\\% \\citep{cap01,cap02,ma05}.  Unfortunately, these suggested changes seriously damage the remarkably good agreement that exists between models of the solar interior and various helioseismological measurements, when older abundances are used \\citep{sb04,jnb05a}.  It has been noted that these problems can be mitigated by increasing the assumed solar Ne abundance by a factor of 2--3 \\citep{hma05,jnb05b}.  The reason why such an increase can be reasonably proposed is that the solar Ne abundance is poorly known.  There are no spectroscopic diagnostics of Ne in the photosphere, which means that the solar Ne abundance must be inferred from coronal or transition region measurements.  But this is dangerous given the existence of fractionation processes that result in differences between coronal and photospheric abundances (see \\S5.2).  When \\citet{jjd05} demonstrated that stellar Ne/O ratios have values consistent with ${\\rm Ne/O}=0.41$, which is significantly higher than the ratio assumed for the Sun (${\\rm Ne/O}=0.18$), they proposed that the stellar measurements represent a better estimate of the solar/stellar photospheric Ne/O ratios than the lower solar ratio.  Not only would this explain the remarkably consistent stellar Ne/O ratios, but it would also solve the discrepancies between solar interior models and helioseismology induced by the revised C, N, and O abundances.  Unfortunately, very recent reassessments of the solar coronal Ne/O ratio based on both full-disk and spatially resolved data seem to confirm that the solar Ne/O ratio is truly a factor of 2 lower than those observed from the stars in the \\citet{jjd05} sample, which are more active than the Sun \\citep{pry05,jts05}. Thus, the evidence at present seems to suggest that the Sun really does have a different coronal Ne/O ratio than active stars.  Furthermore, there are measurements of Ne/O for stars as inactive as the Sun that suggest low, solar-like ratios.  These measurements generally carry large uncertainties because the Ne~IX and Ne~X lines used to provide Ne abundances are generally weak for the cooler coronae of inactive stars.  Nevertheless, \\citet{ar02} quote a value of ${\\rm Ne/O}=0.21$ for Procyon, though \\citet{jsf04} find a higher value of ${\\rm Ne/O}=0.40$.  \\citet{ar03} report ${\\rm Ne/O}=0.18$ and ${\\rm Ne/O}=0.24$ for $\\alpha$~Cen~A and B, respectively, and \\citet{at05} find ${\\rm Ne/O}\\approx 0.14$ for $\\beta$~Com, though with large uncertainties.  A compilation of measurements in \\citet{mg04}, including those mentioned above, provides some evidence for a weak correlation of Ne/O with activity.  As mentioned in \\S4.4, the high Ne/O ratios seen for our K dwarfs are significant because they demonstrate that the high Ne/O ratio found by \\citet{jjd05} applies to stars less active than most of those in their sample.  But our K dwarfs are still significantly more active than the Sun.  There are two possible explanations for the different Ne/O ratios that exist for the Sun and more active stars:  (a) The stellar Ne/O measurements are representative of the true cosmic abundance and the low solar Ne/O measurements are a product of some fractionation mechanism operating in the solar atmosphere but not the active stars, or (b) The solar Ne/O measurement is characteristic of the true cosmic abundance and the high stellar ratios are due to a fractionation mechanism that is not present for the Sun.  At first glance, the latter interpretation seems more likely, partly because the low solar Ne/O abundances are known to exist down to at least transition region temperatures \\citep{pry05}, and partly because it is easier to imagine active stars having strong fractionation effects that result in substantial differences between coronal and photospheric Ne/O ratios. However, this would leave the solar interior problems unresolved.  Without direct photospheric measurements of Ne, it is not possible to say much about solar/stellar Ne abundances with any great amount of confidence.  Perhaps for Ne it is best to measure the cosmic abundance from a very different astronomical source.  We note that \\citet{gg04} have estimated O and Ne abundances in the local ISM using {\\em Ulysses} and {\\em ACE} measurements of pickup ions within our solar system, and they find ${\\rm Ne/O}\\approx 0.41$.  This is in better agreement with the active star ratio than the solar ratio, though it must be noted that the ISM measurements come with their own set of systematic uncertainties and assumptions regarding dust depletion, deflection at the heliopause, and ISM ionization fraction.  \\subsection{Do Winds Correlate with Coronal FIP Bias?}  A major motivation behind our {\\em Chandra} observations was to see whether coronal differences in our K dwarf sample could be correlated with observed differences in the strengths of their winds. However, the only coronal properties that clearly vary in this sample of stars are the abundances.  The overall coronal activity levels of these stars are about the same, and the {\\em Chandra} data cannot distinguish any significant density or temperature differences in the stellar coronae.  Can the observed abundance variations be correlated with wind strength?  The mass loss rates per unit surface area of 36~Oph, $\\epsilon$~Eri, and 70~Oph (in solar units) are 18.1, 49.3, and 85.7, respectively (see Table~1).  (Unfortunately, we only know the combined mass loss rates for the 36~Oph and 70~Oph binaries and not the contributions of the individual stars.)  The 36~Oph binary has the weakest wind, and both stars have coronae that show a modest FIP effect (see Fig.~8).  This suggests a potential connection between weak winds and strong FIP bias.  The stronger wind and weak-or-no FIP effect of $\\epsilon$~Eri would support such a connection, but the fact that 70~Oph consists of both a high FIP-bias star (70~Oph~A) and a no-or-inverse FIP-bias star (70~Oph~B) complicates the interpretation of its high mass loss measurement. If the FIP-bias/wind anticorrelation suggested by 36~Oph and $\\epsilon$~Eri is true, that would imply that 70~Oph~B must be responsible for most of 70~Oph's strong wind.  Small number statistics preclude any definitive conclusions on the nature of any FIP/wind connection, but we note that for the Sun different levels of FIP bias are seen for low and high speed streams \\citep{rvs00}. This suggests that a FIP/wind connection of some sort is not implausible, but it must be noted that our X-ray spectra will be characteristic of brighter, high density plasma still confined in coronal loops rather than outflowing wind material, so any FIP/wind connection would have to be indirect."
    },
    "0601/astro-ph0601284_arXiv.txt": {
        "abstract": "We present new observations of the rapid oscillations in the dwarf nova VW Hyi, made late in outburst. These dwarf nova oscillations (DNOs) increase in period until they reach 33 s, when a transition to a strong 1st harmonic and weak fundamental takes place. After further period increase the 2nd harmonic appears; often all three components are present simultaneously. This 1:2:3 frequency suite is similar to what has been seen in some neutron star and black hole X-Ray binaries, but has not previously been seen in a cataclysmic variable. When studied in detail the fundamental and 2nd harmonic vary similarly in phase, but the 1st harmonic behaves independently, though keeping close to twice the frequency of the fundamental. The fundamental period of the DNOs, as directly observed or inferred from the harmonics, increases to $\\sim$ 100 s before the oscillation disappears as the star reaches quiescence. Its maximum period is close to that of the `longer period' DNOs observed in VW Hyi. The quasi-periodic oscillations (QPOs), which have fundamental periods 400 -- 1000 s, behave in the same way, showing 1st and 2nd harmonics at approximately the same times as the DNOs. We explore some possible models. One in which the existence of the 1st harmonic is due to transition from viewing a single accretion region to viewing two regions, and the rate of accretion onto the primary is modulated at the frequency of the 1st harmonic, as in the `beat frequency model', can generate the suite of DNO frequencies observed. But the behaviour of the QPOs is not yet understood. ",
        "introduction": "Many dwarf novae show rapid (typical periods 5 -- 40 s) low amplitude brightness oscillations during outburst, and some nova-like variables show similar modulations. These have been actively studied since their discovery some thirty years ago (Warner \\& Robinson 1972) and are known collectively as dwarf nova oscillations (DNOs). They show a systematic period-luminosity relationship during outburst, with minimum period at maximum bolometric luminosity, and are of moderate coherence near maximum but usually appear to become increasingly incoherent near the return to quiescence. The DNO periods have been observed to range over a factor of more than two in some dwarf novae. The short periods indicate that the source of these modulations is close to the white dwarf primary in these cataclysmic variables (CVs); this is confirmed by the appearance of oscillations of substantial amplitude in soft X-Ray and EUV flux during the outbursts of, e.g., the nearby dwarf novae VW Hyi and SS Cyg. Evidently the mass transfer from the accretion disc to the primary is strongly modulated at the DNO period.  The DNOs were reviewed in Warner (1995a); updates are given in Warner (1995b), in the first two papers in this series (Woudt \\& Warner 2002: hereafter Paper I; Warner \\& Woudt 2002: hereafter Paper II) and in Warner (2004). A model in which the accretion flow close to the primary is controlled by the magnetic field generated in a rapidly rotating equatorial belt, formed as a result of the high rate of accretion, has been proposed (Warner 1995b; Paper II); this is basically an intermediate polar structure, including a truncated accretion disc, but where only the belt receives the accreting angular momentum; the belt is magnetically coupled to the accretion disc but is able to slip relative to the body of the white dwarf primary. This Low Inertia Magnetic Accretor (LIMA) model (which was first proposed by Paczynski (1978)) allows large variations in DNO periods, and also explains the rapid deceleration of the DNOs observed in VW Hyi near the end of its outburst (Papers I and II) -- from the effect of outward propellering of gas as the belt continues to revolve rapidly after the mass transfer rate ($\\dot{M}$) diminishes. It is probable that the inner edge of the truncated disc is close to the corotation radius (where the Keplerian period in the disc equals the rotation period of the primary), even when propellering is present, as argued by Rappaport, Fregeau \\& Spruit (2004). Furthermore, the rotation period of the field carried round by the equatorial belt will rarely be exactly the same as the period of the inner edge of the disc -- the belt is continually being spun up or down by its magnetic connection to the disc. This corresponds to the model with large $\\Theta$ in the 3D simulations of accretion from a disc onto an inclined dipole by Romanova et al.~(2004), which is the model in which they find rapid quasi-periodic luminosity variations. In these 3D models, unlike the simple 2D LIMA model, the field inflates, opens out and reconnects, but the process is still driven by differential rotation between star and disc (see Uzdensky (2004) for a review).  The optical DNOs are the result of soft X-Ray and EUV emission from the accretion zones creating a `beam' that sweeps around and is reprocessed to longer wavelengths by the surface of the disc or anything else it illuminates. Another DNO component can come from the accretion zone itself.  Independent deduction of the existence of a spinning belt in VW Hyi and in other dwarf novae has been obtained from HST and FUSE spectroscopic observations. Spectra taken after outburst can only be modeled by composite spectra, in the case of VW Hyi comprising a white dwarf with projected rotational velocity $v \\sin i$ in the range 300 -- 400 km s$^{-1}$ and a hotter component of small area and $v \\sin i$ 3000 -- 4000 km s$^{-1}$ (Sion et al.~1996; G\\\"ansicke \\& Beuermann 1996; Godon et al.~2004; Sion et al.~2004a,b,c). The spinning accretion belt lasts for at least 14 days after a superoutburst (Sion et al.~2004b).  There is a second kind of DNO with periods about 4 times those of the standard DNOs, with little or no dependence on accretion luminosity. These are termed longer period DNOs: lpDNOs (Warner, Woudt \\& Pretorius: hereafter Paper III; Pretorius 2004).  In addition to the two types of DNO, there are large amplitude modulations with periods about 15 times larger than those of the DNOs, and much shorter coherence, known as quasi-periodic oscillations (QPOs). In Paper II we proposed that these are caused by traveling waves excited in the inner disc -- reprocessing and obscuring radiation from the bright central regions of the disc. Furthermore, the occasional appearance of `double DNOs' with frequency difference equal to the QPO frequency are attributed to the DNO `beam' sweeping around and being partly reprocessed by the progradely traveling vertically thickened disc. We should therefore in general be prepared to distinguish between `direct' DNOs, which arise from reprocessing from the disc, and what we will call here `synodic' DNOs, which are reprocessed from some moving site (this is discussed in more detail in Section 3). Whereas the direct DNOs are observed to be sinusoidal, the synodic DNOs commonly possess significant amplitude at their first harmonic. Because the synodic DNOs are observed to be reprocessed from a progradely moving source we do not consider a magnetically warped and precessing inner disc, which moves retrogradely (Pfeiffer \\& Lai 2004).  In VW Hyi we observed evolution of both the DNO and the QPO periods but the ratio $R = P_{QPO}/P_{DNO}$ remained almost constant at $\\sim$ 15 (Paper I). It was pointed out that a similar ratio is seen in the high and low frequency QPOs observed in X-Ray binaries (Beloni, Psaltis \\& van der Klis 2002), and that in fact (identifying DNOs and QPOs in CVs as the equivalent of the two kinds of QPOs in X-Ray binaries) the behaviour of VW Hyi forms an extension two orders of magnitude lower in frequency of the two-QPO correlation seen in the latter. An extension to even lower frequencies in other CVs has since been observed (Mauche 2002; Paper III).  The importance of further observations of DNOs and QPOs in CVs, where they are more easily analysed than in the X-Ray binaries, is evident. The rich set of phenomena exhibited by VW Hyi (Papers I and II) makes this a prime target for further study. In Section 2 we report further high speed photometric measurements of VW Hyi made towards the end of outburst, supplementing those already presented in Paper I, and revealing the occurrence of harmonic 1:2:3 ratios of periods late in outburst. In Section 3 we consider possible interpretations in terms of magnetically controlled accretion; in Section 4 we compare our VW Hyi results with those seen in X-Ray binaries and in Section 5 we briefly summarize the current state of play.  \\begin{table*} \\centering \\caption{An overview of a section of our data archive of VW Hyi. Observations in the late stages of decline from outburst with DNOs present.} \\begin{tabular}{@{}lrcrclrrccccc@{}} Run        & Tel. & Filter & HJD Start  & $t_{in}$ & Outburst & $T = 0^*$  & Length  & Range $T$  & \\multicolumn{3}{c}{DNOs$^{**}$} & QPO  \\\\ &      &        & +244\\,0000 & (sec)     &Type$^\\dag$     & (days) & (hours) & (days) &     &     &  &     \\\\[10pt] S0018 & 20-in & -- &  1572.49752 & 2 &  Normal (L) &  1572.75 &  2.04 &--0.25 $\\rightarrow$ --0.17& F &  & &  -- \\\\ S6133 & 74-in & I  & 11785.54592 & 2 & Normal (L) & 11785.75 &  2.84 &--0.20 $\\rightarrow$ --0.09 & F &  & &  -- \\\\ S6184 & 40-in & -- & 11957.27611 & 5  &Normal (L) & 11957.25 &  1.68 & 0.03 $\\rightarrow$ 0.10   & F &   & & $\\surd$ \\\\ S6059 & 40-in & -- & 11580.27697 & 4 & Normal (M) & 11580.20 &  5.27 & 0.08 $\\rightarrow$ 0.30   & F &   & & $\\surd$ \\\\ S0127 & 40-in & -- &  1677.29328 & 4 & Super      &  1677.20 &  3.77 & 0.09 $\\rightarrow$ 0.25   & F &   & & $\\surd$ \\\\ S7621 & 40-in & -- & 13463.22526 & 6 & Normal (L) & 13463.10 &  4.03 & 0.13 $\\rightarrow$ 0.34   & F & 1 & &      \\\\ S2915 & 40-in & -- &  4937.28242 & 4 & Normal (L) &  4937.10 &  0.73 & 0.18 $\\rightarrow$ 0.21   &   & 1 & &  -- \\\\ S1307 & 40-in & -- &  2354.38968 & 4 & Normal (M) &  2354.20 &  1.33 & 0.19 $\\rightarrow$ 0.25   &   & 1 & & $\\surd$  \\\\ S7311 & 74-in & B  & 13139.19081 & 4 & Normal (M) & 13138.95 &  5.41 & 0.24 $\\rightarrow$ 0.47   &   & 1 & 2 & $\\surd$ \\\\ S7342 & 40-in & -- & 13155.19077 & 5 & Normal (M) & 13154.90 &  3.07 & 0.29 $\\rightarrow$ 0.42   & F & 1 & 2 & $\\surd$ \\\\ S6316 & 40-in & -- & 12354.23650 & 5 & Normal (L) & 12353.95 &  5.72 & 0.29 $\\rightarrow$ 0.53   &   & 1 &   & $\\surd$ \\\\ S6138 & 40-in & -- & 11898.27800 &4,5& Normal (M) & 11897.85 &  7.63 & 0.43 $\\rightarrow$ 0.75   & F & 1 &   & $\\surd$ \\\\ S7384 & 74-in & -- & 13234.64681 & 4  &Normal (S) & 13234.20 &  0.78 & 0.45 $\\rightarrow$ 0.48   &   & 1 &   &  --  \\\\ S3416 & 40-in & -- &  5967.55248 & 2 & Normal (M) &  5967.00 &  1.19 & 0.55 $\\rightarrow$ 0.60   &   & 1 &   & $\\surd$ \\\\ S6528 & 74-in & -- & 12520.21754 & 4 & Normal (L) & 12519.65 & 10.20 & 0.57 $\\rightarrow$ 0.99   &   & 1 & 2 &  --  \\\\ S5248 & 40-in & -- &  8202.36304 &4,5& Normal (M) &  8201.75 &  4.66 & 0.61 $\\rightarrow$ 0.81   &   & 1 &   &  --  \\\\ S0019 & 20-in & -- &  1573.44068 & 5 & Normal (L) &  1572.75 &  4.30 & 0.69 $\\rightarrow$ 0.87   &   & 1 & 2 & $\\surd$ \\\\ S2623 & 40-in & -- &  3515.30484 &4,5& Normal (M) &  3514.60 &  3.60 & 0.70 $\\rightarrow$ 0.85   &   & 1 & 2 & $\\surd$ \\\\ S7343 & 40-in & -- & 13155.65679 & 5 & Normal (M) & 13154.90 &  0.21 & 0.76 $\\rightarrow$ 0.77   & F &   & 2 &  --  \\\\ S0484 & 40-in & -- &  2023.30835 & 4 & Super      &  2022.55 &  3.91 & 0.76 $\\rightarrow$ 0.92   & F & 1 &   & $\\surd$ \\\\ S7222 & 74-in & B  & 13007.27570 & 4 & Super      & 13006.40 &  7.54 & 0.88 $\\rightarrow$ 1.19   &   & 1 & 2 & $\\surd$ \\\\ S0129 & 40-in & -- &  1691.29620 & 5 & Super      &  1690.20 &  1.86 & 1.10 $\\rightarrow$ 1.18   & F &   &   &  --  \\\\ S7368 & 74-in & -- & 13169.48671 & 4 & Normal (M) & 13168.35 &  2.42 & 1.14 $\\rightarrow$ 1.24   & F &   & 2 & $\\surd$ \\\\ S7301 & 40-in & -- & 13087.24153 & 5 & Normal (M) & 13086.10 &  2.94 & 1.14 $\\rightarrow$ 1.26   &   &   & 2 & $\\surd$ \\\\ S1322 & 40-in & -- &  2355.38508 & 4 & Normal (M) &  2354.20 &  4.91 & 1.19 $\\rightarrow$ 1.39   &   &   & 2 & $\\surd$ \\\\[5pt] \\end{tabular} {\\footnotesize \\newline $^*$ The time $T = 0$ is based on a template light curve of VW Hyi as defined in Paper I (see figure 5 of Paper I).\\\\ $^{**}$ F: Fundamental, 1: First harmonic, 2: Second harmonic. $^{\\dag}$ S = short, M = medium, L = long.\\hfill \\\\ } \\label{dno4tab1} \\end{table*} ",
        "conclusions": "VW Hyi continues to provide unprecedented behaviour in its rapid variations. No other dwarf nova is yet known that shows such a range of phenomena; whether this is connected with the sustained and increasing prominence of oscillations late in outburst is not known -- no other dwarf nova has shown this either. The rapid increase in period of the DNOs, at a rate greater than that observed in other declining dwarf novae, which we have attributed to a phase of propellering and deceleration of an equatorial accretion belt, has been shown here to be accompanied by the appearance of first and second harmonics in both DNOs and QPOs. These are again unique to VW Hyi as a dwarf nova, but bear similarities to harmonics seen in the QPOs of X-Ray binaries.  Whereas most of the behaviour of the DNOs can be understood in terms of accretion from a disc onto a freely rotating equatorial accretion belt, that of the QPOs has not been fully explained.  Tests of the models suggested in this series of papers could be made from measuring the amplitude and phase changes of DNOs and QPOs observed in eclipsing CVs, and from detection of polarization modulated at the periods of the DNOs. These will require large telescopes to provide the necessary good signal/noise."
    },
    "0601/astro-ph0601417_arXiv.txt": {
        "abstract": "We present some preliminary (half-way) results on our adaptive optics spectroscopic survey of AGN at spatial scales down to 0.085$''$. Most of the data were obtained with SINFONI which provides integral field capability at a spectral resolution of $R\\sim4000$. The themes on which we focus in this contribution are: star formation around the AGN, the properties of the molecular gas and its relation to the torus, and the mass of the black hole. ",
        "introduction": "\\label{dav:sec:sample}  The primary criteria for selecting AGN were that (1) the nucleus should be bright enough for adaptive optics correction, (2) the galaxy should be close enough that small spatial scales can be resolved, and (3) the galaxies should be ``well known'' so that complementary data can be found in the literature. These criteria were not applied strictly, since some targets were also of particular interest for other reasons. The resulting sample of 9 AGN is listed in Table~\\ref{dav:tab:sample}. The observations of these are now completed, and while the data for some objects has been fully analysed, others are still in a preliminary stage. Additional AGN will likely be added once the Laser Guide Star Facility is commissioned.  \\begin{table} \\begin{centering} \\caption{AGN sample} \\label{dav:tab:sample}       % \\begin{tabular}{llrlrll} \\hline\\noalign{\\smallskip} Target & Classification & Dist. & \\ \\ \\ & \\multicolumn{3}{c}{Observations} \\\\ &                & (Mpc)    && Date & \\ \\ \\ & Instrument\\\\ \\noalign{\\smallskip}\\hline\\noalign{\\smallskip} Mkn 231$^1$  & ULIRG / Sy 1 / QSO & 170 && May '02 && Keck / NIRC2 \\\\ NGC 7469$^2$ & Sy 1 & 66 && Nov '02 && Keck / NIRSPAO \\\\ IRAS 05189-2524 \\ \\ & ULIRG / Sy 1 & 170 && Dec '02 && VLT / NACO \\\\ Circinus$^3$ & Sy 2 & 4 && Jul '04 && VLT / SINFONI \\\\ NGC 3227$^4$ & Sy 1 & 17 && Dec '04 && VLT / SINFONI \\\\ NGC 3783 & Sy 1 & 42 && Mar '05 && VLT / SINFONI \\\\ NGC 2992 & Sy 1 & 33 && Mar '05 && VLT / SINFONI \\\\ NGC 1068 & Sy 2 & 14 && Oct '05 && VLT / SINFONI \\\\ NGC 1097 & LINER / Sy 1 & 18 && Oct '05 && VLT / SINFONI \\\\ \\noalign{\\smallskip}\\hline \\end{tabular} \\end{centering} $^1$ Davies et al. 2004a \\cite{dav:dav04a}; $^2$ Davies et al. 2004b \\cite{dav:dav04b}; $^3$ M\\\"uller Sanchez et al. 2006 \\cite{dav:mul06}; $^4$ Davies et al. 2006 \\cite{dav:dav06}; \\end{table}  \\begin{figure} \\centering \\includegraphics[height=4.5cm]{davies_n2992.eps} \\caption{Spectrum of the central 0.5$''$ of NGC\\,2992, taken with SINFONI at a spatial resolution of 0.3$''$ and a spectral resolution of $R\\sim3400$. The stellar absorption features are clear, as is the coronal [Ca\\,{\\sc VIII}], the H$_2$ 1-0\\,S(1), and both narrow and broad Br$\\gamma$.} \\label{dav:fig:n2992}       % \\end{figure}  One immediate result, which has a bearing on the classifications in the table, is the frequent detection of broad Br$\\gamma$ -- i.e. with FWHM at least 1000\\,km\\,s$^{-1}$. An example of this is given in Fig.~\\ref{dav:fig:n2992}. In only 3 galaxies was no broad Br$\\gamma$ detected: Circinus, NGC\\,1068, and NGC\\,1097 (in which even the narrow Br$\\gamma$ is so weak that it is almost lost in the stellar absorption features). ",
        "conclusions": ""
    },
    "0601/astro-ph0601621_arXiv.txt": {
        "abstract": "We compile a number of recent observations to estimate the time-averaged rate of formation or ``buildup'' of red sequence galaxies, as a function of mass and redshift. Comparing this with the mass functions of mergers and quasar hosts, and independently comparing their clustering properties as a function of redshift, we find that these populations trace the same mass distribution, with the same characteristic masses and evolution, in the redshift interval $0<z\\lesssim1.5$. Knowing one of the quasar, merger, or elliptical mass/luminosity functions, it is possible to predict the others.  Allowing for greater model dependence, we compare the rate of early-type ``buildup'' with the implied merger and quasar ``triggering'' rates as a function of mass and redshift, and find agreement. We show that over this redshift range, observed gas-rich merger fractions can account for the entire bright quasar luminosity function, and buildup of the red sequence at all but the highest masses at low redshift ($\\gtrsim10^{11}\\,M_{\\sun}$ at $z\\lesssim0.3$) where dissipationless ``dry'' mergers appear to dominate. This supports a necessary prediction of theories which postulate that mergers between gas-rich blue galaxies produce ellipticals with an associated phase of bright quasar activity, after which the remnant becomes red. All of these populations, regardless of sample selection, trace a similar characteristic ``transition'' mass reflecting the characteristic mass above which the elliptical population is mostly ($\\gtrsim50\\%$) assembled at a given redshift, which increases with redshift over the observed range in a manner consistent with previous suggestions that ``cosmic downsizing'' may apply to red galaxy {\\em assembly} as well as star formation. We show that these mass distributions as a function of redshift do not uniformly trace the all/red/blue galaxy population, ruling out models in which quasar activity is generically associated with either star formation or is long-lived in ``old'' systems. ",
        "introduction": "\\label{sec:intro}  Observations motivate the notion of ``cosmic downsizing'' \\citep[as coined by][]{Cowie96}, with the global star formation rate declining rapidly below $z\\sim2$, and the sites of galactic star formation shifting to smaller masses at lower redshift. Moreover, galaxy surveys such as SDSS, COMBO-17, and DEEP2 demonstrate that the color distribution of galaxies is bimodal \\citep[e.g.,][]{Strateva01,Balogh04}, and that this bimodality extends at least to $z\\sim1$ \\citep[e.g.,][]{Bell04,Faber05}.  It is increasingly established that high mass, red elliptical galaxies have older stellar populations than smaller spheroids \\citep[e.g.,][]{Caldwell03,Nelan05,Gallazzi06}. But, many studies also see a significant population of massive/luminous galaxies in place (i.e.\\ assembled) by $z\\sim2$ \\citep[e.g.,][and references therein]{Papovich06,Renzini06}, with measurements of galaxy stellar mass functions (MFs) and luminosity functions at redshifts $0<z<2$ favoring either a uniform increase or buildup in the numbers of early-type (``red sequence,'' RS) galaxies \\citep[e.g.,][]{Bell04,Faber05} or an anti-hierarchical scenario in which this ``buildup'' at $z\\lesssim1$ occurs primarily at the low-mass end of the RS \\citep{Bundy05.full,Bundy05.masslimit,Zucca05,Yamada05, Borch06,Franceschini06,Pannella06,Cimatti06,Brown06}. The blue, disk dominated, star forming galaxy mass function (dominant at low mass), meanwhile, remains relatively constant, or perhaps declines to $z=0$. As a consequence, the ``transition mass,'' above which the red galaxy population dominates the galaxy MF, decreases with time, tracing this downsizing trend. There is evidence for some evolution at the highest masses as ellipticals grow by spheroid-spheroid or ``dry'' mergers \\citep{vanDokkum05,Bell06}, but this, by definition, proceeds strictly hierarchically, and {\\em cannot} account for the movement of mass onto the RS in the first place or any buildup in the number density of low-mass ellipticals.  Meanwhile, the discovery of tight correlations between the masses of central supermassive black holes (BHs) in galaxies and the bulge or spheroid stellar mass \\citep{Magorrian98}, velocity dispersion \\citep{Gebhardt00,FM00} or concentration \\citep{Graham01} implies that the formation of galaxies and BHs must be linked.  Moreover, the evolution of the quasar luminosity function (QLF) shows a sharp decline after $z\\sim2$, with the density of lower-luminosity AGN peaking at low redshift \\citep[e.g.,][and references therein]{HMS05}. To the extent that BH assembly traces galaxy assembly (i.e.\\ there is weak evolution in the BH-host mass relation, as observed to at least $z\\gtrsim1$ by e.g.\\ \\citet{Shields03,AS05a,Peng06}), this implies early {\\em assembly} times ($z\\gtrsim1$) for many of the most massive systems containing $M_{\\rm BH}\\gtrsim10^{8}\\,M_{\\sun}$ BHs.  A number of theoretical models have been proposed to explain the evolution of these populations with redshift, and their correlations with one another \\citep[e.g.,][]{KH00,Somerville01,WyitheLoeb03,Granato04,Scannapieco05, Baugh05,Monaco05,Croton05,H06b,H06c,H06d,Cattaneo06}. In many of these models, the merger hypothesis \\citep{Toomre77} provides a potential physical mechanism linking galaxy star formation, morphology, and black hole evolution and explaining these various manifestations of cosmic ``downsizing.'' In this scenario, gas-rich galaxy mergers channel large amounts of gas to galaxy centers \\citep[e.g.,][]{BH91,BH96}, fueling powerful starbursts \\citep[e.g.,][]{MH94,MH96} and buried BH growth \\citep[e.g.,][]{Sanders88,BH92} until the BH grows large enough that feedback from accretion rapidly unbinds and heats the surrounding gas \\citep{SR98}, leaving an elliptical galaxy satisfying observed correlations between BH and spheroid mass. Major mergers rapidly and efficiently exhaust the cold gas reservoirs of the progenitor systems, allowing the remnant to rapidly redden with a low specific star formation rate, with the process potentially accelerated by the expulsion of remnant gas by the quasar \\citep[e.g.,][]{SDH05a}. This naturally explains the observed close association between the elliptical and red galaxy populations \\citep[e.g.,][]{Kauffmann03}.  In a qualitative sense, the evolution of the characteristic mass at which these processes occur can be understood as follows. Mergers proceed efficiently at high redshift, occurring most rapidly in the regions of highest overdensity corresponding to the most massive galaxies, building up the high-mass elliptical MF.  However, once formed these galaxies are ``dead'', and mergers involving gas-rich galaxies must transition to lower masses.  Recent hydrodynamical simulations, incorporating star formation, supernova feedback, and BH growth and feedback \\citep{SDH05b} make it possible to study these processes self-consistently and have lent support to this general picture.  Mergers with BH feedback yield remnants resembling observed ellipticals in their correlations with BH properties \\citep{DSH05}, scaling relations \\citep{Robertson05b}, colors \\citep{SDH05a}, and morphological and kinematic properties \\citep{Cox06a,Cox06b}.  The quasar activity excited through such mergers can account for the QLF and a wide range of quasar properties at a number of frequencies \\citep{H05a,H06b}, and with such a detailed model to ``map'' between merger, quasar, and remnant galaxy populations it is possible to show that the buildup and statistics of the quasar and red galaxy populations are consistent and can be used to predict one another \\citep{H06c}.  However, it is by no means clear whether this is, in fact, the dominant mechanism in the buildup of early-type populations and quasars and their evolution with redshift. For example, many semi-analytic models incorporate quasar triggering/feedback and morphological transformation by mergers \\citep{KH00,Volonteri03,Volonteri06, WyitheLoeb03,Somerville04a,Monaco05,Bower06,Lapi06,Menci06}. However, some models tie quasar activity directly to star formation \\citep[e.g.,][]{Granato04}, implying it will evolve in a manner tracing star-forming galaxies, with this evolution and the corresponding downsizing effect roughly independent of mergers and morphological galaxy segregation at redshifts $z\\lesssim2$. Others invoke post-starburst AGN feedback to suppress star formation on long timescales and at relatively low accretion rates through e.g.\\ ``radio-mode'' feedback \\citep{Croton05}, which, if this is also associated with optical QSO modes, would imply quasars should trace the established ``old'' red galaxy population at each redshift. There are, of course, other sources of feedback, with galactic superwinds from star formation presenting an alternative means to suppress subsequent star formation, although the required wind energetics are sufficiently high to prefer a quasar-driven origin \\citep[e.g.,][]{Benson03}. Several models invoke a distinction between ``hot'' and ``cold'' accretion modes \\citep{Birnboim03,Keres05,Dekel06}, in which new gas cannot cool into a galactic disk above a critical dark matter halo mass, potentially supplemented by AGN feedback \\citep{Binney04}, as the dominant distinction between the blue cloud and red sequence, essentially independent of effects on scales within galaxies.  It is also important to distinguish the processes which may be associated with the initial movement of galaxies onto the red sequence from their subsequent evolution. Once morphologically transformed by a gas-rich merger, for example, mass can be moved ``up'' the RS (galaxies increased in mass) by gas-poor mergers, but it cannot be {\\em added} to the red sequence in this manner. It also remains an important cosmological question to understand how, once formed, further growth of ellipticals by accretion or ``cooling flows'' may be halted. The models above invoke various feedback processes, including ``radio mode'' activity \\citep[e.g.,][]{Croton05}, cyclic quasar or starburst-driven feedback \\citep{Somerville01,Granato04,Binney04,Monaco05}, massive entropy injection from a single quasar epoch \\citep[e.g.,][]{WyitheLoeb03,Scannapieco04}, and ``hot mode'' accretion \\citep{Birnboim03} to address this problem. Although critical to our understanding of galaxy formation, these processes must operate over timescales of order the Hubble time for all massive galaxies once formed, and therefore are {\\em not} necessarily associated with the {\\em addition} of mass to the red sequence. As such, the details of these long-term suppression mechanisms should be studied in different (e.g.\\ already formed elliptical) populations, and are outside the scope of this paper.  Observationally, it is still unclear whether mergers can account for the the buildup of elliptical and/or quasar populations \\citep[see, e.g.,][and references therein]{Floyd04,RJ06,Lotz06b}. Even within the context of the merger hypothesis, the relative importance of dissipational (gas rich, disk) vs.\\ dissipationless (gas poor, spheroid-spheroid) mergers is unclear \\citep[e.g.,][]{vanDokkum05,Bell06}, although all measurements agree that the ``dry'' merger rate is much less than the gas-rich merger rate at all observed redshifts \\citep{Bell06,Lotz06b,Bell06b}. This is essentially related to the critical question of whether the buildup of the red sequence and elliptical populations is dominated by the formation or movement of ``new'' early-type galaxies onto that sequence or instead by the hierarchical assembly of small ``seed'' early-types and substructure formed at high redshift (which will also not trigger quasar activity).  Fundamentally, it is not clear and has not yet been tested whether the observed downsizing trends in the transition mass, galaxy stellar populations, quasars, and other populations are in fact {\\em quantitatively} the same trend, or merely {\\em qualitatively} similar. This represents a key test which can distinguish between several of the various scenarios above. Attempting to predict the values of this transition mass in an a priori cosmological manner is inherently model dependent and, at least at low redshift, degenerate between the various models described above. However, if mergers are indeed the critical link in the process causing the flow of galaxies from the blue to red sequence and triggering quasar activity, then it is a strong prediction of these theories, and specifically the modeling of Hopkins et al.\\ (2005a-d,\\,2006a-d) that the same mergers are responsible for the bulk of the bright quasar population and the buildup of the new mass on the red sequence at each redshift. In other words, these downsizing trends must quantitatively reflect one another.  In this picture, the ``transition mass'' ($\\mtrans$) may represent the ``smoking gun'' of mergers causing the flow of galaxies from the blue to red sequence. Therefore, to the extent that $\\mtrans$ traces the mass at which the red sequence is being ``built'' at some $z$, it should also trace the characteristic mass of star-forming galaxies merging at that time, and the characteristic mass of galaxies hosting quasars which are initially triggered by those mergers. Of particular interest, the empirical test of this association does not require the adoption of some a priori model for galaxy formation.  Here, we consider the observed $\\mtrans$ over the interval $0<z<2$, and compare it to the characteristic masses of quasar hosts and merging galaxies over the same range in redshift. We demonstrate that they appear to be evolving in a manner consistent with a merger-driven unification model of quasars, interacting galaxies, and the red galaxy population. Note that we use the term ``quasar'' somewhat loosely, as a proxy for high-Eddington ratio accretion inevitably caused by gas-rich mergers, although there may be other triggering mechanisms as well \\citep{Sanders88,Alexander05a,Alexander05b,Borys05,H06b}. Such activity will of course be significantly weaker in small systems (especially those typical of local ULIRGs) and may not technically qualify as a ``classical'' optical quasar \\citep{H05b}, but this distinction is essentially arbitrary and has little impact on our analysis \\citep[see also][]{H06d}.  We adopt a $\\Omega_{\\rm M}=0.3$, $\\Omega_{\\Lambda}=0.7$, $H_{0}=70\\,{\\rm km\\,s^{-1}\\,Mpc^{-1}}$ cosmology. All stellar masses are rescaled to a \\citet{Salpeter55} IMF. ",
        "conclusions": "\\label{sec:discussion}  We compile a large number of observations of red/elliptical galaxy mass functions, and use these to determine the rate of ``buildup'' of the red sequence (RS) as a function of mass and redshift. Comparing these with observations of other populations allows us to test a number of different models for the possible associations between these populations and the ``transition'' of galaxies from blue, star-forming disks to red, ``dead'' ellipticals.  Independent of the nature of ``downsizing'' in the buildup of RS mass functions (discussed below), the rate of RS ``buildup'' is sufficiently well-determined to place meaningful constraints on a number of models. Dissipationless (gas poor, red, or ``dry'') mergers can account for the buildup of the RS at only the largest masses $\\gtrsim10^{11}\\,M_{\\sun}$ at low redshift ($z\\lesssim0.3$). At higher redshifts ($z\\gtrsim0.5$), the dry merger rate would have to be at least an order of magnitude larger than observationally estimated \\citep{vanDokkum05,Masjedi06,Bell06,Bell06b,Lotz06b} to account for observed RS buildup, even at the highest masses. This is perhaps unsurprising, as these and other observations find the {\\em gas-rich} merger rate/fraction is an order of magnitude or more larger at all but the lowest redshifts. Furthermore, the total mass density on the RS is observed to increase by a factor $\\sim2.5-3$ since $z\\sim1$ \\citep[e.g.,][]{Bundy05.full,Franceschini06,Pannella06,Borch06}, and dry mergers cannot, by definition, move/form ``new'' galaxies and mass on the RS.  However, we find the {\\em total} observed merger population (gas rich+gas poor) agrees very well with that expected if all RS galaxies are formed in mergers. Both the detailed mass distribution and fraction/rates of galaxy mergers are consistent with the rate of RS buildup at all masses and redshifts observed. This merger population is dominated by gas-rich mergers at all masses at high redshifts \\citep[$z\\gtrsim0.5$][]{Bell06b,Lotz06b} and at low masses at low redshifts, morphologically identifiable as bright (i.e.\\ star-forming or starbursting) interacting systems \\citep[e.g.,][]{Bundy05.masslimit,Wolf05}. In detail, completely neglecting dry mergers (or merger mass functions sensitive to them), this agreement is unchanged except for the highest masses at low-$z$ discussed above. There is substantial systematic uncertainty in converting a merger fraction to a merger rate; our comparisons assume a characteristic observable merger timescale $\\sim0.5\\,$Gyr. However, this is a theoretically reasonable timescale (see \\S~\\ref{sec:compare.merger.mass}), and given the scatter in the observations, our conclusions are not changed for systematic shifts within a factor $\\lesssim2$, nor for allowing the merger timescale to scale with halo dynamical times ($\\propto(1+z)^{-3/2}$). Furthermore, this has no effect our comparison of the mass distributions of these populations.  Similarly, we find the rate at which host galaxies trigger quasars, determined as a function of the host stellar mass and redshift from the quasar luminosity function, agrees well with the observed RS buildup at all masses and redshifts observed. There is some discrepancy at the lowest redshifts and highest masses, but this is again where the dry merger contribution can account for the observed buildup, and dry mergers (by definition being gas-poor or gas-free) are not expected to trigger quasar activity. We consider this comparison first in a purely empirical fashion, using observed quasar Eddington ratios and the black hole-host mass relation to estimate quasar host masses as a function of redshift, and then in greater detail adopting the models of quasar light curves and lifetimes as a function of luminosity and host properties from the simulations of merger-induced quasar activity in \\citet{H06b}. The latter introduces some model dependence (although it is consistent with the Eddington ratio and black hole-host mass relation estimates), but allows us to consider this comparison in greater detail and to make specific predictions for the characteristic host masses of quasars as a function of their position on the QLF and for the AGN or ``active'' fraction of galaxies as a function of stellar mass. In either case, the agreement between the rates of quasar formation/triggering as a function of host stellar mass and the buildup of RS galaxies is similar.  We independently test these possible associations by comparing clustering measurements of the relevant populations as a function of redshift, and find similar results. The clustering of quasars and systems ``in transition'' to the RS agree at all redshifts as if they trace the same mass distribution. Clustering properties of merger (ULIRG and SMG) and post-merger (E+A) populations are consistent, but considerably less well-constrained.  Although the above comparisons do not technically depend on it, we determine the ``transition'' mass ($\\mtrans$, $M_{Q}$), i.e.\\ the mass which separates the blue, star-forming disk and red, non star-forming elliptical populations, as a function of redshift. It has been suggested \\citep[e.g.,][]{Bundy05.masslimit} that this represents the characteristic mass at which galaxies are forming on or being added to the RS as a function of redshift, but quantified in this manner, it is not obviously so \\citep[see, e.g.][]{Shankar06}. We therefore also determine \\citep[$M_{50}$, following][]{Cimatti06} the minimum mass above which the RS mass function is $\\gtrsim50\\%$ assembled at a given redshift. Regardless of definition, and furthermore regardless of the criterion used to separate early and late-type populations (whether e.g.\\ a color, star formation rate, or morphology criterion), $\\mtrans$/$M_{Q}$/$M_{50}$ shift to systematically larger masses at higher redshift (significant at $>6\\sigma$), tracing a very similar trend as a function of redshift.  This trend, especially in $M_{50}$ (which is independent of possible evolution in late-type mass functions) suggests that ``downsizing'' applies not just to galaxy {\\em star formation}, but also in some sense to galaxy {\\em assembly}, as suggested by the studies of e.g.\\ \\citet{Bundy05.full,Bundy05.masslimit,Zucca05,Yamada05, Franceschini06,Cimatti06,Fontana06,Brown06}. In greater detail, considering the full rate of RS buildup as a function of stellar mass and redshift, low mass ($\\lesssim10^{11}\\,M_{\\sun}$) galaxies appear to be building up rapidly/continuously at low redshifts ($\\sim7-15\\%$ per Gyr), but the most massive systems do not ($\\sim1\\%$ per Gyr growth at $z\\sim0$). The growth of the most massive systems instead appears to be rapid at significantly higher redshifts (e.g.\\ rising to $\\sim20-50\\%$ per Gyr by $z\\sim1$). Equivalently, the characteristic mass (Schechter function $M_{\\ast}$) defined by this ``formation'' rate appears (albeit at only $\\sim2-3\\,\\sigma$) to increase with redshift in a similar fashion to the ``transition'' mass.  We compare the ``transition'' mass with the characteristic masses of mergers and quasars and again find they trace similar masses as a function of redshift, with ``downsizing'' evident in all three populations ($>3\\sigma$ for mergers, $>6\\sigma$ for quasars), further supported by their observed clustering. We compare with the characteristic (Schechter function) $M_{\\ast}$ of the entire, red, and blue galaxy populations, and rule out at high significance the possibility that ``transition'' mass objects are drawn uniformly from any of these populations. Even with the systematic uncertainties in this mass estimate, it is also clearly distinct as a function of redshift from the characteristic masses of e.g.\\ cluster, radio galaxy, ERO, DRG, or LBG populations \\citep[see e.g.\\ Figure~2 of][]{Farrah06}.  These observations are all consistent with and suggest a scenario in which major mergers, quasars, and the transition from blue disk to red elliptical galaxies are associated. They do not inform us regarding, for example, whether gas exhaustion or stellar or quasar feedback is the specific mechanism for the reddening which accompanies the merger-driven morphological transformation and quasar episode. However, they support the hypothesis that mergers drive the transition from blue disks to red elliptical galaxies, terminating in decaying, feedback-driven bright quasar phases. The transition mass and break in the QLF appear to reflect the characteristic mass of gas-rich objects merging at a given redshift, which may build up the {\\em new mass} on the red sequence at progressively lower masses at lower $z$ as gas supplies are exhausted in more massive systems.  That quasar host masses trace the ``transition'' mass and not, e.g.\\ the blue galaxy population $M_{\\ast}$ (see also Figure~\\ref{fig:compare.models}) rules out the possibility that quasar activity {\\em generically} traces star formation, as variants of e.g.\\ the \\citet{Granato04} models might predict. Likewise, that quasar masses do not trace the red galaxy $M_{\\ast}$, as they would if quasar activity was long-lived or randomly (but uniformly as a function of mass) episodic in high-mass black holes. The former case would be expected from the low-level AGN activity invoked in e.g.\\ \\citet{Croton05}, if ``radio'' and optical or high-Eddington ratio quasar activity were associated, but they are in fact generally believed to be distinct \\citep[e.g.,][]{Ho02,White06,Koerding06}. The latter case implies strong limits on implementations of e.g.\\ the \\citet{Binney04} model, which seek to suppress cooling flows through sporadic but potentially high accretion rate AGN activity.  Although we can rule out some alternatives to the merger scenario, there remain a number of viable variants of ``quenching'' models, in which crossing the critical halo mass $\\mhaloquench$ and entering a ``hot'' accretion regime plays a key role in the ``transition'' to the RS. Especially given that some feedback mechanism is typically required, even in the ``hot'' accretion regime, to prevent the formation of cooling flows, it is easy to imagine a scenario in which, upon entering this regime, new infalling halo gas is shock-heated, but feedback from e.g.\\ a central disk galaxy with a small black hole is inefficient, cold gas reservoirs remain large, and cooling flows can form. Thus the system will not redden until it subsequently undergoes a major merger, which morphologically transforms the system, rapidly exhausts the remaining cold gas reservoir, and triggers a quasar and builds up a massive black hole, injecting some level of feedback and enabling efficient future (e.g.\\ cyclic AGN or radio-mode) feedback. The ``hot'' accretion regime may be a necessary prerequisite for feedback to efficiently prevent subsequent cooling, and as discussed in \\S~\\ref{sec:halomass}, our comparisons are all consistent with this possibility. This is generally similar to the scenario assumed in e.g.\\ \\citet{Croton05}, although they do explicitly incorporate quasar light curves or feedback. Recognizing these distinctions (as opposed to e.g.\\ assuming a system simply ``shuts down'' upon reaching $\\mhaloquench$) will probably have little effect on the $z=0$ predictions of semi-analytic models, since in either case star formation will be effectively suppressed at relatively early times in the most massive systems \\citep[see also][]{Cattaneo06}. However, at higher redshifts when massive objects are still forming, the distinctions will almost certainly be significant.  We note that none of our conclusions conflict with the hypothesis that, once formed, elliptical galaxies can continue to grow by dry mergers. However, they emphasize that the importance of such mergers is restricted to the most massive galaxies at low redshifts. Our results, even the steep evolution of $M_{50}$ implying some ``downsizing'' in red galaxy assembly, are all consistent with (and in fact, marginally {\\em favor}) the relatively low observationally inferred dry merger rate ($\\sim1$ major dry merger since $z\\sim1$). Essentially, ``downsizing'' in galaxy assembly as we have quantified it is not, strictly speaking, ``anti-hierarchical.'' Massive galaxies still continue to build up their populations to the present; it is simply a statement that the {\\em relative} rates of red galaxy formation/assembly decrease or ``slow down'' at late times in the most massive systems. This could be related to pure dark matter processes, for example the rapid evolution in large overdensities could simply exhaust the ``supply'' of galaxies with which to merge, or the cluster environments of massive systems at low redshift attain sufficient circular velocities as to rapidly reduce merger rates \\citep[see also][]{Neistein06}. The evolution of the ``transition'' mass may, alternatively, be a statement that galaxy assembly does not strictly trace halo assembly. There are a number of baryonic processes which make this possible, as it simply requires that the effective baryon conversion efficiencies in galaxies be a function of time, or different for central vs.\\ satellite systems.  Improved measurements of early-type mass functions at high redshift ($z\\sim1$), larger samples of mergers from which to construct merger mass functions, revised or direct determinations of high-redshift conditional mass functions, and direct observations of the masses of quasar hosts will substantially improve the constraints in this paper. Ultimately, the integration of the merger and quasar host mass functions may enable a purely observational comparison with the remnant, red galaxy mass function. Calibration of the observable merger ``timescale'' with realistic high resolution galaxy merger simulations, i.e.\\ calibration of selection efficiencies for observed merger fractions, can further remove the factor $\\sim2$ uncertainty in comparing the rates of elliptical buildup and observed merger populations. The association favored here between mergers, quasars, and elliptical buildup also makes specific predictions for the characteristic masses of E+A galaxies and quasar hosts as a function of redshift, which should be testable in future wide-field surveys.  The scenario we have described does not, of course, imply that mergers, quasars, and remnant ellipticals will necessarily be recognizable as the same, singular objects at a given instant -- in fact, simulations which follow the transition through these stages \\citep[e.g.,][]{H06c}, predict that they will be seen as distinct phases in merger-triggered evolution, and observations tracking e.g.\\ the associations between dynamical merger state and quasar activity \\citep[e.g.,][]{Straughn05} support this distinction. What we ultimately find evidence for here in the masses, luminosities, and clustering properties of mergers, galaxies being ``added'' or ``in transition'' to the red sequence, and quasars is that they are drawn from the same ``parent'' population, and that this population is distinct from the ``quiescent'' all/red/blue galaxy population. Again, none of this strictly implies causality, but it does favor models which associate these populations with the same event, a natural expectation if mergers of gas-rich galaxies trigger quasars and morphologically transform disks to spheroids, moving ``new'' mass to the early-type population and leaving an elliptical, gas-poor, rapidly reddening remnant galaxy."
    },
    "0601/astro-ph0601147_arXiv.txt": {
        "abstract": "{\\it Spitzer Space Telescope} mid-infrared images of the radio galaxy Centaurus A reveal a shell-like, bipolar, structure 500~pc to the north and south of the nucleus. This shell is seen in 5.8, 8.0 and 24$\\mu$m broad-band images. Such a remarkable shell has not been previously detected in a radio galaxy and is the first extragalactic nuclear shell detected at mid-infrared wavelengths. We estimate that the shell is a few million years old and has a mass of order million solar masses.  A conservative estimate for the mechanical energy in the wind driven bubble is $10^{53}$ erg. The shell could have created by a small  few thousand solar mass nuclear burst of star formation. Alternatively, the bolometric luminosity of the active nucleus is sufficiently large that it could power the shell. Constraints on the shell's velocity are lacking. However, if the shell is moving at 1000~km~s$^{-1}$ then the required mechanical energy would be 100 times larger. ",
        "introduction": "The nearest of all the giant radio galaxies, Centaurus~A (NGC~5128) provides a unique opportunity to observe the dynamics and morphology of an active galaxy in detail across the electromagnetic spectrum. For a recent review on this remarkable object see \\citet{israel}. In its central regions, NGC~5128 exhibits a well recognized, optically-dark band of absorption across its nucleus. {\\it Spitzer} images of the galaxy reveal a parallelogram shape \\citep{quillen_irwarp} that has been modeled as a series of folds in a warped thin disk (e.g., \\citealt{bland86,bland87,nicholson92,sparke96,quillen_irwarp}).  In this letter we report on the discovery of a 500~pc sized bipolar shell in mid-infrared images in the center of the galaxy. Shells have been previously seen in active and star forming galaxies (for a recent review on galactic winds see \\citealt{veilleux05}). For example, the Seyfert 2 galaxy NGC~2992 exhibits a figure-eight shaped morphology in the radio continuum \\citep{ulvestad84}. The starburst/LINER galaxies NGC~2782 \\citep{jogee98}, and NGC~3079 \\citep{ford86,veilleux94}, exhibit partially closed shell morphologies in optical emission lines and radio continuum.  The above shells have been detected in radio continuum, optical emission lines and in emission from atomic hydrogen. The only nuclear shell to have been previously detected at mid-infrared wavelengths is the 170~pc large bipolar bubble in the direction of the Galactic center \\citep{bland03}. As pointed out by \\citet{bland03}, infrared observations allow unique constraints on the mass of swept up material. Theoretical models predict that in the early stages of a starburst driven outflow supernovae and stellar winds inject energy into the ISM forming a bubble of hot gas and thermalized ejecta (e.g., \\citealt{castor75,chevalier85,tomisaka88,heckman90}). Active galaxies can also drive winds (e.g., see \\citealt{veilleux05,begelman04}).  Based on the discussion by \\citet{israel}, we adopt a distance to Cen A of $3.4$~Mpc. At this distance $1'$ on the sky corresponds to $\\sim 1$~kpc. ",
        "conclusions": "In this paper we have presented {\\it Spitzer Space Telescope} observations of the nuclear region of the galaxy Centaurus A that reveal a previously undetected 500~pc radius shell above and below the gaseous and dusty warped disk.  The shell resembles a bipolar outflow or bubble, has an axis ratio of $\\sim 0.6$, a position angle of $\\sim 10^\\circ$ and a surface brightness of  $\\sim 10$~MJy~sr$^{-1}$ at 8.0$\\mu$m. It is extended in the direction perpendicular to the gaseous and dusty disk and not in a direction obviously related to the radio jet. We conservatively estimate that the shell contains a million solar mass of hydrogen and that the energy required to create it is $\\sim 10^{53}$ergs. Unfortunately we lack constraints on the shell's velocity. If the shell is expanding at 1000~km~s$^{-1}$ then the energy required could be 100 times larger.   While a blue shifted diffuse component was detected at the nucleus by \\citet{marconi06}, most spectroscopic observations fail to cover the shell or lack the bandwidth or sensitivity to have detected it. Observational spectroscopic studies are needed to confirm the presence of this transient shell and place better constraints on its mass, composition and energetics.  If the energy in the shell is low ($10^{53}$ ergs), then a modest starburst of a few thousand solar masses could have provided the mechanical energy. The orientation of the shell differs from that of the radio jets so the shell may not have been caused by the active nucleus.   However the bolometric luminosity of the active nucleus exceeds that required to expand the bubble so the active nucleus could have been important in its creation. This shell is too small to have evacuated the 0.1 - 0.8 kpc gap in the dust distribution in the disk reported by \\citet{quillen_irwarp}, suggesting that there could be or have been more than one expanding bubble in the heart of this galaxy."
    },
    "0601/astro-ph0601371_arXiv.txt": {
        "abstract": "{ The Galactic synchrotron emission is expected to be the most relevant source of astrophysical contamination in cosmic microwave background polarization measurements, at least at frequencies $\\nu \\lsim 70$~GHz and at angular scales $\\theta\\gsim 30'$. We present a multifrequency analysis of the Leiden surveys, linear polarization surveys covering the Northern Celestial Hemisphere at five frequencies between 408~MHz and 1411~MHz. By implementing specific interpolation methods to deal with these irregularly sampled data, we produced maps of the polarized diffuse Galactic radio emission with a pixel size $\\simeq 0.92^\\circ$. We derived the angular power spectrum (APS) ($PI$, $E$, and $B$ modes) of the synchrotron dominated radio emission as function of the multipole, $\\ell$. We considered the whole covered region and % some patches at different Galactic latitudes. By fitting the APS in terms of power laws ($C_\\ell \\sim \\kappa\\cdot\\ell^{\\alpha}$), we found spectral indices that steepen with increasing frequency: from $\\alpha \\sim -$(1-1.5) at 408~MHz to $\\alpha \\sim -$(2-3) at 1411~MHz for $10 \\lsim \\ell \\lsim 100$ and from $\\alpha \\sim -0.7$ to $\\alpha \\sim -1.5$ for lower multipoles (the exact values depending on the considered sky region and polarization mode). The bulk of this flattening at lower frequencies can be interpreted in terms of Faraday depolarization effects. We then considered the APS at various fixed multipoles and its frequency dependence. Using the APSs of the Leiden surveys at 820~MHz and 1411~MHz, we determined possible ranges for the rotation measure, $RM$, in the simple case of an interstellar medium {\\it slab model}. Also taking into account the polarization degree at 1.4~GHz, it is possible to break the degeneracy between the identified $RM$ intervals. The most reasonable of them turned out to be $RM \\sim 9-17$~rad/m$^2$ although, given the uncertainty on the measured polarization degree, $RM$ values in the interval $\\sim 53-59$~rad/m$^2$ cannot be excluded. ",
        "introduction": "In the last decade an impressive number of experiments has been dedicated to measurements of the cosmic microwave background (CMB) anisotropies, first observed in temperature by COBE (\\cite{smoot1990}). Modern cosmologies predict also the existence of polarization fluctuations, produced at the recombination epoch via Thomson scattering (see, e.g., \\cite{Kosowsky99}). However, the foreseen degree of polarization of the CMB anisotropies should be $\\lsim 10\\%$, implying a signal much weaker than in temperature and therefore extremely difficult to reveal. The first detection of the CMB polarization achieved by DASI (\\cite{Kovac2002}) and the recent measure by BOOMERanG (\\cite{masietal05};~\\cite{montroyetal05}) confirmed the expectations. The polarization data expected by the NASA WMAP~\\footnote{http://lambda.gsfc.nasa.gov/product/map/} satellite will permit us to extend the information about polarization to a wider portion of the sky, though at small scales ($\\theta \\lsim 0.5^{\\circ}$) its sensitivity will likely allow a detection rather than an accurate measure.  In the near future, the ESA {\\sc Planck}~\\footnote{http://www.rssd.esa.int/planck}  satellite will significantly improve these results by observing the whole sky at nine different frequencies between 30 and 857~GHz, both in temperature and polarization, with unprecedented resolution and sensitivity (e.g. see \\cite{tauber04} and references therein). In particular, the sensitivities per pixel of the {\\sc Planck} maps are expected to be between 3-4 and 10 times better than in the WMAP data, producing a significant improvement in terms of angular power spectrum recovery. Consequently, {\\sc Planck} alone will allow an estimate of the cosmological parameters much more precise than that obtained by WMAP and other CMB experiments combined with different cosmological data. Moreover, it will accurately measure the CMB $E$-mode and could potentially observe the $B$-mode as well (e.g. see \\cite{burigana04} and references therein).  Strong efforts have been made to evaluate the impact of the foregrounds on the feasibility of CMB anisotropy measurements, as well as to understand how to separate the different components contributing to the observed maps. The methods elaborated for this purpose can be divided in two categories: blind (Independent Component Analysis, see \\cite{ica}; \\cite{fastica}; \\cite{icapol}; \\cite{delabrouille}) and non-blind (Wiener filtering and Maximum Entropy Methods, see respectively \\cite{wf} and \\cite{mem}). The latter class needs an a priori knowledge of the frequency and spatial dependence of the foreground properties, while a blind approach only requires maps at different frequencies as input. On the other hand, blind methods benefit from having ancillary data (such as realistic templates) to improve the separation robustness and quality~\\footnote{Future component separation methods will probably search for a good compromise between the relaxation of the a-priori assumptions on the independent signals superimposed on the microwave sky and the exploitation of the physical correlation among the various components. Also it will be crucial to analyse microwave data together with ancillary data from radio and infrared frequencies.}.   The Galactic polarized diffuse synchrotron radiation is expected to play a major role at frequencies below 70~GHz on intermediate and large angular scales ($\\theta \\gsim 30'$), at least at medium and high Galactic latitudes where satellites have the clearest view of the CMB anisotropies. At about 1~GHz the synchrotron emission is the most important radiative mechanism out of the Galactic plane, while at low latitudes it is comparable with the bremsstrahlung; however the free-free emission is unpolarized, whereas the synchrotron radiation could reach a theoretical intrinsic degree of polarization of about $75\\%$. Consequently, radio frequencies are the natural range for studying the Galactic synchrotron emission, although it might be affected by Faraday rotation and depolarization. Nowadays\\footnote{During the final phase of the revision process of this work, the DRAO 1.4 GHz polarization survey has been released (\\cite{wolleben05}). The data are available at http://www.mpifr-bonn.mpg.de/div/konti/26msurvey/ or http://www.drao-ofr.hia-iha.nrc-cnrc.gc.ca/26msurvey/ or http://cdsweb.u-strasbg.fr/cgi-bin/qcat/?J/A+A/448/411~. See La Porta et al.~2006 for a first angular power spectrum analysis of the DRAO survey.} the only available data suitable for studying the polarized Galactic synchrotron emission on large scales are the so-called Leiden surveys (\\cite{spo76}). These are linear polarization surveys, extending up to high Galactic latitudes, carried out at five frequencies between 408~MHz and 1411~MHz. We have elaborated specific interpolation methods to project these surveys into maps with a pixel size $\\simeq 0.92^{\\circ}$. These maps can be exploited in different contexts. They are suitable to analyse the statistical properties of the Galactic radio polarized emission. Beside this obvious application, they can be used as input to construct templates for simulation activities in the context of current and future microwave polarization anisotropy experiments. For example they can be utilized to build templates for straylight evaluation (see, e.g., \\cite{challinoretal2000};~\\cite{buriganaetal2001}; ~\\cite{barnesetal2003}) and for component separation analyses.  It is standard practice to characterize the CMB anisotropies in terms of the angular power spectrum~\\footnote{It contains all the relevant statistical information % in the case of pure Gaussian fluctuations.} (APS) as a function of the multipole, $\\ell$ (inversely proportional to the angular scale, $\\ell \\simeq {180}/{\\theta(^{\\circ})}$). The APS is an estimator related to the angular two-point correlation function of the fluctuations of a sky field (\\cite{peebles}). The APS is not able to fully characterize the complexity of the intrinsically non-Gaussian Galactic emission. Nevertheless, it is of increasing use also in the study of the correlation properties of the Galactic foregrounds (see, e.g., \\cite{bacci01} for an application to the polarized radio emission). We stick to the commonly adopted approach and consider the APS of both the polarized intensity, $PI$, and the $E$ and $B$ modes~\\footnote{We keep here the physical dimension of the maps (antenna temperature, typically in K or mK) and consequently express the APS in K$^2$ or mK$^2$.} (see \\cite{kamion} and \\cite{zald}).  In Sect.~2 we briefly summarize the properties of the Leiden surveys and of some selected areas relevant for the analysis at the smaller angular scales. In Sect.~3 we introduce the general criteria adopted to elaborate the class of our interpolation algorithms. The final, optimized version of the algorithm and the accuracy of the derived maps and APSs are described in Sect.~4. In Sect.~5 we present our results in terms of maps, APSs, and a parametric description of the APSs. A brief comparison with a preliminary version of the DRAO polarization survey (\\cite{wolleben03}; \\cite{wollebenPhD}) at 1.4~GHz is given in Sect.~6 in terms of APS. Sect.~7 is devoted to the multifrequency analysis of our results and to their interpretation in the context of an interstellar medium {\\it slab model}. We summarize and discuss our main results in Sect.~8. ",
        "conclusions": "The Leiden surveys have been up to now the only available data at $\\nu \\sim 1$~GHz suitable for a multifrequency study of the polarized diffuse component of the Galactic synchrotron emission on large angular scales. A linear polarization survey of the Southern sky had previously been carried out at 408~MHz using the Parkes telescope (\\cite{mathewson65}); unfortunately these data have never been put in digital form. In Fig.~\\ref{allpol408} an image of the data from this survey has been combined with the Leiden survey at 408~MHz to show an all-sky linear polarization map.  \\begin{figure*} \\centering \\includegraphics[scale=0.8,angle=0]{4321f14.ps} \\caption{Map of polarization vectors at 408 MHz, obtained combining the corresponding figures of the Leiden survey (Brouw \\& Spoelstra 1976) and of the Parkes southern sky survey (\\cite{mathewson65}). Unfortunately, the vector length has a different scale in the two original figures, therefore only a direct comparison of the direction is feasible.} \\label{allpol408} \\end{figure*}  The ``North Polar Spur'', the ``Fan Region'', and the ``Cetus Arc'' (see \\cite{berkhuijsen71} for a sketch of the Galactic Loops these structures are thought to belong to) are visible as areas where the polarization vectors are longer (the length is proportional to the polarized intensity) and more regularly aligned, drawing small circles in the sky. The most plausible theory suggests that these Galactic spurs are the brighter ridges of evolved Supernova Remnant sitting relatively close to the Sun (within a few~$\\times 10^2$~pc) and obscuring our view of the Galactic diffuse non-thermal emission. The APSs obtained for these regions will be characterized by an amplitude considerably higher than the average, representing therefore an upper limit for the polarized signal of the Galactic synchrotron radiation. On the other hand, the APS should be steeper in these regions, since the polarization vectors appear rather ordered and with approximately constant length, implying a relatively low fluctuation power at small scales. These considerations should be kept in mind when extrapolating the results to the microwave range in order to check the impact of the Galactic synchrotron emission on CMB polarization measurements. \\\\ Implementing an interpolation method to deal with the sparse and irregular sampling of the Leiden surveys, we produced maps of $PI$, $Q$ and $U$ with $\\theta_{pixel} \\simeq 0.92^{\\circ}$ at five frequencies between 408 and 1411~MHz. We accomplished several tests in order to check the reliability of the interpolation algorithm and to optimize its quality (see Sects.~3 and 4). We emphasize that the APSs computed from our maps are significantly different from those obtained by other authors (e.g. \\cite{bruscoli02}) for similar areas of the sky, using the same data; namely, at the higher frequencies of the Leiden surveys we find steeper APSs (first recognized in \\cite{buriganalaporta02}). A firm confirmation of our approach comes from the comparison of the APSs obtained for the 1411~MHz map with those of a preliminary version of the DRAO 1420~MHz survey, characterized by much denser sampling and higher sensitivity.  We used the interpolated full coverage maps to study the APS of the polarized synchrotron emission on the multipole range [2,50], being limited at high $\\ell$ by the map sensitivity (related to the accuracy and sampling of the original data). Exploiting some better sampled regions, approximately coinciding with the two most prominent Galactic spurs, we could extend our analysis up to $\\ell \\sim 100$.  At each frequency the APSs of the entire coverage maps turn out to change slope at $\\ell \\sim 10$, $C_{\\ell}^{PI}$,  $C_{\\ell}^{E}$, and $C_{\\ell}^{B}$ presenting a steepening going from the lower to the higher multipoles (see Sect.~5). A general and remarkable result (see Sect.~5) is that the slope of the polarized synchrotron emission APS increases from 408~MHz to 1411~MHz; it changes from $\\alpha \\sim -(1$-1.5) to $\\alpha \\sim -(2$-3) for $\\ell \\sim 10$ to $\\sim 100$ and from $\\alpha \\sim -0.7$ to $\\alpha \\sim -1.5$ for lower multipoles, the exact value depending on the considered sky region and polarization mode. \\\\ One possible explanation for such a behaviour of the APS relies on Faraday depolarization arguments. Fluctuations of the electron density and/or of the magnetic field (in strength and/or direction) might redistribute the synchrotron emission power from the larger to the smaller angular scales, creating fake structures. The effect is more relevant at lower frequencies, where the APS slope is expected to become flatter. We have verified this interpretation by performing a toy-simulation of pure Faraday depolarization in a simple {\\it slab model}. We note that in the better sampled patches the polarization vectors are well ordered and similarly distributed at 820 MHz and 1411 MHz, which is consistent with the slab model assumptions at least for the considered angular scales. We reproduced the effects of Faraday depolarization on simulated polarized intensity maps, representing the intrinsic Galactic synchrotron emission at 408~MHz and 1420~MHz. We constructed an $RM$ map interpolating the data of Spoelstra~(1984) and applied it to transform the input maps according to Faraday depolarization formulae. At 1420~MHz the input and output maps are almost identical, whereas at 408~MHz the large-scale structure of the output map appears depressed and the small scale structures more abundant. Consequently, the corresponding 408~MHz APSs are flatter than the original ones (see Sect.~7.2 for a more precise description of the simulation process). It is therefore preferable to extrapolate the results (spectral index and amplitude of the APS) obtained at 1411~MHz (significantly less affected by Faraday depolarization effect than those at 408~MHz) to estimate the synchrotron APS at microwave frequencies.   By interpreting the ratios of the APS amplitudes at 820 and 1411~MHz (the latter smoothed to the lower angular resolution of the former) in terms of synchrotron emission depolarized by Faraday effect, we have identified possible $RM$ ranges (see Sect.~7.1). This seems a reasonable approach at least in the case of patch 1,~2,~and~3, since the polarization vectors are mostly aligned, indicating rather homogeneous physical conditions. Furthermore, taking into account the upper limit to $RM$ derived from the available information on the degree of polarization at 1.4~GHz and considering that the maximum theoretical degree of synchrotron polarization is $\\sim 75\\%$, we could break or, at least, reduce the degeneracy between the identified $RM$ intervals. For the three better sampled patches, the most reasonable value of $RM$ are $\\sim 9-17$~rad/m$^2$. However, given the uncertainty on the measured polarization degree, $RM$ values in the interval $\\sim 53-59$~rad/m$^2$ cannot be excluded. Higher frequencies observations (at $\\nu \\gsim 2$-3~GHz) would be extremely useful to break the degeneracy left by the available $\\sim 1$~GHz data."
    },
    "0601/astro-ph0601692_arXiv.txt": {
        "abstract": "The Perseus molecular cloud complex is a $\\ga$30\\,pc long chain of molecular clouds most well-known for the two star-forming clusters NGC\\,1333 and IC\\,348 and the well-studied outflow source in B5. However, when studied at mid- to far-infrared wavelengths the region is dominated by a $\\sim$10\\,pc diameter shell of warm dust, likely generated by an H{\\sc ii} region caused by the early B-star HD\\,278942. Using a revised calibration technique the COMPLETE team has produced high-sensitivity temperature and column-density maps of the Perseus region from IRAS Sky Survey Atlas (ISSA) 60 and 100\\,\\um\\ data.  In this paper, we combine the ISSA based dust-emission maps with other observations collected as part of the COMPLETE\\footnote{{\\bf C}o{\\bf O}rdinated {\\bf M}olecular {\\bf P}robe {\\bf L}ine, {\\bf E}xtinction and {\\bf T}hermal {\\bf E}mission; http://cfa-www.harvard.edu/COMPLETE} Survey, along with archival H$\\alpha$ and MSX observations.  Molecular line observations from FCRAO and extinction maps constructed by applying the NICER method to the 2MASS catalog provide independent estimates of the ``true'' column-density of the shell. H$\\alpha$ emission in the region of the shell confirms that it is most likely an H{\\sc ii} region located behind the cloud complex, and 8\\,\\um\\ data from MSX indicates that the shell may be interacting with the cloud.  Finally, the two polarisation components seen towards background stars in the region by \\citet{gbmm90} can be explained by the association of the stronger component with the shell. If confirmed, this would be the first observation of a parsec-scale swept-up magnetic field. ",
        "introduction": "\\label{intro} The interaction of newly formed stars with their parent cloud, particularly in the case of massive stars, can have significant consequences for the subsequent development of the cloud. Hot, massive stars will cause cloud disruption through their ionizing flux as well as through the momentum in their stellar winds. Alternatively, given the right conditions in the cloud, stellar winds may also trigger the collapse of cloud cores and induce further star formation.  One region where a stellar wind may be triggering star formation is the Perseus Molecular Cloud complex, a well-studied chain of molecular clouds, extending approximately 30\\,pc in length.  The complex has a total mass of $\\sim$1.3$\\times$10$^4$\\Msolar, assuming it is at a distance of 260\\,pc (although actual distance estimates range from 230\\,pc \\citep{cernis90} to 350\\,pc \\citep{hj83}, and it is clear from the large ($\\ga$ 10\\,\\kms) velocity range of molecular gas that a single distance to the entire complex is unlikely). Star formation is ongoing in several parts of the complex, most obviously around the two reflection nebulae \\objectname{IC\\,348} and \\objectname{NGC\\,1333}.  Several surveys \\citep[e.g.][]{llm93,asr94,ll95,luhman} have established that there is a population of pre-main sequence stars located both within the clusters and throughout the complex, but relatively few high-mass stars have been found.  The Perseus complex is not just interesting for its star-forming properties. An almost complete shell of enhanced emission can be seen in IRAS data toward the center of the molecular cloud complex (figure \\ref{ext_temp}).  The shell, referred to in previous papers as \\objectname{\\hbox{G\\,159.6$-$18.5},} has a diameter of 0.75\\degrees, or 10\\,pc at the distance of the molecular cloud. Although its existence has been known for about 15 years (it was first described by \\citet{ps89} in a conference proceedings and further discussed by \\citet{fpjd94}) there has been little investigation into the nature or source of the shell, and it has been mostly ignored by studies of star-formation in the region. Based on radio data, \\citeauthor{fpjd94} argued that the feature is caused by a supernova remnant, which if true, would be one of the closest to the Sun, and have the highest known Galactic latitude for such an object. More recently \\citet{dZ99} associated the shell with the B star HD\\,278942, located at its geometric center. \\citet{ander2000} performed a multi-wavelength study of the star and its surroundings. They reclassified HD\\,278942 as an O9.5-B0 V star with an age of 8\\,Myr, and found weak radio continuum emission with a flat spectral index, consistent with an optically thin H{\\sc ii} region filling the shell. However, partly due to the lack of short-spacing information in their interferometer maps leading to large uncertainties on the derived radio-continuum fluxes, they could not rule out a negative spectral index, leaving the possibility that the radio continuum is due to synchrotron emission from the interaction between a stellar wind or supernova remnant and the molecular cloud. \\citeauthor{ander2000}'s other observations concentrated on the central star HD\\,278942, and hence provide little insight into the nature of shell itself.  The Perseus molecular cloud complex is a target of the Spitzer Space Telescope (SST) Legacy program ``From Molecular Cores to Planet Forming Disks'' \\citep[hereafter c2d;][]{evans03}, while the two clusters NGC\\,1333 and IC\\,348 are included in an SST Guaranteed Time Observer (GTO) program. These programs aim to determine the distribution of young stars and clusters, and investigate their association with known dense cores. The ongoing COMPLETE Survey \\citep{data}, is a large international effort which has coordinated its observations with c2d.  COMPLETE aims to obtain and compare high quality molecular line, submillimeter continuum, far-infrared column-density and near-infrared extinction data over the extents of three of the star-forming molecular clouds targeted by c2d, including the whole of Perseus.  As part of COMPLETE, we have recently used 60 and 100\\,\\um\\ IRAS maps of the Perseus region to create new high-sensitivity temperature and column-density maps \\citep{schnee_iras}. Here we intercompare the recalibrated far-IR results with COMPLETE near-infrared extinction and \\xco\\ maps of the same region \\citep{data}, an H$\\alpha$ image from the VTSS\\footnote{Virginia Tech Spectral Line Survey; http://www.phys.vt.edu/~halpha/} Survey, and the MSX 8\\,\\um\\ image. A detailed picture of the shell which, would not be possible by looking at each data set alone emerges and suggests a complex picture of Perseus as a group of cold molecular clouds being impacted gently from behind by a stellar wind bubble. ",
        "conclusions": "\\label{discuss}  \\subsection{Could HD\\,278942 be the driving source of the shell?} The fact that the shell is visible in the 2MASS extinction map and is filled with preferentially heated dust emission tells us that it is a bubble of heated material, enclosed in a colder shell of increased density (visible only on the edges). Based on multiwavelength data, \\citet{ander2000} suggested that the shell was the result of an expanding stellar wind from HD\\,278942. This is supported by the fact that the shell is filled with H$\\alpha$ emission, as is shown in figure \\ref{halpha_iras}.  Historically, the combined effects of the shell and the high reddening towards HD\\,278942 (A$_V$ = 7.4\\,mag; \\citealt{ander2000}) has made spectral classification of the star problematic, with values in the literature ranging from F2 (Hipparcos) through B3\\,III \\citep{cernis93} to O9.5$\\rightarrow$B0\\,V \\citep{ander2000}. The situation is further complicated by the possibility that HD\\,278942 may be a photometric B3\\,III + F5\\,I binary \\citep{cernis93}. Recent spectra obtained by \\citet{steen} confirm that HD\\,278942 is most likely an early--mid B-star. Assuming a local density, n$_H$, of 1\\,cm$^{-3}$, such a star could produce an H{\\sc ii} region of several parsecs in size \\citep{osterbrock}. Given the uncertainty in the star's spectral classification and the local density, it is therefore plausible that HD\\,278942 created the ~10\\,pc diameter shell.   \\subsection{Evidence for Two Dust Populations in Perseus} \\label{twodust}  In order to compare the properties of the shell with other regions of the cloud, we selected by eye an annular region of the IRAS column-density image which contains all of the shell emission.  Figure \\ref{annulus} shows this region, as a hatched annulus overlaid on the IRAS column density image. In the remainder of the text we will refer to the hatched region of the annulus as the ``shell component'', the circular region contained within the annulus as ``inside shell'' and all points in the image that are outside of the annulus as the ``cloud component''. Also shown are two small grey boxes, which indicate the positions of the two young clusters NGC\\,1333 and IC\\,348 which we excluded from our analysis.  \\begin{figure} \\caption{IRAS extinction map of Perseus. The hatched area indicates the region we define as the ``shell component'', points within the circular region contained within the annulus as ``inside shell'' and points outside the outer circle are defined as the ``cloud component.''  The two grey boxes show the small regions around IC\\,348 (left) and NGC\\,1333 (right) which were excluded from the analysis in section \\ref{twodust}. \\label{annulus}} \\end{figure}   Our new IRAS calibration technique \\citep{schnee_iras} leads to temperature and column-density maps which look somewhat different to those presented in \\citet{ander2000}. They interpreted their temperature map as showing that the ring is at a temperature minimum of 24\\,K with the dust interior to the shell at a higher temperature. The fact that they also found the cloud material outside the shell marginally warmer than the shell itself (making the shell a temperature minimum) is probably an artifact of the way they created their temperature map.   We find an average temperature in the shell component of 28.9\\,K, with a maximum of 32\\,K.  The distribution of temperatures within the shell component and cloud component are shown in figure \\ref{temp_hist}. It is clear that the shell component is significantly warmer than the molecular cloud material which surrounds it. This is supported by a two-sided Kolmogorov-Smirnov (KS) test, which gives a probability of $<10^{-5}$ that the two samples are drawn from the same population.  \\begin{figure} \\plotone{f7.eps} \\caption{ Histogram of dust temperature determined from IRAS 60 and 100\\um\\ emission. The thick line represents the cloud component, while the thin line shows the shell component. Note that small areas containing the young clusters IC\\,348 and NGC\\,1333 were excluded (see fig. \\ref{annulus}). The histograms are normalised due to the large difference in the number of points in the two components (65000 points for the cloud versus 6800 for the shell component). \\label{temp_hist}} \\end{figure}    In figure \\ref{ext_temp_scatter} we show a scatter plot of the colour temperature versus column-density derived from recalibrated IRAS 60 and 100\\,\\um\\ emission \\citep{schnee_iras}. Three distinct populations are seen in the emission, indicated by the green points (cloud component), red points (shell component) and blue points (inside shell). Clearly the shell and its interior are warmer and denser on average than the surrounding cloud.  The difference in density can be seen more clearly in the histograms shown in the left panel of figure \\ref{ext_hist}. The cloud component has a mean extinction of 1.6\\.mag., while the mean of extinction in the shell component is almost double that (3.1\\,mag.)  Again, a two-sided KS test gives a probability of $\\ll 10^{-5}$ that the two populations (cloud and shell) are drawn from the same parent population.   \\begin{figure} \\caption{Column density vs. temperature derived from IRAS 60 and 100\\,\\um\\ emission. Green points indicate cloud component points (i.e. outside the outer circle of the annulus shown in figure \\ref{annulus}).  Red dots indicate points within the hatched annulus (shell component) and blue dots indicate points from the interior of the shell (i.e. inside the inner circle of the annulus). \\label{ext_temp_scatter}} \\end{figure} \\begin{figure} \\plottwo{f9a.eps}{f9b.eps} \\caption{Normalised histograms of extinction from IRAS and 2MASS/NICER.  In both panels, the thick line represents the cloud component, and the grey filled histogram shows the shell component as traced by IRAS (hatched region in figure \\ref{annulus}).  The thin line in the right histogram shows the shell component as traced by 2MASS/NICER.  Note that small areas containing the young clusters IC\\,348 and NGC\\,1333 were excluded in both cases (see fig. \\ref{annulus}), and that the histograms are normalised due to the large difference in the number of points in the two components (65000 points for the cloud versus 6800 for the shell component). \\label{ext_hist}} \\end{figure}   Also shown in the right panel of figure \\ref{ext_hist} is a histogram of extinction as derived from 2MASS/NICER for the shell component.  If we assume that the 2MASS/NICER derived extinction is a better estimate of the ``true'' column-density of material towards any line-of-sight then it appears that IRAS is systematically overestimating the column-density of material in the shell.  Another way to look at this is figure \\ref{2mass_co_IRAS_scatter} where we show a scatter plot of column density derived from IRAS against column density from 2MASS/NICER. A solid line in the figure shows a linear least-squares fit to the data and the dashed line indicates a 1:1 relation between the two properties. Clearly, the shell component (red points) fall preferentially above both the fitted line and the 1:1 line, again indicating IRAS seems to be overestimating the column density in the shell region.   \\begin{figure} \\caption{IRAS column density vs. 2MASS extinction. The dashed line shows a 1:1 relation, and the solid line shows a fit to the data. The data clearly differ from the 1:1 line. Note that for clarity only 10\\% of the green (cloud component) points are plotted, but the fit is heavily weighted by this component. Symbols are the same as figure \\ref{ext_temp_scatter}. \\label{2mass_co_IRAS_scatter}} \\end{figure}   The ``shell'' (red) and ``cloud'' (green) populations seen in figure \\ref{ext_temp_scatter} are less apparent when we make a similar plot of temperature determined from IRAS against extinction determined from 2MASS/NICER, as is done in figure \\ref{2mass_iras_temp}. Although still warmer on average, the shell component does not show a significant difference in column-density when compared to the cloud component.  \\begin{figure} \\caption{ Temperature from IRAS 60 and 100\\,\\um\\ emission vs. extinction from 2MASS/NICER. Symbols are the same as figure \\ref{ext_temp_scatter}. \\label{2mass_iras_temp}} \\end{figure}  We interpret the ``high contrast'' appearance of the shell on the IRAS based column-density maps as caused by the presence of two different dust populations along the line of sight, the cool component associated with the molecular cloud complex, and a warmer component due to the shell. This type of situation was discussed in an appendix of \\citet{lwgb89}, who showed that the result of assuming a single dust temperature along a line of sight is an {\\em underestimate} of the total optical depth (and hence column density). However, in this case we know the ``true'' column-density of shell material from the 2MASS/NICER map is in fact less than the column density we determine from IRAS, i.e.  the line-of-sight dust distribution seems to cause an {\\em overestimate} of the 100\\,\\um\\ opacity within the shell component, as can be seen in the right panel of figure \\ref{ext_hist}.  An alternative explanation is that our assumption of a uniform dust-emissivity co-efficient for all of the region is incorrect. It is possible that the grain-size distribution and/or composition of the dust in the shell may be unusual, as it has been processed by the radiation field of the massive star.  The hypothesis of two cloud populations in the direction of Perseus is not new. \\citet{ut87} suggested the possibility of two clouds along the line of sight to explain the presence of two distinct velocity components seen in their CO map of Perseus, and in several papers about photometric distances and extinctions toward field stars in the directions of IC\\,348, NGC\\,1333 and the dark cloud \\objectname{Barnard 1} \\citep[][hereafter CS]{cernis90,cernis93,cernis03}, \\citeauthor{cernis93}  and collaborators have found an apparent jump in extinction towards Perseus at a distance of $\\sim$140-160\\,pc, placing some kind of absorbing material at this distance. They attribute this material to an extension of material from the Taurus molecular cloud complex along the line of sight to Perseus (their suggested morphology is shown in figure 9 of \\citealt{cernis93}).  In a study of the optical polarization of background stars, \\citet[][hereafter GBMM]{gbmm90} found a bimodal distribution of polarisation angles in the Perseus region and attributed it to two separate dust distributions along the line of sight, each associated with a different magnetic field orientation. At the time, comparing only to molecular-line maps, they could detect no meaningful spatial distinction between the two polarisation components.  \\citetalias{cernis03} have 18 stars in common with \\citetalias{gbmm90}, for 15 of which they were able to determine distances and extinctions, based on photometric spectral-typing. These 15 stars' properties are summarised in table \\ref{pol_tab} which is reproduced in most part from \\citetalias{cernis03}. We have split the stars in the table into two groups, those with A$_V <$5 and those with A$_V >$5.  From table \\ref{pol_tab}, it is clear that the stars with {\\em high} polarisation percentage (bold entries in column 6 of the table) appear almost exclusively in the {\\em high} extinction ($\\ga$0.5\\,mag) group, whilst stars with low polarisation percentage are more common at lower extinction. The lower extinction stars also tend to be at nearer distances.  \\citetalias{cernis03} interpreted these results as further evidence for their 140\\,pc material, proposing that stars with low polarisation are located between the 140\\,pc material and the cloud complex, while the more distant stars show high polarisation and high extinction because they are behind two layers of absorbing material (the 140\\,pc material and the molecular cloud).  \\begin{deluxetable}{ccccccc} \\tablecolumns{7} \\tablewidth{0pc} \\tablecaption{Distances, extinctions and polarisations of stars in Perseus, adapted from tables 1 and 3 of \\citet{cernis03}.\\label{pol_tab}} \\tablehead{ \\colhead{(1)} & \\colhead{(2)} &\\colhead{(3)} &\\colhead{(4)} &\\colhead{(5)} &\\colhead{(6)}&\\colhead{(7)}\\\\ \\colhead{ID\\tablenotemark{a}}&\\colhead{RA (J2000)} & \\colhead{Dec (J2000)} &R &\\colhead{A$_V$} & \\colhead{P\\tablenotemark{b}} & \\colhead{$\\Theta$}\\\\ \\colhead{}&\\colhead{hh:mm:ss} & \\colhead{dd:mm.m} &pc &\\colhead{mag} & \\colhead{\\%} & \\colhead{deg}} \\startdata \\sidehead{Stars With A$_V <$0.7} 30&03:33:04 &30:34.5   &    268      &   0.21 &   0.51  &   87   \\\\ 40&03:33:55 &31:10.0   &    310      &   0.29 &   {\\bf1.24}  &   72   \\\\ 51&03:33:41 &31:07.9   &    195      &   0.37  &  0.42  &   80   \\\\ 71&03:35:41 &30:55.7   &    254      &   0.17  &  0.66  &   74   \\\\ 72&03:35:52 &31:15.0   &    141      &   0.42  &  0.34  &   70   \\\\ 88&03:36:36 &31:04.5   &    134      &   0.25  &  0.34  &   110  \\\\ \\tableline \\sidehead{Stars With A$_V >$ 0.7} 15&03:31:54 & 30:55.8  & 700  &  {2.33}  &     {\\bf1.14}  &   17   \\\\ 24&03:32:44 & 30:30.3   & 300           &  {2.99}  &     0.21  &   114  \\\\ 31&03:33:10 & 30:50.3   & 860           &  {2.18}  &     0.43  &   143  \\\\ 74&03:35:55 & 31:33.8   & 440           &  {1.38}  &     {\\bf1.46}  &   146  \\\\ 92&03:36:42 & 31:31.1    & 1430          &  0.71  &     {\\bf1.12}  &   156  \\\\ 111&03:37:33& 30:58.9  & 264           &  {2.47}  &     {\\bf4.69}  &   145  \\\\ 112&03:37:45& 31:07.0 & 5200  &  {1.24}  &     {\\bf5.06}  &   152  \\\\ 124&03:38:24& 31:12.0  & 245           &  {3.56}  &     {\\bf9.14}  &   150  \\\\ 125&03:38:36& 31:56.5  & 340           &  {1.66}  &     {\\bf1.23}  &   137  \\\\ HD\\,278942&03:39:55  &31:55:33.2 & 207\\tablenotemark{c} &{\\bf7.4}\\tablenotemark{d,\\rm{e}} & {\\bf5.9}\\tablenotemark{d} & 154\\tablenotemark{d}\\\\ \\enddata \\tablenotetext{a}{Notation of \\citetalias{cernis03}} \\tablenotetext{b}{Entries shown in {\\bf bold face} have polarisation percentage $>$ 1\\%.} \\tablenotetext{c}{Hipparcos \\citep{hipparcos}.} \\tablenotetext{d}{\\citet{ander2000}.} \\tablenotetext{e}{Contains a significant circumstellar component \\citep{ander2000}.} \\end{deluxetable}  In figure \\ref{pol} we show polarization vectors from \\citetalias{gbmm90} overlaid on our IRAS column-density and \\xco\\ maps. The weaker polarisation component identified by \\citetalias[red vectors;][]{gbmm90} is clearly aligned along the major axis of the molecular cloud as stated in that paper, but on comparison with our new column-density maps it now becomes evident that much of the high polarisation-percentage component is aligned with the warm shell and associated arm to the west.  \\begin{figure} \\caption{IRAS colum density (left) and extinction from 2MASS (right) overlaid with polarisation vectors from \\citetalias{gbmm90}.  Blue vectors have polarisation strength, P $>$1.2\\% and red vectors have P $<$1.2\\%). The stronger polarisation component (blue) appears aligned with the warm dust emission seen in the IRAS map, while the weaker component mostly lies along the chain of molecular clouds, most evident in the 2MASS extinction map. \\label{pol}} \\end{figure}   In fact, the polarisation strength ($\\la$1.2\\%) measured for the weaker component is typical for nearby molecular clouds \\citep[e.g.][]{gbmm90,bhatt04}. We therefore propose that the weaker component is due to magnetic fields within the molecular cloud material, whilst the stronger component, is due to a swept-up field associated with the shell material. This is consistent with \\citetalias{cernis03}'s measurements, as the distances to the low-polarisation stars (non-bold entries in column 6 of table \\ref{pol_tab}) place them close to the distance of the molecular cloud complex, whilst the higher-polarisation stars (bold entries in column 6 of the table) are generally more distant than 300\\,pc. \\citet{ander2000} measured a polarisation of 5.9$\\pm$0.03\\% at a position angle of 154\\fdg0$\\pm$0\\fdg2 towards HD\\,278942, also consistent with its polarisation being associated with the shell in our scenario (although we note that some fraction of the polarisation is thought to be circumstellar).  Although we do not know its distance, the shell cannot be the source of \\citetalias{cernis03}'s proposed 140\\,pc material, as it must lie {\\em behind} the molecular cloud complex.  We know this, because in figure \\ref{halpha_2mass}, which shows contours of extinction from 2MASS/NICER overlaid on the H$\\alpha$ image, a finger of extinction (contours), is exactly co-incident with a dark ``shadow'' to the southeast in the H$\\alpha$ emission. The same shadow can be seen in the 1907 photograph in Barnard's Atlas \\citep{barnard}, against a circular shaped nebulosity which is located within the circumference of the shell. This would place the shell behind at least the eastern end of the Perseus molecular cloud complex, and therefore more distant than 260$\\pm$20\\,pc, the distance to IC\\,348 \\citep{cernis93} which is embedded in that portion of the molecular cloud\\footnote{This makes either the distance determination to HD\\,278942, or its association with the shell somewhat problematic, as its Hipparcos distance of 207$\\pm$52\\,pc \\citep{hipparcos,ander2000} would place it on the near side of the cloud containing IC\\,348, while its high reddening and polarisation would suggest it is behind a significant column of dense material.}. We also do not believe that the shell could be far behind the molecular cloud complex, as several of the more highly polarised stars in table \\ref{pol_tab} have distances $\\la$300\\,pc.  \\begin{figure} \\caption{ H$\\alpha$ emission \\citep{fink04,vtss}, overlaid with contours of extinction from 2MASS/NICER. The extinction from the molecular clouds obscures the southwest quadrant of the ionized emission, indicating that the H$\\alpha$ emission, and hence the shell, is behind the cloud complex. \\label{halpha_2mass}} \\end{figure}    \\subsection{Is the shell impacting on the molecular cloud?} \\label{impact}  The final piece of the puzzle of the Perseus shell is whether or not it is interacting with the molecular cloud. In earlier sections we have argued that it is likely to be located behind, but still close to the cloud complex. The most compelling evidence for an interaction is given by a comparison of the MSX 8\\,\\um\\ image with extinction from 2MASS as was shown in figure \\ref{msx}.  The shell is clearly detected in 8\\,\\um\\ emission, and in particular has three bright ``knots'' of emission\\footnote{see also \\citealt{kraemer} for images at all 4 MSX bands.}, labelled A, B and C in figure \\ref{msx}). An extinction of A$_V \\ga$ 100 is required to make the cloud opaque at 8\\,\\um, so the bright knots are likely to be real and not just a result of varying extinction in front of the emission.  In fact, the three bright knots are spatially coincident with slight enhancements in the extinction, as indicated by the contours on figure \\ref{msx}. Emission in the 8\\,\\um\\ band of MSX has been modelled as a combination of grey-body dust emission and emission from the unidentified infrared bands (UIBs), usually attributed to polycyclic aromatic hydrocarbons (PAHs) excited in shocks \\citep[e.g.][]{go}, and hence the spatial coincidence between the 8\\um\\ knots and the extinction enhancements are most easily explained by a layer of swept up material in front of a shock, providing compelling evidence for an interaction. Further evidence is seen in the \\xco\\ emission, where the spectrum at ``knot C'' shows a strong second component at 5\\,\\kms\\ (figure \\ref{co_spec}) well separated spatially from emission due to the molecular clouds at those velocities (figure \\ref{co_chann}).  The second component is blueshifted with respect to the ``ambient'' molecular gas, consistent with an interaction with the front side of an expanding shell.  \\begin{figure} \\caption{A ``sampler'' of \\xco\\ spectra from the Perseus region. The bottom spectrum shows the mean spectrum of all positions where \\xco\\ was detected at the 5$\\sigma$ level within the boundaries shown in figure \\ref{co}. This spectrum is multiplied by a factor of 4. Shown above the average spectrum are three spectra from the positions of the 8\\um\\ knots labelled A, B and C in figure \\ref{msx}. The component at 5\\,\\kms\\ in knot C is due to an unusual-looking small feature which appears unconnected either spatially or in velocity to the rest of the cloud complex (also see figure \\ref{co_chann}), and hence may be due to an interation between the shell and the cloud complex. At the top, three spectra from the position of peak \\xco\\ emission in the known star-forming clumps NGC\\,1333 ($\\alpha_{J2000}$=03$^{\\rm h}$29$^{\\rm m}$09\\fs5, $\\delta_{J2000}$=+31\\degrees21$^{\\rm m}$45$^{\\rm s}$), IC\\,348 ($\\alpha_{J2000}$=03$^{\\rm h}$44$^{\\rm m}$03\\fs9, $\\delta_{J2000}$=+32\\degrees04$^{\\rm m}$23$^{\\rm s}$) and B\\,1 ($\\alpha_{J2000}$=03$^{\\rm h}$34$^{\\rm m}$36\\fs0, $\\delta_{J2000}$=+31\\degrees23$^{\\rm m}$34$^{\\rm s}$) are shown for comparison. The dashed lines show the central velocity of the two components of knot C. \\label{co_spec}} \\end{figure} \\begin{figure} \\caption{ The bottom panel shows total integrated \\xco\\ intensity near the shell. The cross indicates the position of HD\\,278942. The three triangles show the positions of the three 8\\um\\ knots which were also indicated in figure \\ref{msx} (A, B and C, from left to right).  The top panel shows channel maps in intervals of 2\\,\\kms\\ of the same region. Velocities are given in the top right corner of each panel. The bright feature at the center of the 4--6\\,\\kms\\ channel (top right) is due to the 5\\,\\kms\\ component seen in the spectrum at the position of 8\\,\\um\\ knot 'C' (figure \\ref{co_spec}), but is not physically connected to the emission associated with the molecular clouds at similar velocities. \\label{co_chann}} \\end{figure}  CO spectra in the Perseus region are extremely complex and confused, with a velocity gradient of almost 10\\,\\kms across the cloud and multiply peaked, optically thick lines throughout. We have attempted to make \\co\\ observations of the top half the shell, where contamination from the molecular clouds should be low in order to perform a more detailed kinematical analysis of the shell, but as can be seen from the 2MASS/NICER extinction data, the column density there is low, and we have been unable to detect \\co\\ at sufficient sensitivity to enable such an analysis.  By assuming a gas-to-dust ratio $\\frac{N_H}{A_V}$, we can calculate the mass of material in the shell from the 2MASS extinction map, using the relation presented by \\citet{dickman}: \\begin{equation} M = (\\alpha d)^2 \\mu \\frac{N_H}{A_V}\\sum_{i} A_V(i), \\end{equation} where $\\alpha$ is the angular size of a map pixel, $d$ is the distance (300\\,pc), $\\mu$ is the mean molecular weight corrected for helium abundance and $\\frac{N_H}{A_V}=1.87\\times 10^{21}$cm$^{-2}$mag$^{-1}$ \\citep{sm79}, although this number could be higher by a factor of $\\sim$2, \\citet[see e.g.][]{gh94}. Due to the contamination from cloud emission in the lower half of the ring, we calculated the mass in the top half of the shell and scaled this by a factor of two to estimate a total mass of the material in the shell of 760\\Msolar. If we assume we are seeing the limb brightened edge of a complete spherical shell then the total mass, $M_{TOT}$ such a shell would contain is given by: \\begin{equation} M_{TOT} = M_{RING} \\left\\{\\frac{R_2^3 - R_1^3}{\\left[R_2^2-R_1^2\\right]^{3/2}}\\right\\}, \\end{equation} where $M_{RING}$ is the mass enclosed within the 2-dimensional annulus indicated in figure \\ref{annulus} and $R_1$ and $R_2$ are the inner and outer radius of that annulus respectively. In this case $R_1 = 2.05$ and $R_2 = 4.29$\\,pc (d=300\\,pc), leading to a total mass of the shell of 1000\\,\\Msolar. Assuming a typical B-star lifetime of 8\\,Myr and a radius of 4.3\\,pc gives a time-averaged expansion velocity of 0.5\\,\\kms\\ for the shell, and a total outward momentum of 520\\,\\Msolar \\kms, much less than 10$^5$\\,\\Msolar \\kms\\ we would estimate based on the properties of the central star (see equation 18 of \\citealt{matzner}). This would suggest that the shell is in the process of merging with the background gas, in which case it is surprising to see such a coherent ring. This discrepancy could be explained if the shell was initally expanding more quickly into a higher density material. For instance, the same diameter could be reached in 1\\,Myr with an expansion velocity of 4\\,\\kms."
    },
    "0601/astro-ph0601001_arXiv.txt": {
        "abstract": "We derive the frequencies of hot Jupiters (HJs) with 3--5 day periods and very hot Jupiters (VHJs) with 1-3 day periods by comparing the planets actually detected in the OGLE-III survey with those predicted by our models.  The models are constructed following \\citet{gouldandmorgan} by populating the line of sight with stars drawn from the {\\it Hipparcos} catalog.  Using these, we demonstrate that the number of stars with sensitivity to HJs and VHJs is only 4--16\\% of those in the OGLE-III fields satisfying the spectroscopic-followup limit of $V_{\\rm max}<17.5$.  Hence, the frequencies we derive are much higher than a naive estimate would indicate. We find that at 90\\% confidence the fraction of stars with planets in the two period ranges is $(1/310)(1^{+1.39}_{-0.59})$ for HJs and $(1/690)(1^{+1.10}_{-0.54})$ for VHJs. The HJ rate is statistically indistinguishable from that found in radial velocity (RV) studies.  However, we note that magnitude-limited RV samples are heavily biased toward metal-rich (hence, planet-bearing) stars, while transit surveys are not, and therefore we expect that more sensitive transit surveys should find a deficit of HJs as compared to RV surveys. The detection of 3 transiting VHJs, all with periods less than 2 days, is marginally consistent with the complete absence of such detections in RV surveys. The planets detected are consistent with being uniformly distributed between 1.00 and 1.25 Jovian radii, but there are too few in the sample to map this distribution in detail. ",
        "introduction": "}   More than 150 extrasolar planets have been detected to date, the majority by radial velocities \\citep{butler02,mayor04}, but also by pulsar timing \\citep{wf92}, microlensing \\citep{bond04,udalski05}, and transits \\citep{bulge,alpha,udalski02c,systematics,konacki03a, konacki,konacki-113,tr-113-132,bouchy,tr-111,alonso04}.  These planets and planetary systems exhibit a wide and unanticipated variety of characteristics, making the individual detections very interesting in themselves, to both scientists and the general public.  However, it is also important to derive from the ensemble of planet detections (and non-detections), the frequency of planets as functions of various parameters, such as mass, semi-major axis, and metalicity. This obviously requires a careful accounting of not only the catalog of detections, but also the stars for which planets would have been detected if they had had planets with various specified characteristics.  In deriving such rates from a given detection technique, the idiosyncrasies of that technique must be taken into account. \\citet{gaudi00} developed a method for deriving rates for microlensing searches, which was first implemented by \\citet{gaudi02} and later improved upon by others (\\citealt{dong05} and references therein). Rates have been derived from radial velocity (RV) detections by \\citet{nelson98}, \\citet{cumming99}, and \\citet{cumming04}, while \\citet{santos05} and \\citet{fischervalenti05} have derived relative frequencies for RV planets as functions of planet and host properties. The detection efficiencies of various surveys for transiting planets in clusters have been calculated by \\citet{gilliland00}, \\citet{mo05}, \\citet{weldrake05}, and \\citet{burke05}.  The specific difficulty in deriving rates from transit surveys of field stars is that the luminosity and radius distributions of the target stars are not known \\citep{gouldandmorgan,brown03}. This problem is made substantially worse by the fact that most transit surveys are directed toward fields in the Galactic plane where there is a huge background of high-luminosity/large-radius sources, which dominate the star counts but whose planets would be almost completely undetectable. These difficulties are exacerbated by strong selection effects in field surveys such that the number of expected detections is a strong function of the planet and host star parameters \\citep{scaling,gaudi05,gaudi05a}.  These obstacles must be overcome if the full value of the field transit surveys is to be realized.  The OGLE survey of Galactic-plane fields toward Baade's Window and Carina raises many questions that are difficult to address in the absence of systematic rate studies.  For example, why does OGLE detect VHJs even though RV surveys do not?  Why does OGLE seem to detect so few HJs compared to RV?  Why does OGLE detect comparable numbers of HJs and VHJs?  Without a well-quantified rate study, it is impossible to determine if these apparent discrepancies are statistically significant and, if so, what is their physical origin.  The only studies to derive quantitative rate information from field transit surveys \\citep{gaudi05,pont} avoided this problem by focusing on the question of relative frequencies of different classes of planets, rather than determining absolute frequencies.  This enabled them to address some but not all of these questions.  Moreover, using relative rates exacerbates uncertainties due to Poisson fluctuations.  Hence, a more sophisticated treatment is needed.  Here we derive absolute frequencies of close-in planets from the OGLE transit survey.  To determine the distributions of stars in the field as a function of luminosity and radius, we build on the work of \\citet{gouldandmorgan}, who used the {\\it Hipparcos} catalog to model local field populations.  We generalize their approach to take account of possible variations in extinction and mean density along various lines of sight.  In \\S\\ref{sec:ogledata}, we review the OGLE-III transit data on which our analysis is based.  In \\S\\ref{sec:philosophy}, we give a broad overview of how transiting planets are selected from the data and how these procedures are simulated in our models.  In \\S\\ref{sec:planetids}, we give a comprehensive summary of the selection procedures, and in \\S\\ref{sec:selection} we give an in depth account of the most important aspects of modeling the resulting selection effects.  In \\S\\ref{sec:calculations}, we present our method for modeling the target populations of the OGLE-III survey, and in \\S\\ref{sec:numprobed} we use our models to evaluate the number of stars actually probed for planets.  This turns out to be far fewer than the number in the field satisfying the survey's magnitude limit. In \\S\\ref{sec:freqplan}, we evaluate the frequency of planets derived from the OGLE-III survey and show that the rates are consistent those derived from RV surveys.  \\S\\ref{sec:consistency} presents two Kolmogorov-Smirnov tests that demonstrate that the distributions of detections as functions of signal-to-noise ratio and $I$-band magnitude are consistent with the model predictions.  \\S\\ref{sec:limbdarkening} gives our prescription for incorporating the effects of limb darkening and ingress/egress.  These effects, which we demonstrate to be only a few percent, are actually incorporated in the main analysis, but their description is deferred to this section for clarity of presentation. In \\S\\ref{sec:resonances} and \\S\\ref{sec:unresolved}, we demonstrate that effects due to ``resonances'' near integer-day periods and due to unresolved binaries are small, of order a few percent.  In \\S\\ref{sec:scaling}, we investigate how well the scaling relations derived analytically by \\citet{scaling} for planet senstivity as functions of planet radius and semimajor axis, apply to the OGLE-III sample.  Finally in \\S\\ref{sec:conclusions}, we summarize our conclusions. ",
        "conclusions": "}  We have measured the frequency of HJs ($3\\,{\\rm days}< P < 5\\,{\\rm days}$) and VHJs ($1\\,{\\rm day}< P < 3\\,{\\rm days}$) by calculating the numbers of stars probed for planets in the OGLE-III Carina and Galactic bulge transit fields, and comparing these to the 5 planets actually detected. At 90\\% confidence, we find that the HJ frequency is $(1/310)(1^{+1.39}_{-0.59})$ and the VHJ frequency is $(1/690)(1^{+1.10}_{-0.54})$. The HJ rate is consistent with the results of RV studies, and the VHJ rate is marginally consistent with the complete lack of RV detections for $P<2\\,{\\rm days}$. Because RV surveys are heavily biased toward metal-rich stars, which are known to harbor substantially more planets than average, we expect them to yield a higher rate of planet detection.  However, the statistics are too sparse yet to be sensitive to this anticipated effect.  We find that the number of stars probed (to the limit set by the followup threshold of $V_{\\rm max}<17.5$) is less than number of stars in the field to this limit by roughly a factor 6 toward Carina and a factor 25 toward the bulge.  The remainder of the stars have physical radii that are too large to permit a detectable transit signal. Hence, a careful determination of the number of {\\it accessible} stars in a given field is required to produce reliable estimates of the planet frequency.  The five OGLE-III planets are consistent with being uniformly distributed between 1.00 and 1.25 Jovian radii, although the number of detections is too small to probe this distribution in detail. The survey has deteriorating sensitivity at smaller radii, but would have detected planets of larger radius, had a substantial such population been present.  We place a 95\\% conficence upper limit of $F<1/400$ on planets with $1.3<r/r_J<1.5$ in the period range of 1--5 days.  Additionally, we have investigated the impact of four effects on the senstivity of the OGLE-III survey to planets with periods of 1--5 days and found this to be of order a few percent in each case. These include limb darkening, ingress/egress, the periodic nature of ground-based observations, and unresolved binaries.  The first two were specifically included in our simulations while the latter two were not.  Finally, we have also shown that the scaling relations predicted by \\citet{scaling} are approximately valid for planets probed in the OGLE-III survey, but that the actual power-law indices can deviate from the predicted ones by up to 25\\% over the parameter ranges we explored."
    },
    "0601/astro-ph0601237_arXiv.txt": {
        "abstract": "{We present new observational results on the RR Lyrae $K$-band Period-Luminosity relation ($PLK$). Data on the Galactic globular clusters NGC 3201 and NGC 4590 (M68), and on the Large Magellanic Cloud cluster Reticulum are shown. We compare the observed slopes of the $PLK$ relations for these three clusters with those predicted by pulsational and evolutionary models, finding a fair  agreement. Trusting on this finding we decided to adopt these theoretical calibrations to estimate the distance to the target clusters, finding a good agreement with optical-based RR Lyrae distances, but with a smaller formal scatter. ",
        "introduction": "RR Lyrae stars are well-known Population II variable stars, widely used as distance indicators, since their mean $V$-magnitude is almost constant and they are bright enough to be detected at moderately large distances. Usually their mean $V$-magnitude is calibrated as a linear funtion of their abundance [Fe/H], but this calibration is still hampered by several theoretical and observational uncertainties (see, e.g., \\citealt{CacciariClementini} and \\citealt{Bonoreview}). However, in the $K$-band the RR Lyrae stars follow a quite tight Period-Luminosity relation ($PLK$, \\citealt{l90}). This relation shows several advantages when compared to the optical calibration: the $PLK$ relation is only moderately affected by evolutionary effects, its intrinsic spread is of the order of 0.03 mag (\\citealt{B01} (B01); \\citealt{B03} (B03)), $K$-band magnitudes are only marginally affected by reddening, and the light curves have low amplitudes, allowing the estimate of the mean magnitude even with a few observations. Moreover, empirical templates \\citep{jones96} can be adopted to improve the estimate of the mean magnitude when a single observation is available. However, theoretical $PLK$ relations show a small dependence on the metallicity, with a slope of the order of $0.2$ mag/dex. We present current results on the Galactic globular clusters NGC 3201 and NGC 4590 (M68), and on the Large Magellanic Cloud (LMC) cluster Reticulum. This work is part of a large observational project, which includes data for 17 Galactic globular clusters, 5 LMC clusters, and field RR Lyrae stars in the Baade Window, in the Sagittarius Stream, and in the main body of the LMC. ",
        "conclusions": "The observed slopes for the three targets are listed in Table 1. Current slopes have been estimated using the entire sample of variables and only for fundamental variables. The slopes predicted by pulsational and evolutionary models are also listed. In Table 2 we compare the estimated distances using the four most recent calibrations of the $PLK$ available in the literature, and for comparison the optical-based distances. Values listed in this table indicate that B03 calibration gives longer distances than B01. This difference can be accounted for with the metallicity coefficient ($+0.167$ and $+0.231$ in the B01 and B03 calibrations, respectively). Finally, we note that $PLK$ distances are characterized by a smaller formal error than optical distances. It is worth noting that the $PLK$ relations estimated by \\citet{santino} and by \\citet{marcio} require knowledge of the HB morphological type. This parameter is completely unknown for field RR Lyrae stars. On the other hand, pulsationally based $PLK$ relations only need knowledge of the metallicity and, once calibrated, can be adopted for individual RR Lyrae stars. To this aim, we have already collected accurate near-infrared data for a large sample of globular clusters that cover a wide range of metal contents and HB morphologies, and firmer conclusions will be drawn when the photometry of this dataset is completed."
    },
    "0601/astro-ph0601570_arXiv.txt": {
        "abstract": "{Possible signatures of primordial magnetic fields on the Cosmic Microwave Background (CMB) temperature and polarization anisotropies are reviewed. The signals that could be searched for include excess temperature anisotropies particularly at small angular scales below the Silk damping scale, B-mode polarization, and non-Gaussian statistics. A field at a few nG level produces temperature anisotropies at the $5 \\mu$K level, and B-mode polarization anisotropies 10 times smaller, and is therefore potentially detectable via the CMB anisotropies. An even smaller field, with $B_0 < 0.1$ nG, could lead to structure formation at high redshift $z > 15$, and hence naturally explain an early re-ionization of the Universe. ",
        "introduction": "Magnetic fields are ubiquitous in astrophysical systems but their origin is still a mystery. One possibility is that they arise due to the dynamo amplification of small seed fields. The dynamo paradigm has been extensively studied (cf. Brandenburg \\& Subramanian 2005 for a recent review), but potential problems remain to be surmounted. These involve the question of whether mean field dynamo coefficients for the galactic dynamo are catastrophically quenched or not, and whether fields generated by the fluctuation dynamo in clusters have a sufficient degree of coherence to explain the observations (Subramanian, Shukurov \\& Haugen 2006; Shukurov, Subramanian \\& Haugen 2006; Schekochihin et al. 2005).  An interesting alternative is that the observed large-scale magnetic fields are a relic from the early Universe, arising perhaps during inflation or some other phase transition (Turner \\& Widrow 1988; Ratra 1992; and reviews by Widrow 2002; Giovannini 2005a). It is well known that scalar (density or potential) perturbations and gravitational waves (or tensor perturbations) can be generated during inflation. Could magnetic field perturbations could also be generated? Indeed mechanisms involving, say, string theory, for the generation of primordial fields have been reviewed in this meeting (Gasperini 2006 and references therein). These generically involve the breaking of the conformal invariance of the electromagnetic action, and the predicted amplitudes are rather model dependent. Nevertheless, if a primordial magnetic field with a present-day strength of $B \\sim 10^{-9}$~G and coherent on Mpc scales is generated, it can strongly influence Galaxy formation. An even weaker field, sheared and amplified due to flux freezing, during galaxy and cluster formation may ease the problems of the dynamo. It is then worth considering if one can detect or constrain such primordial fields.  Here we review possible probes of primordial magnetic fields, which use CMB temperature and polarization anisotropies. Indeed there are also a number of puzzling features associated with current CMB observations, which are not fully understood. These include (a) the excess power in temperature anisotropies detected by the Cosmic Background Imager (CBI) experiment on small angular scales (Readhead et al. 2004), (b) the observations by the Wilkinson Microwave Anisotropy Probe (WMAP) satellite for an unexpectedly high redshift of re-ionization (Kogut et al. 2003), (c) the low CMB quadrupole anisotropy seen compared to theoretical expectations (cf. Spergel et al. 2003), and (d) various asymmetry and alignment effects (cf. de Oliveira-Costa et al. 2004). Such features also encourage us to keep open possibilities of new physical effects, perhaps having to do with primordial fields! ",
        "conclusions": "We briefly reviewed some of the possible ways one could detect/constrain primordial magnetic fields using the CMB anisotropies and polarization. This endeavor would be even more fruitful if there were a compelling mechanism for primordial magnetognesis, which also produced strong enough ($B_{-9} \\sim 1$) and coherent enough (ordered on Mpc scales) fields. At this juncture, theoretical predictions are highly parameter dependent, and so we have taken a more pragmatic approach of assuming that such a field could be generated in the early Universe and asking what it would imply for the CMB and structure formation in general.  For a field of $B \\sim 3$~nG and a nearly scale-invariant spectrum one predicts CMB temperature anisotropies with a $\\Delta T \\sim 5 \\mu$K, at $l< 100$ and $l> 1000$ and polarization anisotropies with $\\Delta T_P \\sim 0.4 \\mu$K at $l > 1000$. Especially interesting is that the vector modes induced by primordial fields can contribute significantly below the Silk scale, where the conventional scalar modes are exponentially damped. Further, the magnetically induced signal at small angular scales will be dominated by B-mode polarization. There do exist intriguing results from the CBI experiment for the presence of a temperature excess at small angular scales. Some part of this excess could arise due to the influence of primordial fields. This excess is conventionally explained as arising due to the SZ effect, but the power in density fluctuations on cluster scales (conventionally measured by $\\sigma_8$), has to be pushed to be in the upper range of values ($\\sigma_8 \\sim 1$), allowed by current CMB and large-scale structure data. Since the SZ signal is frequency dependent, a crucial test would be to compare CMB observations at different frequencies.  Also if B-type polarization at these scales is detected this could be a good indicator of the magnetic field effects. Clearly it will be important to make further observations at small angular scales, especially at different frequencies. It is also important to study the statistics of the CMB anisotropies, since the magnetically induced signals are predicted to be strongly non-Gaussian. If magnetic fields have helicity, this would induce further interesting effects, which have been reviewed in this meeting (Kahniashvilli 2006 and references therein).  We have also emphasized another interesting consequence of the existence of primordial magnetic fields, which indirectly affects the CMB anisotropies. Due to their presence the first collapsed objects of dwarf galaxy masses and smaller can form at high $z > 15$, even for $B \\sim 0.1$~nG. This can potentially lead to the early re-ionization indicated by the present WMAP polarization data in a very natural manner."
    },
    "0601/astro-ph0601093_arXiv.txt": {
        "abstract": "This paper discusses several aspects of current research on high energy emission from supernova remnants, covering the following main topics: 1) The recent evidence for magnetic field amplification near supernova remnant shocks, which makes that cosmic rays are more efficiently accelerated than previously thought. 2) The evidence that ions and electrons in some remnants have very different temperatures, and only equilibrate through Coulomb interactions. 3) The evidence that the explosion that created Cas A was asymmetric, and seems to have involved a jet/counter jet structure. And finally, 4), I will argue that the unremarkable properties of supernova remnants associated with magnetars candidates, suggest that magnetars are not formed from rapidly ($P\\approx 1$~ms) rotating proto-neutron stars, but that it is more likely that they are formed from massive progenitors stars with high magnetic fields. ",
        "introduction": "Supernova remnants (SNRs) are important high energy sources. Not only are they thought to be the major source of Galactic cosmic rays of energies up to at least $10^{15}$~eV, they are, until a supernova occurs in the Galaxy, also the only Galactic sources with which can get a direct view on supernova explosions and nucleosynthesis.  Moreover, SNRs are sources in which interesting physical processes take place, some of which can now be spatially or spectroscopically resolved with the current generation of high spatial and spectral resolution X-ray telescopes, \\chandra, and \\xmm.  \\begin{figure*} \\centerline{ \\parbox{0.45\\textwidth}{ \\includegraphics[width=0.45\\textwidth]{sn1006_acis_rgs_fov.ps}} \\parbox{0.5\\textwidth}{ \\includegraphics[angle=-90, width=0.5\\textwidth]{sn1006_ovii_broadening.eps}} } \\caption{On the left: Map of O\\,VII emission made from several \\chandra-ACIS observations. The lines indicate the region observed by the \\xmm\\ RGS instrument. The target was the bright knot in the northeast. Right: Detail of the RGS1 spectrum of the northeastern knot, showing O\\,VII He$\\alpha$ line emission. The dashed line is the best fit model without line broadening, whereas the solid line shows the model including thermal line broadening \\citep{vink03b}. \\label{fig_sn1006}} \\end{figure*} ",
        "conclusions": "I have shown that X-ray studies of SNR provides us with important information on collisionless shock physics, cosmic ray acceleration and magnetic field amplification by SNRs.  Moreover, by studying SNRs we can learn about the details of the supernova explosions that caused them. For example, the kinematics and spatial distribution of metals in Cas A reveal that the explosion was intrinsically asymmetric, and was accompanied by the emergence of a jet/counter jet system. And I have discussed that the most remarkable property of SNRs associated with magnetars is that they are unremarkable: Their energies are similar to those of other SNRs, which suggests that magnetars were {\\em not} formed from proto-neutron stars with period close to 1~ms."
    },
    "0601/astro-ph0601436_arXiv.txt": {
        "abstract": "I compare the burst detection sensitivity of {\\it CGRO}'s BATSE, {\\it Swift}'s BAT, the {\\it GLAST} Burst Monitor (GBM) and {\\it EXIST} as a function of a burst's spectrum and duration.  A detector's overall burst sensitivity depends on its energy sensitivity and set of accumulations times $\\Delta t$; these two factors shape the detected burst population.  For example, relative to BATSE, the BAT's softer energy band decreases the detection rate of short, hard bursts, while the BAT's longer accumulation times increase the detection rate of long, soft bursts. Consequently, {\\it Swift} is detecting long, low fluence bursts (2-3$\\times$ fainter than BATSE). ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/nlin0601006_arXiv.txt": {
        "abstract": "Explicit solutions for ballistic aggregation of dust-like matter, whose particles stick inelastically upon collisions, are considered. This system provides a model of large-scale structure formation in cosmology within the Zel'dovich approximation.  In particular we show the equivalence of two different representations of solutions proposed in \\citep{S86,ERS96} for a flat 1D flow, extend these representations to cylindrically or spherically symmetric flows, and provide explicit counterexamples showing how exactly these representations break down in the case of non-symmetric flow. ",
        "introduction": "\\label{sec:intro}  We consider here explicit solutions for ballistic aggregation of dust-like matter, with particles that stick upon collisions absolutely inelastically.  This system may be considered as a model for the large structure formation in the Universe within the Zel'dovich approximation~\\citep{Z70}.  First we briefly recall how this model comes about.  Consider a flat Einstein--de Sitter universe containing massive dust-like matter that moves in a self-consistent gravitational field. This matter participates in two types of motion, the homogeneous cosmological Hubble expansion and the evolution of local perturbations of density and velocity, which leads to formation of large-scale inhomogeneities.  To describe the latter kind of dynamics in coordinates comoving with the Hubble expansion, Ya.~B.~Zel'dovich proposed~\\citep{Z70} to use a nonlinear function of time, the perturbation growth factor, instead of the cosmological time.  In the new time, the gravitational force turns out to be approximately balanced by the `Hubble drag' caused by the continuous expansion of the spatial scale.  Within this approximation fluid elements move ballistically, and the evolution of local perturbations is described by the following system of equations: \\begin{gather} \\label{eq:mass} \\pd\\rho t + \\nabla\\cdot(\\rho\\uu) = 0,\\\\ \\label{eq:momentum} \\pd{(\\rho\\uu)}t + \\nabla\\cdot(\\rho\\uu\\otimes\\uu) = 0, \\end{gather} expressing conservation of respectively mass and momentum.  Here \\(t\\) is the new time parameter, \\(\\rho(\\xx, t)\\) is the mass density field, \\(\\uu(\\xx, t)\\) is the velocity field, \\(\\xx\\) is the spatial position comoving with the Hubble expansion, and \\(\\uu\\otimes\\uu\\) denotes the tensor with components~\\(u_iu_j\\).  \\begin{figure} \\centering\\includegraphics[width=6cm]{AGS.1} \\caption{Matter concentration in regions of multi-streamed flow (2D numerical simulations of~\\citep{MS89}).} \\label{fig:aggreg} \\end{figure}  When trajectories of fluid elements cross, the velocity field ceases to be single-valued, and several streams passing through the same spatial location form.  Numerical simulations show (see e.g.\\ \\citep{MS89,WG90}) that the gravitational interaction between streams confines the matter to domains of relatively small width and high density (Fig.~\\ref{fig:aggreg} on page~\\pageref{fig:aggreg}), causing it to aggregate in wall-like, filament-like, and cluster-like structures.  This process is called \\emph{ballistic aggregation}.  In \\citep{GS84,GSS89} (see also the monograph \\cite{GMS92} and \\cite{GM03}) the Burgers equation \\begin{equation*} \\pd\\uu t + (\\uu\\cdot\\nabla)\\uu = \\mu\\Delta\\uu \\end{equation*} was proposed to describe ballistic aggregation.  Indeed, in the vanishing viscosity limit \\(\\mu\\to 0\\) solutions to the Burgers equation develop singularities in the form of infinitesimally thin walls, filaments, and point clusters, whose shape is tantalizingly similar to the large-scale cosmological structure.  However this agreement is limited, as the Burgers equation is written in terms of the velocity field and fails to account for conservation of mass or momentum.  Quantitatively, direct \\(N\\)-body simulations of self-gravitating matter disagree with solutions of the Burgers equation for identical initial data.  These limitations can be overcome in the case of 1D flow with flat symmetry (note that in the 1D case a thorough analytical study is possible even without the aggregation approximation \\cite{AFGM03}). Several authors \\citep{S86,MP94,ERS96,BG98} independently introduced exact solutions for ballistic aggregation in a 1D flow of particles that stick upon collisions with exact conservation of mass and momentum.  In particular, in \\citep{S86} and~\\citep{ERS96} these solutions were expressed in the form of variational principles, so that the density and velocity fields at any given time could be constructed by minimization of certain functionals.  It is the purpose of the present work to explore the possibility of extending these results to the case where there is no flat symmetry. In Section~\\ref{sec:flat} we prove the equivalence of the variational principles proposed in~\\citepalias{ERS96} and~\\citepalias{S86}\\footnote{In the sequel, \\citep{S86} and~\\citep{ERS96} are referred to as \\citepalias{S86} and \\citepalias{ERS96} correspondingly.} in the 1D case with flat symmetry.  To make the exposition self-contained we recall the necessary details of constructions of \\citepalias{S86} and~\\citepalias{ERS96}, in particular because Ref.~\\citepalias{S86} is rather difficult to obtain.  In Section~\\ref{sec:spherical} the variational principles are extended to the case of cylindrical and spherical symmetry.  Both variational principles \\citepalias{S86} and~\\citepalias{ERS96} can be extended formally to the case of 2D or 3D non-symmetric flow, but the corresponding constructions no longer coincide.  Moreover, there exist counterexamples showing that neither variational principle is valid in its extended version already for non-symmetric 2D flows.  The multidimensional extensions of the variational principles \\citepalias{S86} and~\\citepalias{ERS96} and explicit counterexamples are presented in Section~\\ref{sec:counterexamples}.  The second and third authors are grateful to the French Ministry of education for its support.  The work of SG was also supported by RFBR (grant 05--02--16517) and the Russian Leading Schools in Science support program (NSh-5200.2006.2).  The work of AS was supported by RFBR (grant 05--01--00824). ",
        "conclusions": ""
    },
    "0601/astro-ph0601399_arXiv.txt": {
        "abstract": "The development of prebiotic homochirality on early-Earth or another planetary platform may be viewed as a critical phenomenon. It is shown, in the context of spatio-temporal polymerization reaction networks, that environmental effects --- be them temperature surges or other external disruptions --- may destroy any net chirality previously produced. In order to understand the emergence of prebiotic homochirality it is important to model the coupling of polymerization reaction networks to different planetary environments. ",
        "introduction": "According to the prebiotic soup hypothesis \\cite{Bada00},  the early Earth had the supply of organic compounds needed to jump-start polymerization  reactions that, through gradual complexification, led to the first biochemical networks displaying some of the characteristics attributed to life,  such as metabolic activity and replication. Although the road is still obscure \\cite{Orgel98}, the situation is not all bleak. In 1953, Stanley Miller simulated tentative conditions of early-Earth in the laboratory to  obtain amino acids from simple chemical compounds. That same year, Frank \\cite{Frank53} proposed that auto-catalytic polymerization could explain the emergence of biomolecular homochirality, a clear signature of terrestrial and, possibly, all life \\cite{Bonner}: terrestrial amino acids belonging to proteins are overwhelmingly left-handed, while sugars are right-handed.  If a bottom-up approach to the early development of life is adopted,  the homochirality of life's biochemistry must have emerged dynamically, as reactions among the simplest molecular building blocks occurred with high enough yield. Alternatively, one may assume that, somehow, only monomers of a single chirality were present in the prebiotic soup: they were made that way or brought here during the intense bombardment of Earth's infancy, that lasted to about 3.8 Gyr ago \\cite{Gomes05}. We would, however, still need to understand how homochirality developed elsewhere in the cosmos and not here, and whether it developed with the same chiral bias in more than one place.  Here, we consider the homochirality of life as an emergent process that took place on early-Earth and, possibly, other planetary platforms. As a starting point, we use the reaction-network model proposed by Sandars \\cite{Sandars03},  which includes enantiometric cross-inhibition. As shown in the interesting work of Brandenburg and Multam\\\"aki \\cite{BM}  (BM), Sandar's polymerization reaction network can be reduced to an effective spatio-temporal mean-field model, where the order parameter is the chiral asymmetry between left and right-handed polymers. To this, we add the effects of an external environment, showing that they can be crucial in the final determination of the net value of enantiometric excess, if any. ",
        "conclusions": ""
    },
    "0601/astro-ph0601350_arXiv.txt": {
        "abstract": "The ROTSE-IIIc telescope at the H.E.S.S. site, Namibia, obtained the earliest detection of optical emission from a Gamma-Ray Burst (GRB), beginning only 21.8~s from the onset of \\emph{Swift} GRB~050801.  The optical lightcurve does not fade or brighten significantly over the first $\\sim250$ s, after which there is an achromatic break and the lightcurve declines in typical power-law fashion.  The \\emph{Swift}/XRT also obtained early observations starting at 69~s after the burst onset.  The X-ray lightcurve shows the same features as the optical lightcurve.  These correlated variations in the early optical and X-ray emission imply a common origin in space and time.  This behavior is difficult to reconcile with the standard models of early afterglow emission. ",
        "introduction": "Gamma-ray bursts (GRBs) are the most luminous explosions in the universe, but the origin of their emission remains elusive.  With the launch of the \\emph{Swift} $\\gamma$-ray Burst Explorer~\\citep{gcgmn04} in late 2004, great progress has been made in the study of the early afterglow phase of GRBs. However, only a small number of bursts have been imaged simultaneously in both the optical and X-ray bands in the first minutes after the burst~\\citep{nkgpg05,qryaa05,rykaa05,bbbbc05}.  In this letter, we report on the earliest detection of optical emission, starting at 21.8 seconds after the onset of GRB~050801 with the ROTSE-IIIc (Robotic Optical Transient Search Experiment) telescope located at the H.E.S.S. site in Namibia.  This is the most densely sampled early lightcurve yet obtained.  It does not fade or brighten significantly over the first $\\sim250$ seconds, after which there is a break and the lightcurve declines in a typical power-law fashion.  The \\emph{Swift}/XRT also obtained early observations starting at 69 seconds after the burst onset.  The X-ray lightcurve shows the same features as the optical lightcurve.  These correlated variations in the early optical and X-ray emission imply a common origin in space and time.  This behavior differs from that seen in GRB~050319~\\citep{qryaa05}, GRB~050401~\\citep{rykaa05}, and GRB~050525a~\\citep{bbbbc05}.  It is difficult to explain this behavior with standard models of early afterglow emission without assuming there is continuous late time injection of energy into the afterglow. ",
        "conclusions": "In the standard fireball model of GRB afterglow emission, the spectral energy distribution of GRB afterglows can be fit by a broken power-law with spectral segments $F_\\nu \\propto \\nu^\\beta$ (for a review, see \\citet{p05}).  The spectral index obtained by comparing the de-extincted optical (see Figure~\\ref{fig:optandxraylc}) to X-ray flux density during the second XRT integration is $\\beta_{\\mathrm{opt-X}} = -0.92\\pm0.05$, consistent with the X-ray only spectral index of $\\beta_{\\mathrm{X}} = -0.87\\pm0.15$ [0.2-10 keV]. To test for evolution in the broadband spectral index, we have compared the optical and X-ray lightcurves during the first 7000 s (top panel of Figure~\\ref{fig:optandxraylc}).  The optical to X-ray flux ratio is consistent with a constant value ($\\chi^2=15.9$ with 16 degrees of freedom).  Across the break at 250 s, both $\\beta_{\\mathrm{opt-X}}$ and $\\beta_{\\mathrm{X}}$ are unchanged, and therefore the break is achromatic.  Furthermore, there is no evidence of a spectral change in the UVOT images~\\citep{bbhgc05}, although the time resolution is insufficient to constrain the time of the break.  Many X-ray lightcurves have been seen to steepen around 1000 s - 5000 s post-burst with no change in the X-ray spectral index~\\citep{nkgpg05}, For the few bursts with sufficient early optical and X-ray coverage~\\citep{qryaa05, bbbbc05}, this behavior has not been mirrored in the optical band.  The tight correlation between the optical and X-ray emission suggests that they share the same origin in space and time.  The standard fireball model of GRB afterglows can explain the behavior of the optical and X-ray lightcurve after 250~s.  The observed spectral parameters and decay indices are most consistent with a fireball expanding adiabatically into a constant density medium, with the typical synchrotron frequency $\\nu_\\mathrm{m}$ below the optical band, and the cooling frequency $\\nu_\\mathrm{c}$ above the X-ray band. For example, this can be produced by the following parameters: the electron energy index $p = 2.8$; the isotropic equivalent energy $E \\sim 10^{53}\\,\\mathrm{erg}$ at a redshift of $z \\sim 0.5$; the circumburst density $n \\sim 0.7\\,\\mathrm{cm}^{-3}$; the energy fraction in the electrons $\\epsilon_e \\sim 0.07$; and the energy fraction in the magnetic field $\\epsilon_B \\sim 0.0002$. These values of the electrons and magnetic energy are consistent with those deduced for other bursts albeit on the lower side.  If the ejecta were expanding into a $1/r^2$ density profile (a so-called ``wind'' medium), the fireball model predicts a relationship between the spectral and temporal behavior that is inconsistent at the $4\\sigma$ level with the observations after 250~s.  We now investigate the possible explanations of the flat early lightcurve and the origin of the break at 250~s.  First, any spectral transition (e.g. $\\nu_m$ crossing the optical band) would fail to explain the achromatic nature of the break. Achromatic breaks observed in other afterglows have been interpreted as geometric, when the edge of a conical jet becomes visible to the observer and the jet starts to spread~\\citep{hbfsk99,sgkpt99}. At 250 s, this would be the earliest such ``jet break'' detected.  In the fireball model, the post jet break afterglow is expected to decay as $t^{-p}$, where $p$ is the electron energy index with $N_e \\propto E^{-p}$, provided that $p>2$~\\citep{sph99}. A hard electron index of $p<2$ predicts a post-jet decay even steeper than $t^{-p}$~\\citep{dc01}.  Therefore, the observed post-break temporal decay implies $p \\le 1.3$ which predicts a significant pre-break decay~\\citep{dc01} that is inconsistent with the observed pre-break flatness as well as the observed spectral index $\\beta$.  Therefore, the achromatic evolution of GRB~050801 cannot be explained with a jet break.  We have investigated whether the early afterglow is consistent with the predictions of a structured jet viewed off axis~\\citep{gk03}.  In this case, it is difficult to create a sharp early break; under such conditions, the post-break evolution should track closely with the electron energy index $p$, which is inconsistent with observations as described above.  Such an early break at 250~s, can perhaps be explained as the onset time of the afterglow. If the reverse shock is non relativistic (as indicated by the relatively short duration of the burst, see \\citet{sari97}) then self similar expansion starts once the mass collected from the environment is a factor $\\gamma$ smaller than that in the ejecta: \\begin{equation} t_{{\\rm afterglow}}=100\\,{\\rm s}\\,(1+z) \\left( \\frac{E}{10^{53}\\,{\\rm erg}} \\right)^{1/3} \\left( \\frac{n}{1\\,{\\rm cm^{-3}}} \\right) \\left(\\frac{\\gamma}{100} \\right)^{-8/3}. \\end{equation} A value of the initial Lorentz factor $\\gamma$ just below a hundred would therefore be consistent with an onset time of 250~s. However, it is difficult to reconcile the flat part before 250~s as the rise of the afterglow. During the onset, since the fireball is coasting with a constant Lorentz factor, the bolometric luminosity is given by $L_\\mathrm{B} \\propto t^2 n$, the surface area times the density. For a constant density a sharp rise $\\propto t^2$ is therefore expected. For a wind-like decreasing density, the lightcurve should be flat as observed. However, as stated before, a wind density profile seems inconsistent with the behavior after 250~s.  Continuous energy injection has been suggested as a source of early X-ray light\\-curve flat\\-tening~\\citep{nkgpg05}. This injection could be observed if the initial fireball ejecta had a range of Lorentz factors, with the slower shells catching up with the decelerating afterglow~\\citep{rm98,sm00}.  However, we require a very steady injection of energy to produce the observed lightcurve, flat for more than a decade in time.  If we adopt this explanation, the afterglow must start before our first optical observation, implying an initial Lorentz factor of more than 200, and energy injection rate which is roughly constant over a decade in time, and which shuts off suddenly at 250~s.  Flat or very slowly decaying optical lightcurves have been seen in a number of other early afterglows (eg, GRB~030418~\\citep{rspaa04}, GRB~050319~\\citep{qryaa05}, and GRB~041006~\\citep{msmy04,ysr04}).  Early X-ray lightcurves detected by \\emph{Swift} are typically more complex, with rapidly fading sections and short timescale flares~\\citep{nkgpg05}. The early afterglow of GRB~050801, flat in both optical and X-rays, is, so far, unique. It is inconsistent with the standard fireball model for early afterglow emission, unless continuous energy injection is involved. Further \\emph{Swift} prompt GRB detections, combined with rapid follow-up by \\emph{Swift} and ground-based telescopes, will provide further opportunities to explore the origin of this type of early afterglow behavior."
    },
    "0601/astro-ph0601166_arXiv.txt": {
        "abstract": "The relevance of encounters on the destruction of protoplanetary discs in the Orion Nebula Cluster (ONC) is investigated by combining two different types of numerical simulation. First, star-cluster simulations are performed to model the stellar dynamics  of the ONC, the results of which are used to investigate the frequency of encounters,  the mass ratio and separation of the stars involved, and the eccentricity of the encounter orbits. The results show that interactions that could influence the star-surrounding disc are more frequent than previously assumed in the core of the ONC, the so-called Trapezium cluster. Second, a parameter study of star-disc encounters is performed to determine the upper limits of the mass loss of the discs in encounters.  For simulation times of $\\sim$\\,1-2\\,Myr (the likely age of the ONC) the results show that gravitational interaction might account for a significant disc mass loss in dense clusters.  Disc destruction is dominated by encounters with high-mass stars, especially in the Trapezium cluster, where the fraction of discs destroyed due to stellar encounters can reach 10-15\\%. These estimates are in accord with observations of \\cite{lada:aj00} who determined a stellar disc fraction of 80-85\\%. Thus, it is shown that in the ONC - a typical star-forming region - stellar encounters do have a significant effect on the mass of protoplanetary discs and thus affect the formation of planetary systems. ",
        "introduction": "According to current knowledge planetary systems form from the accretion discs around young stars. These young stars are in most cases not isolated but are part of a cluster. It is still an open question as to how far encounters with the surrounding stars of the cluster influence planet formation. The fact is that these discs disperse over time either by photo-evaporation, encounter-induced disc mass loss or some other means. Simple calculations seem to indicate that encounters do not play an important role. Nevertheless, \\cite{scally:mnras01} found that a little less than a third of stars in the central core of the ONC would suffer an encounter within 100\\,AU, and thus, for a significant minority of stars in the Trapezium cluster, star disk encounters are of some importance. However, in order to quantify how many stars would be expected to lose most of their disc material, it is necessary to treat the disc mass loss - and its dependence on the mass of the perturbing object - in a more sophisticated way than was done by \\cite{scally:mnras01}.  Like \\cite{scally:mnras01} the ONC is used as model cluster for the following reasons: The ONC is thought to be a typical environment for star-formation and its high density suggests that stellar encounters might be relevant for the evolution of circumstellar discs. In addition, it is one of the best-studied regions in our galaxy, so that observational constraints significantly reduce the modelling parameters. \\subsection{Structure and Dynamics of the ONC}  The ONC is a rich stellar cluster with about 4000 members with masses $M^*$ above 0.08\\,\\Msun\\ in a volume $\\sim$ 5\\,pc across \\citep{hillenbrand:apj98,hillenbrand:apj00}. Most of the objects are T~Tauri stars, but there is also strong evidence for the existence of several protostars. The mean stellar mass is about 0.5\\,\\Msun\\ \\citep{scally:mnras05}, the half-mass radius roughly 1\\,pc. Recent studies on the stellar mass distribution \\citep{hillenbrand:apj00,luhman:apj00,muench:apj02,slesnick:apj04} reveal no significant deviation from the field star IMF \\cite{kroupa:mnras93} \\be \\xi(M^*)= \\left\\{ \\begin{array}{lrll} 0.035 M^*\\,^{-1.3} & \\mbox{ if } & 0.08&\\le M^*\\,<0.5, \\\\ 0.019 M^*\\,^{-2.2} & \\mbox{ if } & 0.5&\\le M^*\\,<1.0, \\\\ 0.019 M^*\\,^{-2.7} & \\mbox{ if } & 1.0&\\le M^*\\,<\\infty. \\end{array} \\right. \\label{eq:imf} \\ee  The shape of the system is not perfectly spherical but elongated in the north-south direction. The reason for this asymmetry is the gravitational potential of a massive molecular ridge in the background of the cluster, OMC~1, being part of the much larger complex of the Orion Molecular Cloud. The mean age of the whole cluster has been estimated to be about 1-2\\,Myr, though with a significant age spread of the individual stars. Today, star formation is no longer occurring in the cluster itself, only in the background molecular cloud. After a short period of intense star-formation the ONC has expelled most of the residual gas by now.  The density and velocity distribution of the ONC resembles an isothermal sphere. From the outer edge the number of stars falls linearly with decreasing radius $r$ down to $\\sim$\\,0.1\\,pc; inside this cluster core, the distribution function becomes flatter \\citep{jones:aj88,mccaughrean:aj94,hillenbrand:aj97,hillenbrand:apj98,mccaughrean:msngr02}. The central number density $\\rho_{\\rm{core}}$ in the inner 0.053\\,pc reaches $4.7\\times10^4\\,\\mbox{pc}^{-3}$ \\citep{mccaughrean:aj94} and makes the ONC one of the densest star forming regions in the Galaxy. The velocity dispersion is nearly constant for all cluster radii. In their proper motion study of the ONC, \\cite{jones:aj88} have obtained a three-dimensional velocity dispersion of $4.3\\,$km\\,s$^{-1}$, thus the crossing time is $ t_{\\rm{cross}}=2R_{\\rm{hm}}/{\\sigma}\\approx0.5~\\,\\mbox{Myr}$. Another crucial quantity in stellar dynamics is the virial ratio $Q_{\\rm{vir}}$, which  has been estimated for the ONC as \\begin{eqnarray} Q_{\\rm{vir}}=\\frac{R_{\\rm{hm}}\\sigma^2}{2GM}\\approx1.5~, \\end{eqnarray} where $R_{\\rm{hm}}$ is its half-mass radius. This indicates that the ONC is not only far from virial equilibrium, but even seems to be gravitationally unbound ($Q_{\\rm{vir}}>1)$. However, this statement has to be treated with care because errors in the observational parameters can easily account for an error of over 50 per cent in this calculation. Besides, the estimated total mass of the ONC of 2000\\,\\Msun\\ is only a lower limit since a substantial amount could be present in undetected low-mass binary companions or gas. Furthermore, the contribution of the OMC~1 to the overall gravitational potential is still unknown and the elongated shape of the cluster indicates that it is not negligible. However, as long as measurements of velocities and masses lack higher precision, one cannot constrain the cluster dynamics to contraction, equilibrium or expansion.  Like many other stellar aggregations, the ONC shows mass segregation, with the most massive stars being confined to the inner cluster parts. The Trapezium cluster, a subgroup of about 1000 stars in a volume 0.6\\,pc across, represents the denser core of the Orion Nebula Cluster. It contains four luminous O and B stars at the very center, designated as the Trapezium. Their most prominent member, \\COri, is a O6 star with a mass of about 50\\,\\Msun, a luminosity of $4\\times10^5\\,\\Lsun$ and a surface temperature of $4\\times10^4\\mbox{K}$.     Apart from its high density and young age, the evidence for protoplanetary discs around many stars in this cluster makes the ONC the ideal candidate for the present investigation. Whereas the first identification of ``peculiar stellar objects'' dates back to the paper from \\citet{laques:aap79}, it took more than a decade to recognize them as circumstellar discs which are ionized by the intense radiation of the Trapezium stars. \\citet{o'dell:apj93} designated these bright objects as ``proplyds''. At greater distances from the cluster centre they also detected their dark counterparts: discs in silhouette which are visible due to the bright nebular background. Thus far, about 200 bright proplyds and 15 silhouette discs have been revealed in several HST studies of the ONC \\citep{o'dell:apj93,o'dell:aj96,bally:aj98a,bally:aj00}, nearly all located in the Trapezium cluster due to selection effects.  The most recent study on circumstellar discs in the Trapezium \\citep{lada:aj00} used the L-band excess as detection criterium. They analyze 391 stars and find a fraction of 80-85\\% to be surrounded by discs. This is in agreement with an earlier investigation of the larger ONC in which \\citet{hillenbrand:aj97} states a disc fraction of 50-90\\%. The disc sizes established so far vary between 50\\,AU and 1000\\,AU, with a typical value of 200\\,AU for low-mass stars. The inferred disc masses are only accurate to an order of magnitude but seem not to exceed a few percent of the central stellar mass, which classifies them as low-mass discs. The disc surface densities are generally well described by power-law profiles, $\\Sigma(r) \\propto r^{-a}$, with $0.5\\le{a}\\le1.5$.   \\subsection{Previous work}  At present it is not possible to perform numerical particle simulations where the stars including their surrounding discs are sufficiently resolved to determine the effect of encounters on the discs quantitatively. Therefore \\cite{scally:mnras01} treated the dynamics of the stars in the cluster and star-disc encounters and photoevaporation effects in separate investigations and combined the results to determine the disc destruction rate. The term star-disc encounters means here encounters where only one of the stars is surrounded by a disc, in contrast to disc-disc encounters denominating encounters where both stars are surrounded by discs.  In order to determine the disc destruction rate a number of assumptions were made by \\cite{scally:mnras01}:  \\begin{enumerate} \\item The discs around the stars do not alter the stellar dynamics in any significant way. \\item The discs are of low mass. \\item Only two stars are involved in an encounter event and three-(or even more) body events are so rare that they can be neglected. \\item The encounters were modelled as coplanar and prograde. \\item The disc mass loss is deduced from parameter studies of star-disc encounters. \\item The closest encounter is the most destructive one. \\item A parameter study of encounters between equal mass stars is used.  \\end{enumerate}  From the last point \\cite{scally:mnras01} concluded that only in penetrating encounters a significant amount of mass can be stripped from the disc. Assuming a typical disc size of 200\\,AU, they considered the stars with separations of less than 100\\,AU. Under these assumptions \\cite{scally:mnras01} found that just 3-4\\% of all discs have the potential to be destroyed by encounters.  \\subsection{Aim and structure of present work}  In this work we concentrate on the effect of encounters on the disc dispersal and ignore photo-evaporation. We performed an investigation similar to that of \\cite{scally:mnras01} but drop the last two assumptions of above list.  In contrast to \\cite{scally:mnras01} i) we record the entire path history of each star, so the effect of repeated encounters is included, ii) extend the parameter study to include encounters with massive perturbers and iii) record the most forceful encounter. It will be demonstrated that these modifications alter significantly the result, so that the conclusion that encounters can be excluded as a disc destruction mechanism should be revised.  The paper is organized the following way: In Section~2 we begin with simulations of the ONC. This is  followed by an investigation of the mass loss in star-disc encounters in Section~3, resulting in a fitting formula for the disc mass loss. In Section~4 the results of Section~2 and 3 are combined to determine the disc dispersal as a function of time. In Section~5 it will be discussed how the disc mass loss is influenced by the assumptions made. ",
        "conclusions": "This investigation for the ONC cluster, combining cluster simulations with encounter investigations, shows that potentially up to 10-15\\% of the discs in the Trapezium cluster could have been destroyed by encounters. Our more sophisticated treatment of disc mass loss expected from multiple stellar encounters implies that it is plausible that the 15-20\\% of discless stars observed in the Trapezium \\citep{lada:aj00} may, in a large fraction of cases, result from star-disc collisions.    One important result is that the most massive bodies dominate the disc mass loss, with significant interaction even beyond a separation of ten disc radii for a ONC-like entity. This is particulary so for the Trapezium, where some dozen massive stars are surrounded by hundreds of lighter bodies. Consequently, it is the upper end of a cluster's mass distribution that to a large degree determines the fate of the circumstellar discs in its vicinity and thus there are in principle two quantities that are mainly regulating the effect of stellar encounters on the mass-loss from protoplanetary discs: namely the local stellar density which determines the encounter probability, and the upper limit of the mass range, which affects the maximum strength of the perturbing force.  As the number of massive members in a stellar group seems to be correlated to its initial density \\citep[see][]{testi:aap97,bonnel:mnras04} and the IMF appears to be uniform for all Galactic environments \\citep[e.g][]{kroupa:mnras93,muench:apj00} of star-formation, the dependency of the disc mass loss due to encounters is mainly reduced to one parameter, namely the density distribution of the considered stellar system. This relation would be worth investigating in future.  As the possible disc destruction rate is so high, it is obvious that the remaining discs are considerably affected by encounters - that is the mass distribution and the cut-off radius of the disc. These properties are two ingredients vital for the understanding of the formation of planetary systems.  Further studies of the temporal evolution of these properties in the ONC are on the way."
    },
    "0601/astro-ph0601216_arXiv.txt": {
        "abstract": "Gamma-Ray Bursts (GRBs) are unique probes of the cosmic star formation history and the state of the intergalactic medium up to the redshifts of the first stars. In particular, the ongoing {\\it Swift} mission might be the first observatory to detect individual Population~III stars, provided that the massive, metal-free stars were able to trigger GRBs. {\\it Swift} will empirically constrain the redshift at which Population~III star formation was terminated, thus providing crucial input to models of cosmic reionization and metal enrichment. ",
        "introduction": "Gamma-Ray Bursts (GRBs) are believed to originate in compact remnants (neutron stars or black holes) of massive stars. Their high luminosities make them detectable out to the edge of the visible universe \\cite{CL00,LR00}. GRBs offer the opportunity to detect the most distant (and hence earliest) population of massive stars, the so-called Population~III (Pop~III henceforth), one star at a time.  In the hierarchical assembly process of halos which are dominated by cold dark matter (CDM), the first galaxies should have had lower masses (and lower stellar luminosities) than their low-redshift counterparts.  Consequently, the characteristic luminosity of galaxies or quasars is expected to decline with increasing redshift. GRB afterglows, which already produce a peak flux comparable to that of quasars or starburst galaxies at $z\\sim 1-2$, are therefore expected to outshine any competing source at the highest redshifts, when the first dwarf galaxies have formed in the universe.  The first-year polarization data from the {\\it Wilkinson Microwave Anisotropy Probe} ({\\it WMAP}) indicates an optical depth to electron scattering of $\\sim 17\\pm 4$\\% after cosmological recombination \\cite{Kog03,Spe03}. This implies that the first stars must have formed at a redshift $z\\sim 20$, and reionized a substantial fraction of the intergalactic hydrogen around that time \\cite{C03,CFW03,SL03,WL03,YBH04}. Early reionization can be achieved with plausible star formation parameters in the standard $\\Lambda$CDM cosmology; in fact, the required optical depth can be achieved in a variety of very different ionization histories since {\\it WMAP} places only an integral constraint on these histories \\cite{HH03}. One would like to probe the full history of reionization in order to disentangle the properties and formation history of the stars that are responsible for it. GRB afterglows offer the opportunity to detect stars as well as to probe the ionization state \\cite{BarL04} and metal enrichment level \\cite{FL03} of the intervening intergalactic medium (IGM).  GRBs, the electromagnetically-brightest explosions in the universe, should be detectable out to redshifts $z>10$ \\cite{CL00,LR00}. High-redshift GRBs can be identified through infrared photometry, based on the Ly$\\alpha$ break induced by absorption of their spectrum at wavelengths below $1.216\\, \\mu {\\rm m}\\, [(1+z)/10]$. Follow-up spectroscopy of high-redshift candidates can then be performed on a 10-meter-class telescope. Recently, the ongoing {\\it Swift} mission \\cite{Geh04} has detected a GRB originating at $z\\simeq 6.3$ (e.g., \\cite{Hai05}), thus demonstrating the viability of GRBs as probes of the early universe.  There are four main advantages of GRBs relative to traditional cosmic sources such as quasars:  \\noindent {\\it (i)} The GRB afterglow flux at a given observed time lag after the $\\gamma$-ray trigger is not expected to fade significantly with increasing redshift, since higher redshifts translate to earlier times in the source frame, during which the afterglow is intrinsically brighter \\cite{CL00}. For standard afterglow lightcurves and spectra, the increase in the luminosity distance with redshift is compensated by this {\\it cosmological time-stretching} effect.  \\noindent {\\it (ii)} As already mentioned, in the standard $\\Lambda$CDM cosmology, galaxies form hierarchically, starting from small masses and increasing their average mass with cosmic time. Hence, the characteristic mass of quasar black holes and the total stellar mass of a galaxy were smaller at higher redshifts, making these sources intrinsically fainter \\cite{WL02}.  However, GRBs are believed to originate from a stellar mass progenitor and so the intrinsic luminosity of their engine should not depend on the mass of their host galaxy. GRB afterglows are therefore expected to outshine their host galaxies by a factor that gets larger with increasing redshift.  \\noindent {\\it (iii)} Since the progenitors of GRBs are believed to be stellar, they likely originate in the most common star-forming galaxies at a given redshift rather than in the most massive host galaxies, as is the case for bright quasars \\cite{BarL04}. Low-mass host galaxies induce only a weak ionization effect on the surrounding IGM and do not greatly perturb the Hubble flow around them. Hence, the Ly$\\alpha$ damping wing should be closer to the idealized unperturbed IGM case and its detailed spectral shape should be easier to interpret. Note also that unlike the case of a quasar, a GRB afterglow can itself ionize at most $\\sim 4\\times 10^4 E_{51} M_\\odot$ of hydrogen if its UV energy is $E_{51}$ in units of $10^{51}$ ergs (based on the available number of ionizing photons), and so it should have a negligible cosmic effect on the surrounding IGM.  \\noindent {\\it (iv)} GRB afterglows have smooth (broken power-law) continuum spectra unlike quasars which show strong spectral features (such as broad emission lines or the so-called ``blue bump'') that complicate the extraction of IGM absorption features. In particular, the continuum extrapolation into the Ly$\\alpha$ damping wing (the Gunn-Peterson absorption trough) during the epoch of reionization is much more straightforward for the smooth UV spectra of GRB afterglows than for quasars with an underlying broad Ly$\\alpha$ emission line \\cite{BarL04}.  Although the nature of the central engine that powers the relativistic jets of GRBs is still unknown, recent evidence indicates that long-duration GRBs trace the formation of massive stars (e.g., \\cite{T97,Wij98,BN00,Kul00,BKD02,Nat05}) and in particular that long-duration GRBs are associated with Type Ib/c supernovae \\cite{Sta03}. Since the first stars in the universe are predicted to be predominantly massive \\cite{ABN02,BCL02,BLar04}, their death might give rise to large numbers of GRBs at high redshifts.  In contrast to quasars of comparable brightness, GRB afterglows are short-lived and release $\\sim 10$ orders of magnitude less energy into the surrounding IGM. Beyond the scale of their host galaxy, they have a negligible effect on their cosmological environment\\footnote{Note, however, that feedback from a single GRB or supernova on the gas confined within early dwarf galaxies could be dramatic, since the binding energy of most galaxies at $z>10$ is lower than $10^{51}~{\\rm ergs}$ \\cite{BarL01}.}. Consequently, they are ideal probes of the IGM during the reionization epoch.  Their rest-frame UV spectra can be used to probe the ionization state of the IGM through the spectral shape of the Gunn-Peterson (Ly$\\alpha$) absorption trough, or its metal enrichment history through the intersection of enriched bubbles of supernova (SN) ejecta from early galaxies \\cite{FL03}.  Afterglows that are unusually bright ($>10$mJy) at radio frequencies should also show a detectable forest of 21~cm absorption lines due to enhanced HI column densities in sheets, filaments, and collapsed minihalos within the IGM \\cite{FL02}.  Another advantage of GRB afterglows is that once they fade away, one may search for their host galaxies. Hence, GRBs may serve as signposts of the earliest dwarf galaxies that are otherwise too faint or rare on their own for a dedicated search to find them. Detection of metal absorption lines from the host galaxy in the afterglow spectrum, offers an unusual opportunity to study the physical conditions (temperature, metallicity, ionization state, and kinematics) in the interstellar medium of these high-redshift galaxies. A small fraction ($\\sim 10$) of the GRB afterglows are expected to originate at redshifts $z>5$ \\cite{BL02,BL06}.  This subset of afterglows can be selected photometrically using a small telescope, based on the Ly$\\alpha$ break at a wavelength of $1.216\\, \\mu {\\rm m}\\, [(1+z)/10]$, caused by intergalactic HI absorption.  The challenge in the upcoming years will be to follow-up on these candidates spectroscopically, using a large (10-meter class) telescope.  GRB afterglows are likely to revolutionize observational cosmology and replace traditional sources like quasars, as probes of the IGM at $z>5$.  The near future promises to be exciting for GRB astronomy as well as for studies of the high-redshift universe. ",
        "conclusions": ""
    },
    "0601/astro-ph0601020_arXiv.txt": {
        "abstract": "Quasi-Periodic Oscillations (QPOs) have been observed during three powerful magnetar flares, from SGR0526-66, SGR1806-20 and SGR1900+14. These QPOs have been commonly interpreted as being driven by the mechanical modes of the magnetar's solid crust which are excited during the flare. Here we show that this interpretation is in sharp contradiction with the conventional magnetar model. Firstly, we show that a magnetar crustal mode decays on the timescale of at most a second due to the emission of Alfven waves into the neutron-star interior. A possible modification is then to assume that the QPOs are associated with  the magnetars' {\\it global}  modes. However, we argue that at the frequencies of the observed QPOs, the neutron-star core is likely to support a continuum of Magneto-Hydrodynamic (MHD) normal modes. We demonstrate this on a completely solvable toy model which captures the essential physics of the system. We then show that the frequency of the global mode of the whole star is likely to have a significant imaginary component, and its amplitude is likely to decay on a short timescale. This is not observed. Thus we conclude that either (i) the origin of the QPO is in the magnetar's magnetosphere, or (ii) the magnetic field has a special configuration: either it is expelled from the magentar's core prior to the flares, or it's poloidal component has very small coherence length. ",
        "introduction": "In a prophetic paper, Duncan and Thompson (1992) argued that magnetars---a class of neutron stars with  super-strong ($10^{14}$---$10^{15}$G) magnetic fields---must exist. Treated at first with some skepticism, the magnetar paradigm has proved extremely successful in explaining the rich phenomenology of Anomalous X-ray Pulsars (AXPs) and Soft Gamma-ray Repeaters (SGRs) (e.g., Thompson and Duncan 1995, TD). In particular one phenomenon which has an appealing explanation within the magnetar paradigm are Giant Flares from SGRs. They are thought to be powered by an impulsive release of  magnetic energy stored in the neutron star. The release may be triggered by a fracture in the magnetically-stressed crust (TD) or by a sudden reconnection in a twisted magnetosphere (Lyutikov 2003). So far 3 Giant Flares have been observed: from SGR0526-66 (Mazets et al.~1979), SGR1900+14 (Hurley et al~1999, Feroci et al.~1999), and SGR1806-20 (Hurley et al.~2005, Palmer et al.~2005). In each of the flares' light-curves, there is intriguing evidence for  QPOs with frequencies of tens of Hertz. Firstly, the $43.5$Hz QPO was reported by Barat et al.~(1983) in the 1979 flare from SGR0526-66. Then recently Israel et al.~(2005) have found several QPOs in the 2004 Giant Flare from SGR1806-20, with frequencies from $18$ to $90$ Hz. These QPOs were detected with very high signal-to-noise and lasted for $\\sim 100$ seconds; their presence was recently confirmed by Watts and Strohmayer (2005, WT) who used  data from a different satellite. WT have also detected a relatively high-frequency QPO, of $625$Hz. Lastly, prompted by the Israel et al.~observations, Strohmayer and Watts (2005) have found multiple QPO in the light-curve of the 1998 Giant Flare from SGR1900+14.  The QPOs have been widely associated with  torsional modes of the neutron-star crust [Barat et al.~(1983) made the first suggestion, followed by a more detailed analysis of Duncan (1998), Israel et al.~(2005), Strohmayer and Watts (2005) and Piro (2005)]. Indeed, it is attractive to associate a stable QPO with a mechanical mode of a neutron star, and the frequencies of crustal torsional modes are of the right order of magnitude\\footnote{The $625$Hz QPO found by WT is already somewhat problematic for the crustal-mode picture. One can associate it with a crustal shear mode which has 1 radial node (n=1; see e.g.~Piro 2005). However, a large multitude of $l<5$, $n=1$ shear modes exist around that frequency, with spacings of several Hz. There is no physical reason why just one of them should be excited, and not many.} However,  we  argue below that this interpretation is not viable within the magnetar paradigm. In the next section, we show that the crustal mode of a magnetar is not stable but decays within seconds by emitting Alfven waves into the core. This in direct contradiction with the observations. Chris Thompson (private communications) has suggested that one may associate the QPOs with the  MHD-elastic modes of the {\\it whole} neutron star. However, in section 3 we argue that the frequency of such global mode is likely to have an imaginary component, and  the mode decays on a short timescale. This is due to the continuous spectrum of MHD modes supported by the liquid core. ",
        "conclusions": ""
    },
    "0601/astro-ph0601034_arXiv.txt": {
        "abstract": "We present Keck Interferometer observations of TW Hya that spatially resolve its emission at 2 $\\mu$m wavelength.  Analyzing these data together with existing $K$-band veiling and near-infrared photometric measurements, we conclude that the inner disk consists of optically thin, sub-micron-sized dust extending from $\\sim 4$ AU to within 0.06 AU of the central star.  The inner disk edge may be magnetospherically truncated. Even if we account for the presence of gas in the inner disk, these small dust grains have survival times against radiation blow-out that are orders of magnitude shorter than the age of the system,  suggesting continual replenishment through collisions of larger bodies. ",
        "introduction": "TW Hya is a nearby \\citep[$\\sim 51$ pc;][]{MAMAJEK05}, young \\citep[$\\sim 10$ Myr;][]{WEBB+99} star surrounded by an accretion disk that evinces a large inner hole as judged from the observed spectral energy distribution \\citep[SED;][]{CALVET+02}.  Unusually low excess emission at wavelengths $\\lambda \\la 10$ $\\mu$m can be modeled with an optically thick disk whose inner edge is located $\\sim 4$ AU from the central star. While observations of 10 $\\mu$m silicate emission \\citep{SLR00,UCHIDA+04} together with non-zero excess at 2 $\\mu$m \\citep{JV01} suggest the presence of at least some dust grains with sizes less than a few microns at stellocentric distances $R \\la 4$ AU, this inner disk material appears optically thin, and has been estimated to constitute less than a lunar mass \\citep{CALVET+02}. Detection of warm gas \\citep{HERCZEG+04,RETTIG+04} and accretion signatures \\citep{MUZEROLLE+00,AB02} confirm that the region inside 4 AU is not devoid of material.  At $R \\sim 4$ AU, dust temperatures ($\\sim 100$ K) are substantially lower than sublimation temperatures for silicate grains \\citep[$\\ga 1500$ K; e.g.,][]{POLLACK+94}. This suggests the optically thick outer disk is truncated at 4 AU by a mechanism other than dust sublimation.  Large holes inferred from SEDs are commonly attributed to planets, which may clear gaps about their orbits. A planet impedes accretion of material outside its orbit, while inner disk material is free to drain onto the central star \\citep{GT82,BRYDEN+99,RICE+03}.  However, a viscous outer disk causes inward migration of planets and their associated gaps \\citep[e.g.,][]{LP86,WARD97}. Thus,  inner holes would be filled in on the viscous timescale unless the disk is less massive than the planet, in which case the timescale is lengthened by the mass ratio between the planet and the disk \\citep{CHIANG03}. Since the outer disk of TW Hya is massive \\citep[$\\ga 0.1$ M$_{\\odot}$;][]{WEINBERGER+02,WILNER+05}, planets are unlikely to preserve inner clearings over the lifetime of the system ($\\sim 10$ Myr) unless the outer disk is unusually inviscid \\citep[$\\alpha \\la 10^{-5}$;][]{SS73}.  An alternative explanation for SED-inferred holes is that dust grains have grown larger than a few microns, depleting the population of small grains that would produce the near-IR emission.  Scattered light and long-wavelength emission from the outer ($R > 10$ AU) disk of TW Hya suggest substantial grain growth, up to cm sizes \\citep{WEINBERGER+02,WILNER+05}, supporting the hypothesis of grain coagulation.  The small amount of sub-micron-sized dust required to explain emission at $\\lambda \\la 10$ $\\mu$m and the spectral shape of the 10 $\\mu$m silicate feature \\citep{CALVET+02,UCHIDA+04} may represent the tail at small sizes of a grain size distribution that peaks at sizes much larger than a micron.  Furthermore, as we discuss below, small dust grains are short-lived in the TW Hya disk and demand continual replenishment, possibly from collisions of larger parent bodies. The population of sub-micron-sized dust in the inner disk may not be primordial.  Here we present observations with the Keck Interferometer that spatially resolve the inner disk around TW Hya for the first time. Previous observations at sub-millimeter to centimeter wavelengths spatially resolved emission at larger radii ($> 10$ AU) and enabled powerful constraints on the outer disk structure and dust properties \\citep{QI+04,WILNER+05}. With our near-IR interferometric observations, we extend this analysis to the inner disk.  By combining spatially resolved measurements with spectral information, we determine the radial distribution, temperature, and approximate grain sizes of dust.  We confirm that the inner disk is populated by small amounts of sub-micron-sized dust, and show that the inner radius\\footnote{In the remainder of the paper, the ``inner radius'' refers to the inner edge of the optically thin inner disk, not the inner edge of the optically thick outer disk. We take the latter to be located at $R \\sim 4$ AU based on previous modeling by \\citet{CALVET+02}.} of this optically thin disk occurs further from the star than previous spatially unresolved observations imply. ",
        "conclusions": "We observed TW Hya with the Keck Interferometer and found the 2 $\\mu$m emission to be spatially resolved.  We modeled the interferometric data together with previous measurements of the $K$-band veiling and near-IR fluxes, and inferred that the inner disk consists of optically thin dust extending from the edge of the optically thick outer disk \\citep[$R \\sim 4$ AU;][]{CALVET+02} to $R_{\\rm in} = 0.06$ AU of the central star. This inner radius is larger than expected from dust sublimation; the truncation may be magnetospheric in origin.  The near-IR emitting dust is composed of sub-micron-sized particles which are extremely short-lived; this dust may be replenished by erosive collisions of larger parent bodies in the inner disk.  \\medskip \\noindent{\\bf Acknowledgments.} The near-IR interferometry data presented in this paper were obtained with the Keck Interferometer (KI) of the W.M. Keck Observatory, which was made possible by the generous financial support of the W.M. Keck Foundation and is operated as a scientific partnership between the California Institute of Technology, the University of California, and NASA. The authors thank the entire KI team for making these observations possible, and acknowledge the cultural role and reverence that the summit of Mauna Kea has always had within the indigenous Hawaiian community. The authors are grateful to M. Sitko for providing the 3--5 $\\mu$m photometry used in this paper, and to J. Najita and J. Carr for useful comments about the warm CO gas."
    },
    "0601/astro-ph0601202_arXiv.txt": {
        "abstract": "In this paper, we address whether the growth of supermassive black-holes has kept pace with the process of galaxy assembly. For this purpose, we first searched  the Hubble Ultra Deep Field (HUDF)  for ``tadpole galaxies'', which have a knot at one end plus an extended tail. They appear dynamically unrelaxed --- presumably early-stage mergers --- and make up $\\sim$6\\% of the field galaxy population. Their redshift distribution follows that of field galaxies, indicating that --- if tadpole galaxies are indeed dynamically young --- the process of galaxy assembly generally kept up with the reservoir of field galaxies as a function of epoch.  \\n Next, we present a search for HUDF objects with point-source components that are optically variable (at the \\cge 3.0$\\sigma$ level) on timescales of weeks--months. Among 4644 objects to \\iAB$\\simeq$28.0 mag (10$\\sigma$), 45 have variable point-like components, which are likely weak AGN. About $\\sim$1\\% of all field objects show variability for 0.1\\cle z\\cle 4.5, and their redshift distribution is similar to that of field galaxies. Hence supermassive black-hole growth in weak AGN likely also kept up with the process of galaxy assembly. However, the faint AGN sample has almost no overlap with the tadpole sample, which was predicted by recent hydrodynamical numerical simulations. This suggests that tadpole galaxies are early-stage mergers, which likely preceded the ``turn-on'' of the AGN component and the onset of visible point-source variability by \\cge 1 Gyr. \\begin{keyword} galaxies: mergers \\sep galaxies: formation \\sep galaxies: active galactic nuclei \\sep Supermassive Black Holes \\end{keyword} ",
        "introduction": "Introduction} \\label{introduction}  From the WMAP polarization results (Kogut \\etal 2003), population III stars likely existed at z$\\simeq$20. These massive stars (\\cge 250\\Mo) are expected to produce a large population of black holes (BH; \\Mbh\\cge 150\\Mo; Madau \\& Rees 2001). Since there is now good dynamical evidence for the existence of supermassive (\\Mbh$\\simeq$10$^{6}$--10$^{9}$\\Mo) black holes (SMBH's) in the centers of galaxies at z$\\simeq$0 (Kormendy \\& Richstone 1995; Magorrian, Tremaine, \\& Richstone 1998; Kormendy \\& Gebhardt 2001), it is important to understand how the SMBH's seen at z$\\simeq$0 have grown from lower mass BH's at z$\\simeq$20. A comprehensive review of SMBH's is given by Ferrarese \\& Ford (2004). \\ One sugges-\\vfill\\eject  \\n tion is that they ``grow'' through repeated mergers of galaxies which contain less massive BH's, so the byproduct is a larger single galaxy with a more massive BH in its center. The growth of this BH may then be observed via its AGN activity. If this scenario is valid, there may be an observable link between galaxy mergers and increased AGN activity (Silk \\& Rees 1998). Therefore, studying this link as a function of redshift could give insight into the growth of SMBH's and its relation to the process of galaxy assembly.  Recent numerical simulations addressed some long-standing issues in the dissipational collapse scenario by including previously-neglected energetic feedback from central SMBH's during the merging events (\\eg, Robertson \\etal 2005). They emphasize the relationship between the central BH mass and the stellar velocity dispersion, which confirms the link between the growth of BH's and their host galaxies (di Matteo \\etal 2005; Springel \\etal 2005ab). The present study provides observational support for these models at cosmological redshifts. ",
        "conclusions": "The fact that about 6\\% of all field galaxies are seen in the tadpole stage is an important constraint to hierarchical simulations. Springel \\etal (2005ab) predict a tadpole-like stage $\\sim$0.7--1.5 Gyr after a major merger begins, suggesting that the tadpole morphology represents an early-merger stage of two galaxies of roughly comparable mass. If this 6\\% indicates the fraction of time that an average galaxy in the HUDF spends in an early-merger stage during its lifetime, then every galaxy may be seen  temporarily in a tadpole stage for $\\sim$0.8 Gyr of its lifetime,  and may have undergone $\\sim$10-30 mergers during its lifetime (Straughn \\etal 2006).  More complex mergers involving multiple components may lead to irregular/peculiar and train-wreck type objects, and the luminous diffuse objects or clump-clusters, which dominate the galaxy counts at faint magnitudes (Driver \\etal 1998). Given that tadpoles only trace a certain type and stage of merging galaxies, the above statistics are a lower limit on the number of all mergers.  The question arises if tadpole galaxies and objects with point-sources that show signs of variability are drawn from the same population. Among our 165 tadpole galaxies, none coincide with the sample of 45 variable objects or with the CDF-S X-ray sources S (Alexander \\etal 2005). At most one or two of the variable candidates resemble the tadpole galaxies of Straughn \\etal (2006).  A factor of three of all AGN may have been missed, since their UV--optical flux was obscured by a dust-torus. In the AGN unification picture, AGN cones are two-sided and their axes are randomly distributed in the sky, so that an average cone opening-angle of $\\omega$ implies that a fraction 1--sin($\\omega$) of all AGN will point in our direction. If $\\omega$$\\simeq$45$^\\circ$ (\\eg, Barthel 1989), then every optically detected AGN (QSO) represents 3--4 other bulge-dominated galaxies, whose AGN reflection cone didn't shine in our direction. Hence, their AGN may remain obscured by the dust-torus. Such objects could be visible to Chandra in X-rays or to Spitzer in the mid-IR, although the available Chandra and Spitzer data are not deep enough to detect all HUDF objects to AB$\\simeq$28 mag.  Together with the factor of \\cge 2 incompleteness in the HUDF variability sample due to the limited time-baseline sampled thus far, the actual fraction of weak AGN present in these dynamically young galaxies may thus be a factor of $\\gtrsim$6--8$\\times$ larger than the 1\\% variable AGN fraction that we found in the HUDF. Hence, perhaps as many as $\\gtrsim$6--8\\% of all field galaxies may host weak AGN, only $\\sim$1\\% of which we found here, and another \\cge 1\\% could have been found if longer time-baseline had been available. Another factor of 3--4 of AGN are likely missing because they are optically obscured, The next generation of X-ray and IR telescopes (Windhorst \\etal 2006ab) and longer optical time-baselines are needed to detect all weak AGN in the HUDF.  Recent state-of-the-art hydrodynamical models (di Matteo \\etal 2005; Springel \\etal 2005ab; Hopkins \\etal 2005) suggest that during (major) mergers, the BH accretion rate peaks considerably {\\it after} the merger started, and {\\it after} the star-formation rate (SFR) has peaked. Their models suggest that, for massive galaxies, a tadpole stage is seen typically about 0.7 Gyr after the merger started, but $\\sim$1 Gyr before the SMBH accretes most of its mass, which is when the galaxy displays strong visible AGN activity. Since the lifetimes of QSO's and radio-galaxies are known to be \\cle (few$\\times$10$^7$)--10$^8$ years (Martini 2004; Grazian \\etal 2004, Jakobsen \\etal 2003), these models thus imply that the AGN stage is expected to occur considerably (\\ie, \\cge 1--1.5 Gyr) \\emph{after} the early-merger event during which the galaxy is seen in the tadpole stage.  The observed lack of overlap between the HUDF tadpole sample and the weak variable AGN sample thus provides observational support for this prediction. Hopkins \\etal (2005) have quantified the timescales that quasars will be visible during merging events, noting that for a large fraction of the accretion time, the quasar is heavily obscured. In particular, their simulations show that during an early merging phase --- our observed tadpole phase --- the intrinsic quasar luminosity peaks, but is completely optically obscured. Only after feedback from the central quasar clears out the gas, will the object become visible as an AGN. This should be observable by Spitzer in the mid-IR as a correspondingly larger fraction of IR-selected obscured faint QSO's. To study the relation between galaxy assembly and SMBH growth in detail, we need deeper surveys at longer wavelengths with the James Webb Space Telescope (JWST; Windhorst \\etal 2006a). The JWST photometric {\\it and} PSF stability are crucial for this, since many of our HUDF objects show significant variability of less than a few percent in total flux.  This research was partially funded by NASA grants GO-9793.08-A and AR-10298.01-A, NASA JWST grant NAG5-12460, the NASA Space Grant program at ASU, and the Harriet G. Jenkins Predoctoral Fellowship Program."
    },
    "0601/astro-ph0601458_arXiv.txt": {
        "abstract": "We present a theoretical study of double compact objects as potential short/hard gamma-ray burst (GRB) progenitors. An updated population synthesis code {\\tt StarTrack} is used to calculate properties of double neutron stars and black-hole neutron star binaries. We obtain their formation rates, estimate merger times and finally predict their most likely merger locations and afterglow properties for different types of host galaxies. Our results serve for a direct comparison with the recent {\\em HETE-II} and {\\em SWIFT} observations of several short bursts, for which afterglows and host galaxies were detected. We also discuss the possible constraints these observations put on the evolutionary models of double compact object formation. We emphasize that our double compact object models can successfully reproduce at the same time short GRBs within both young, star-forming galaxies (e.g., GRB~050709 and GRB~051221A), as well as within old, elliptical hosts (e.g.,  GRB~050724 and probably GRB~050509B). ",
        "introduction": "The recent detections of afterglows for short/hard GRBs have made possible a breakthrough in the study of this class of bursts. {\\em Swift} (Gehrels et al. 2005; Barthelmy et al. 2005) and {\\em HETE-II} (Villasenor et al. 2005) observations led to precise localizations of several short GRBs, and subsequent follow-up optical observations allowed for tentative connections with their host galaxies and a measurement of their redshifts.  GRB~050509B was found to lie in the vicinity of a large elliptical galaxy, with no current star formation and at a redshift of 0.225 (Gehrels et al. 2005).  GRB~050709 was found in the outskirts of a dwarf irregular galaxy with ongoing star formation at redshift 0.1606 (Hjorth et al. 2005; Fox et al. 2005; Covino et al. 2006), as was GRB~051221A (Soderberg and Berger 2005, Berger and Soderberg 2005) at redshift 0.5465. GRB~050724 was found within a small elliptical galaxy with no current star formation at a redshift of 0.258 (Berger et al. 2005a).  GRB~050813 was found close to a cluster of galaxies at redshift 1.7-1.9 (Berger 2005b), but no host galaxy has been proposed.  Compact object mergers from double neutron star (NS-NS) and black hole neutron star (BH-NS) binaries have been proposed for the first time by Paczynski (1986) and then further discussed as the central engines of short GRBs by number of authors (e.g., Eichler et al. 1989; Narayan, Paczynski \\& Piran 1992). Observational constraints on these populations may be obtained only for NS-NS systems since only such binaries are currently observed as binary pulsars. Merger rates derived from the observed sample of a handful of Galactic relativistic NS-NS binaries have been presented by Kalogera et al.\\ (2004 and references therein). Double compact objects with both neutron stars and black holes can be studied via population synthesis methods (e.g., Lipunov, Postnov \\& Prokhorov 1997; Portegies Zwart \\& Yungelson 1998; Belczynski, Kalogera \\& Bulik 2002c) There were several early population synthesis studies in the context of potential GRB progenitors (e.g., Bloom, Sigurdsson \\& Pols 1999; Belczynski \\& Bulik 1999; Fryer, Woosley \\& Hartmann 1999). In particular, it was found that NS-NS and BH-NS mergers are expected to take place outside host galaxies with long delay times.  However, it has only recently been recognized that the population of the double compact objects may be more diverse than was previously believed (Belczynski \\& Kalogera 2001; Belczynski, Bulik \\& Rudak 2002b; Perna \\& Belczynski 2002; Belczynski, Bulik \\& Kalogera 2002a). In addition to classical well-recognized channels, these newly recognized formation channels lead to the formation of tighter double compact objects, with short lifetimes and therefore possible prompt mergers within hosts.  The new formation scenarios were independently confirmed by detailed evolutionary calculations (Ivanova et al. 2003; Dewi \\& Pols 2003).  In this study we perform an updated analysis of double compact object mergers using population synthesis methods. Over the last several years and since our previous studies, the {\\tt StarTrack} population synthesis code has undergone major revisions and updates, many of which are guided and tested through comparisons with either observations or detailed evolutionary calculations (see Belczynski et al. 2005). Additionally, this new study is motivated by the recent short GRB observations and their likely connection with double compact object mergers. In \\S\\,2 we present the overview of our claculations. In \\S\\,3 the double compact object formation, merger rates, locations and their afterglow  properties are presented. Finally, in \\S\\,4 we discuss our results in context of the recent short-GRB observations. ",
        "conclusions": "We find two distinct populations of binary compact objects in terms of their binary-orbit characteristics and associated merger times due to differences in their evolutionary history. One is a classical population of rather wide, long-lived systems, with merger times of $\\sim 100-15,000$\\,Myr, and the other consists of tight, short-lived systems with merger times of $\\sim 0.001-0.2$\\,Myr. In a typical binary evolution model, there are roughly similar numbers of short- and long-lived NS-NS systems, while there are $\\sim$~20\\% and $\\sim$~80\\% of short- and long-lived BH-NS systems, respectively. The four known Galactic double neutron stars belong to the long-lived classical systems (see Fig.~\\ref{tmerger}), and, given the small number of NS-NS systems, we do not expect to see the short-lived systems, since they merge very soon after their formation.  We find that most of the double compact binaries ($\\gtrsim$~80\\%) that are formed in {\\em large} galaxies merge within them, independent of the galaxy type. For {\\em small} galaxies, we find that $\\sim$~20\\%, 50\\%, 70\\% mergers take place within elliptical, spiral and starburst hosts respectively (see Fig.~\\ref{locations}). The combination of binary compact object lifetimes and their merger locations leads to a testable prediction for the location of mergers in relation to the host galaxies. For starburst galaxies (which on average have small masses) where stellar populations are young, we expect mergers from short-lived double compact objects. These are expected to take place inside or in the vicinity of their hosts. It is worth noting that such locations, and more importantly, the mere existence of double compact object mergers associated with star-forming, young galaxies, is not predicted by any of the previous studies that considered only classical formation scenarios (long-delayed mergers). If it is shown that the star-forming host galaxy of GRB~050709 does not contain a dominant old underlying stellar population, this association will stand as the ``smoking gun'' evidence for the new short-lived double compact object formation channel, originally identified by Belczynski \\& Kalogera (2001) and Belczynski et al.\\ (2002a,c).  In elliptical galaxies, which have little or no ongoing star formation, we expect to find mergers of long-lived double compact objects at present. For massive hosts, most mergers should occur within the hosts, while for smaller galaxies outside the hosts. GRB~050509B was tentatively associated with a large elliptical galaxy, and its error circle indicates that this burst took place either within or in the outskirts of the host, in agreement with our findings.  GRB~050911 appears to be a case of a ``naked'' burst (Page et al.\\ 2005).  It could possibly be explained by our models with a double compact object binary originating from a small elliptical galaxy and merging far outside the host without any detectable afterglow. On the other hand GRB~050724 was found inside a small elliptical.  Our models are consistent with this observation, since $\\sim$~20\\% of mergers are predicted to occur within 5 kpc of the center for small ellipticals.  However, if more short GRBs are found at such intra-galactic locations, a number of new possibilities may be favored: BH-NS mergers with small or zero BH kicks, or in general small kicks for both NS and BH (e.g., Podsiadlowski et al.\\ 2004; Dewi et al.\\ 2005). Such alternative models would then have to be investigated.  We note that, at present, the necessity of imparting such small kicks to the majority of NS is an open question (c.f., Willems et al.\\ 2004; Chaurasia \\& Bailes 2005; Ihm et al.\\ 2005).  This, combined with our results for the locations of the mergers, indicates that short GRBs in the outskirts of galaxies need not be associated with globular clusters, as recently claimed by Grindlay et al.\\ (2006).  We find that short GRB progenitors originate most probably from a diverse population of compact objects which are formed in old but also in young stellar environments. Although the absolute formation rates and redshift distributions implied by a mixture of galaxy types must be investigated and compared with theoretical predictions in more detail (O'Shaughnessy et al., in preparation), at present we are able to account for the origin (and associated delay times) of several recently observed short GRBs.  Some of the recent work (e.g., Nakar, Gal-Yam \\& Fox 2005) stands in apparent contradiction with our findings, presenting claims that the current  models of double compact objects cannot explain the new short GRB observations. Also, long delay times ($\\sim 6$ Gyr) are inferred from the associations of short GRBs with old elliptical galaxies (Nakar et al.\\ 2005). We note that at present these should be considered with caution. First, they are based on the analysis of a very small number of bursts and may suffer from small number statistics; a significant fraction of short delay bursts cannot be ruled out with high statistical significance. The comparison in Nakar et al.\\ (2005) has been performed using some trial functions of delays, which do not necessarily correspond to the actual delays obtained in detailed population synthesis calculations. Second, their analysis neglects the information about the types of individual short GRB host galaxies and their particular star formation histories. Third, the models of Nakar et al.\\ (2005) cannot explain the observed Galactic population of relativistic double neutron stars, with merger times of $\\sim 100-3000$ Myr, that are consistent with our models (see Fig.~\\ref{tmerger}). Finally, the estimate of the delays may be complicated by evolutionary effects. We do know that the population of long GRBs is affected by cosmological evolutionary effects, for example, related with the metallicity evolution of the Universe, and so may be the population of short GRBs. Thus a direct comparison with the total star formation rate as performed by Nakar et al.\\ (2005) may be misleading. The influence of such evolutionary effects and of the contribution of different types of galaxies and their star formation history will be examined in detail in O'Shaughnessy et al.\\ (2006).  Additionally, the findings of Nakar et al.\\ (2005) should be weighted by the observations of GRB~050724 and GRB~051221A in young star forming environment, as well as the detection of GRB~050813 at z=1.7-1.9 where there ought be no short GRBs, if the delays are as long as they claim. If more short GRBs are found at high redshifts ($z \\gtrsim 1$), this will provide further support for delay times, that at least for some short GRBs, are shorter than $\\sim 6$ Gyr. It has also been suggested that there is an observational bias against detecting short GRBs at high redshifts with {\\em SWIFT}; the minimum flux for detection is apparently smaller in the observed sample than for GRBs with secure redshift determination, and therefore it favors redshift determinations for closest (low redshift) bursts (Hopman et al.\\ 2006). If this is the case, then it is clear that the observed redshift distribution leads to an overestimate in delay times in the analysis of Nakar et al.\\ (2005).  We conclude that the current observations point toward a short-hard GRB progenitor population that is diverse in terms of merger times and locations, as suggested by our models. Soon, with more observations of short GRBs, we will be able to critically test the population synthesis models. Specifically, we will investigate effects of low metallicity, binary populations dominated by roughly equal-mass binaries as suggested by Pinsonneault \\& Stanek (2006) and low natal kicks proposed by Pfahl et al.\\ (2002) on progenitors of double neutron star systems, and their relevance for both short GRB and gravitational-wave observations (Belczynski et al.\\ 2006, in preparation)."
    },
    "0601/astro-ph0601491_arXiv.txt": {
        "abstract": "The present paper reports on the {\\it RXTE} observations of the binary X-ray pulsar 4U~0115+63, covering an outburst in 1999 March-April with 44 pointings. The 3$-$30 keV PCA spectra and the 15$-$50 keV HEXTE spectra were analyzed jointly for the cyclotron resonance features. When the 3$-$50 keV luminosity at an assumed distance of 7 kpc was in the range (5$-$13)$\\times10^{37}$ erg s$^{-1}$, harmonic double cyclotron features were observed in absorption at $\\sim$11 and $\\sim$22 keV, as was measured previously during typical outbursts. As the luminosity decreased below $\\sim$5$\\times10^{37}$ erg s$^{-1}$, the second resonance disappeared, and the fundamental resonance energy gradually increased, up to $\\sim$16 keV at 0.16$\\times10^{37}$ erg s$^{-1}$. These results reconfirm the report by Mihara et al.\\ (2004) using {\\it Ginga}, who observed a single absorption at $\\sim$16 keV in a minor ($\\sim10^{37}$ erg s$^{-1}$) outburst of this object. The luminosity-dependent cyclotron resonance energy may be understood as a result of a decrease in the accretion column height, in response to a decrease in the mass accretion rate. ",
        "introduction": "Accreting binary pulsars are considered to have strong surface magnetic fields in the range of several times $10^{12}$ G. One of the methods to accurately measure their fields is to observe cyclotron resonance scattering features (CRSFs) in their X-ray spectra, because the resonance energy $E_a$ and the magnetic field strength $B$ are related with each other as $ E_a = 11.6\\  B_{12} (1+z_{g}) ^{-1}\\ [{\\rm keV}] $, where $B_{12}$ is the magnetic field strength in units of $10^{12}$ Gauss and $z_g$ is the gravitational redshift. So far, CRSFs have been detected from $\\sim$ 15 X-ray pulsars, all in absorption, with balloons (e.g., Tr\\\"umper et al.\\ 1978), {\\it HEAO-1} (Wheaton et al.\\ 1979; White et al.\\ 1983), {\\it Ginga} (e.g., Clark et al.\\ 1994; Mihara 1995; Makishima et al.\\ 1999), {\\it RXTE} (e.g., Coburn et al.\\ 2002; Heindl et al.\\ 2001), {\\it BeppoSAX} (e.g., Santangelo et al.\\ 1999; Orlandini et al.\\ 1998), and other missions.   The recurrent transient 4U~0115+63, with the 3.6 s pulsations (Rose et al.\\ 1979), is one of the X-ray pulsars whose CRSF has been studied in great detail. The optical companion is an O9e star, V635 Cassiopeae (Unger et al.\\ 1998), with the orbital period of 24.3 days. The distance to 4U~0115+63 is estimated as 7 kpc (Negueruela and Okazaki\\ 2001). Its CRSF was first discovered at $\\sim$23 keV in the {\\it HEAO}-1 A4 spectra by Wheaton et al.\\ (1979). Using the {\\it HEAO}-1 A2 data obtained in the same outburst, White et al.\\ (1983) suggested that the 23 keV feature is in fact the second harmonic resonance, with the fundamental resonance at $\\sim$11 keV. The suggested double harmonic structure was reconfirmed with {\\it Ginga} by Nagase et al.\\ (1991). Moreover, the third harmonic feature was found with {\\it RXTE} by Heindl et al.\\ (1999), and the fourth harmonic with {\\it BeppoSAX} by Santangelo et al.\\ (1999). In the {\\it BeppoSAX} data, the resonance energies were 12.7, 24.2, 35.7 and 49.5 keV, or nearly harmonic ratios of 1:1.9:2.8:3.9. Thus, 4U~0115+63 is one of the most suitable objects to study the physics of cyclotron resonance in the polar caps of binary X-ray pulsars.   This source has been observed at various X-ray luminosity levels, and the fundamental resonance energy has been measured repeatedly at $\\sim$11 keV. However, in a minor outburst in April 1991 observed with {\\it Ginga}, when the 2$-$60 keV luminosity (1.9$\\times10^{37}$ erg s$^{-1}$) was $\\sim$7 times lower than those typical at the outburst peaks, a drastic change in the CRSF was detected: instead of the familiar double absorption features at $\\sim$11 keV and $\\sim$22 keV, it exhibited a single deep and wide absorption at $\\sim$16 keV (Mihara 1995; Mihara et al.\\ 1998; Makishima et al.\\ 1999; Mihara et al.\\ 2004, hereafter Paper 1). Such a large change in $E_a$, presumably depending on the luminosity, had been observed neither from 4U~0115+63 itself previously, nor from other sources.   In Paper 1, we tentatively concluded that the fundamental resonance energy, at $\\sim$11 keV in normal outbursts, increased to $\\sim$16 keV, because a lower luminosity would make the accretion column shorter, and hence increase the magnetic field intensity at the column top. However, there has remained an alternative possibility that the second harmonic resonance, normally at $\\sim$22 keV, decreased to $\\sim$16 keV in the 1991 outburst. Furthermore, even if the former interpretation is correct, the observed change in $E_a$ was considerably larger than is predicted by an accretion column model by Burnard et al.\\ (1991). To make the results in Paper 1 less ambiguous, we need to more densely sample the spectra as the luminosity changes over a wide range.   In this paper, we analyzed the {\\it RXTE} data of 4U~0115+63 acquired in the 1999 March-April outburst, and studied the resonance energy as a continuous function of the luminosity in the rising and declining phases of the outburst. We have indeed detected the clear luminosity-dependent changes in the CRSF, and successfully reconfirmed the inference made in Paper 1. All of the errors appeared in this paper are 90 \\% confidence levels. ",
        "conclusions": "\\label{sec:4}  We have analyzed the 34 PCA$+$HEXTE datasets, covering the whole 1999 March-April outburst of 4U~0115+63 in which the $3-50$ keV source luminosity changed over $L_{\\rm x} = (0.17-14) \\times 10^{37}$ erg s$^{-1}$. When $L_{\\rm x} > 7 \\times 10^{37}$ erg s$^{-1}$, we observed the familiar double CRSFs with $E_{\\rm a1} \\simeq 11$ keV. As $L_{\\rm x}$ decreased across a rather narrow range of $( 5 \\pm 2) \\times 10^{37}$ erg s$^{-1}$, the second harmonic resonance disappeared, and the fundamental resonance energy increased from $E_{\\rm a1}\\sim 11$ keV to $\\sim 16$ keV. These results reconfirm Paper 1, and unambiguously reveal that the resonance energy {\\it increases} as the source gets less luminous.   \\subsection{The resonance energy variations}  Except for some hysteresis effects between the rising and decay phases of the outburst, $E_{\\rm a1}$ changes roughly as a single-valued function of $L_{\\rm x}$. In particular, the intensity-sorted 3$-$30 keV PCA spectra and the date-sorted PCA$+$HEXTE data imply consistent results. Furthermore, the {\\it Ginga} results reported in Paper 1 fall on the same $E_{\\rm a1}$ vs. $L_{\\rm x}$ relation. Therefore, the phenomenon is inferred to have a good reproducibility as a function of the luminosity.   The CRSF is known to depend also on the pulse phase, presumably because different scattering regions characterized by slightly different magnetic field strengths are observed as the neutron star rotates. When studying luminosity-related changes in the resonance energy, such pulse-phase dependent effects must be considered. In the case of 4U~0115+63, however, the phase dependent variation of $E_{\\rm a1}$ in the 1999 outburst was relative small, $\\sim10$\\% (Santangelo et al.\\ 1999). Furthermore, the phase-resolved analysis performed in Paper 1, on the two different luminosity levels, did not affect the main results from the phase-averaged spectroscopy. We therefore concentrate, in the present paper, on the phase-averaged analysis.   As mentioned in Paper 1, luminosity-dependent changes of the accretion column height provide the most likely explanation to our results, because the column is expected to become taller as the source luminosity increases (e.g., Burnard et al. 1991). Assuming that the CRSF is formed at a height $h_{\\rm r}$ above the neutron star surface in the accretion column, and the magnetic field strength there follows the dipole law, we then expect \\begin{equation} E_{\\rm a1} \\propto (R_{\\rm NS} + h_{\\rm r})^{-3} (1+z_{\\rm g})^{-1} \\label{eq3} \\end{equation} with $R_{\\rm NS}$ the radius of the neutron star. The factor $(1+z_{\\rm g})$ describes the gravitational redshift, but below we assume $z_{\\rm g} =0$ for simplicity. Then, the relative resonance height, $h_{\\rm r}/R_{\\rm NS}$, can be related to the resonance energy as \\begin{equation} \\frac{h_{\\rm r}}{R_{\\rm NS}} \\approx \\ \\left( \\frac{E_{\\rm a}}{E_{\\rm 0}} \\right)^{-1/3} - 1 ~, \\label{eq4} \\end{equation} where $E_{\\rm 0}$ is the resonance energy to be observed on the neutron star surface.   Substituting the observed value of $E_{\\rm a1}$ into Equation (\\ref{eq4}), we have calculated $h_{\\rm r}/R_{\\rm NS}$ as a function of the X-ray luminosity. The results are shown in Figure \\ref{fig9}. Since there is no {\\it a priori} knowing of $E_0$, we employed two different values of $E_0$; 18 keV and 20 keV, the former close to the observed maximum value of $E_{\\rm a1}$. Except for some differences between the two assumptions on $E_0$, the height $h_{\\rm r}$ thus increases in a rough proportion to the X-ray luminosity, up to $\\sim 7 \\times 10^{37}$ erg s$^{-1}$ where $h_{\\rm r}$ appears to saturate at $\\sim 0.2 R_{\\rm NS}$. In addition,  $h_{\\rm r}/R_{\\rm NS}$ may approach a finite value of $0.03\\sim0.07$, instead of zero, as $L_{\\rm x}$ decreases below $\\sim 2 \\times 10^{37}$ erg s$^{-1}$.   The results presented in Figure \\ref{fig9} may be compared to a theoretical prediction by Burnard et al.\\ (1991), who calculated the height of the accretion column $h_{\\rm top}$ as \\begin{equation} \\frac{h_{\\rm top}}{R_{\\rm NS}} \\approx \\frac{L_{\\rm x}}{L_{\\rm Edd}^{\\rm eff} \\ H_\\perp} ~~ . \\label{eq5} \\end{equation} Here, $L_{\\rm Edd}^{\\rm eff}$ is the Eddington luminosity along the magnetic field which is identical to the conventional Eddington Luminosity for a $1.4 M_\\odot$ neutron star, $L_{\\rm Edd}^{\\rm eff} =  2.0 \\times 10^{38}$ ergs s$^{-1}$, and $H_\\perp$ is the ratio of the Thomson cross section to the Rosseland-averaged electron scattering cross section for radiation flows across the magnetic fields. The dashed line in Figure \\ref{fig9} shows this prediction, assuming $H_\\perp=1.23$  (Paper 1) and $R_{\\rm NS}=10$ km. Except for the observed saturation toward the highest luminosity and $ \\lesssim 2 \\times 10^{37}$ erg s$^{-1}$, the value of $h_{\\rm r}$ measured at each luminosity level thus corresponds to $\\sim 70\\%$ of the predicted $h_{\\rm top}$. This is quite reasonable, because the resonance energy is expected to sample the magnetic field strength which is measured at, or slightly below, the top of the column (i.e., $h_{\\rm r} \\lesssim h_{\\rm top}$).   So far, we have assumed that the observed X-ray intensity changes are caused solely by actual variations of the intrinsic source luminosity. However, the intensity may also be affected by other extrinsic factors, such as partial obscuration by materials around the accretion disk. Such effects are known, e.g., in Her X-1 (Mihara et al.\\ 1991). In the present outburst of 4U~0115+63, Heindl et al.\\ (1999) detected occasional appearance of ``mHz QPO'', with a period of $\\sim500$ sec and intensity changes up to $\\sim40$ \\%, and suggested that the phenomenon is possibly due to source obscuration by some ionized materials. If so, we would need a caution in interpreting the results from our intensity-sorted analysis.   In order to investigate effects of the mHz QPO on our results, we inspected PCA light curves acquired over the period used for our intensity-sorted analysis (March 27 through April 8). As shown in the inset to Figure \\ref{fig10},  a relatively clear QPO was found on March 31, although the period is about 1000 sec instead of the 500 sec reported by Heindl et al.\\ (1999). Then, we sliced the light curves as shown there, and produced a pair of PCA spectra corresponding to peaks and valleys of the QPO. Figure \\ref{fig10} presents these spectra and their ratio, as well as residuals when fitted individually with the NPEX model. The resonance energy thus clearly increases at the QPO valley. By fitting the spectra with the NPEX times CYAB2 model (with $W_2$ again fixed at $W_1$), we have constrained the resonance centroid as $10.4_{-0.7}^{+0.2}$ keV at the QPO peak with $L_{\\rm x} = 8.2 \\times 10^{37}$ erg s$^{-1}$, and $11.8_{-0.5}^{+0.9}$ keV at the QPO bottom with $L_{\\rm x} = 5.2 \\times 10^{37}$ erg s$^{-1}$. Thus, the energy shift is statistically significant. Furthermore, these two data points line up in Figure \\ref{fig8} closely on the general $L_{\\rm x}$ vs. $E_{\\rm a1}$ correlation from the day-sorted and intensity-sorted analysis. We hence conclude that our basic results remain unaffected, and that the QPO, at least on this occasion, is likely to reflect real changes in the intrinsic luminosity rather than some obscuration effects. This inference is reinforced by the relatively flat spectral ratio, because the ratio should show an increased low-energy absorption if the  obscuration were due to neutral material, or Fe-K edge feature if it were highly ionized.     \\subsection{Behavior of the other parameters}  So far, many authors (e.g. Paper 1 and Coburn et al.\\ 2002) studied relations among the continuum and cyclotron line parameters. Our work provides a valuable opportunity to investigate how these parameters in a single system change in correlated ways, when the resonance energy varies. Among various correlations, a particularly interesting one is that between $\\tau$ and the $W/E_{\\rm a}$ ratio; in fact, Coburn et al.\\ (2002) found a positive correlation between them over a large sample of phase-averaged spectra of X-ray pulsars. Later, Kreykenbohm et al.\\ (2004) noticed that the same correlation holds for the phase-resolved spectra of GX~301-2, and argued that the correlation may be explained if the accretion column has a tall cylindrical shape rather than flat coin-shaped geometry. As presented in Figure \\ref{fig11}, the fundamental and second harmonic widths of 4U~0115+63 from the present observations both depend positively on their respective resonance depths, in agreement with the correlation found by Coburn et al.\\ (2002). These results, together with the argument by Kreykenbohm et al.\\ (2004), strengthen the tall cylindrical column geometry which we invoked in \\S~4.1. For reference, we did not find particular correlations between $E_{\\rm a}$ and $kT$, or between $E_{\\rm a}$ and $W$.   Figure~\\ref{fig12} shows the spectral parameters as a function of the X-ray luminosity. There, the positive $\\tau$ vs. $W/E_{\\rm a}$ correlation of Figure~\\ref{fig11} has been decomposed into luminosity dependent changes in $W/E_{\\rm a}$ (panel a) and  $\\tau$ (panels c, d). Up to $\\sim 4 \\times 10^{37}$ erg s$^{-1}$, $\\tau_1$ thus stays relatively constant, with a hint of mild increase presumably due to the increased column density in the emission region. Over the same luminosity range, the fractional width $W_1/E_{\\rm a1}$ increases more clearly. This would not be due to changes in the Doppler broadening (e.g. M\\'{e}sz\\'{a}ros 1992), since the NPEX $kT$ parameter shown in Figure~\\ref{fig12}b, which is thought to approximate the electron temperature in the emission region (Mihara et al.\\ 2004; Makishima et al.\\ 1999), stays rather constant. Instead, the increase in the $W_1/E_{\\rm a1}$ ratio could be a result of the luminosity-correlated elongation in the accretion column, which would cause a larger range of magnetic field intensities to participate in the resonance formation.   Beyond $\\sim 4 \\times 10^{37}$ erg s$^{-1}$, both $\\tau_1$ and  $W_1/E_{\\rm a1}$ start decreasing clearly, while the 2nd harmonic resonance develops rapidly both in depth ($\\tau_2$, Figure \\ref{fig12}d) and relative width ($W_2/E_{\\rm a2}$ in Figure \\ref{fig12}a). In short, the CRSF makes a transition from the single feature at low luminosities to the harmonic double feature at higher luminosities. Since this occurs approximately over the luminosity range where the rapid change in $E_{\\rm a1}$ takes place, we consider that the single-to-double transition of the CRSF has the same origin as the resonance energy shift.   The behavior of the fundamental and 2nd harmonic parameters at the single-to-double transition may reflect basic differences in their elementary processes. The fundamental resonance has a very large cross section, but it acts as scattering rather than absorption, because an electron which is excited by absorbing a photon of energy $\\sim E_{\\rm a1}$ will immediately return to the ground state by emitting a photon of nearly the same energy. The second harmonic resonance, though with a much smaller cross section, will in contrast act as pure absorption, because the electron excitation/deexcitation in this case occurs via absorption of a photon of energy $\\sim 2 E_{\\rm a1}$ and cascade emission of two photons of energies $\\sim E_{\\rm a1}$ each. The emitted photons will fill up the fundamental resonance, making it shallower (so-called two-photon effects; Alexander \\& Meszaros 1991). In the present case, the accretion column may be effectively thick to the fundamental resonance ($\\tau_1 > 1$) essentially at all the observed luminosities, but presumably optically thin ($\\tau_2  \\ll 1$) to the second resonance photons when the luminosity is low and hence the accretion column is short. As the source luminosity increases, the column becomes taller and opaque to the second resonance, leading to the emergence of the second resonance feature. At the same time, the two-photon effect would reduce $\\tau_1$, just as seen in Figure \\ref{fig12}c. A similar effect was observed in the {\\it INTEGRAL} spectrum of GX~$301-2$ by Okada et al.\\ (2004), who reported that $\\tau_1$ of its $\\sim 35$ keV CRSF decreased toward higher luminosities.     \\subsection{The case of X0331+53}  Although 4U~0115+63 thus provides the first clear example of luminosity-dependent changes of the cyclotron resonance energy, there has emerged another promising case. This is the transient X-ray pulsar X0331+53 (V0332+53), from which a very prominent CRSF was detected with {\\it Ginga} at $E_{\\rm a} = 28.5 \\pm 0.5$ keV together with a hint of the second harmonic (Makishima et al.\\ 1990). This was observed in an outburst when the $2-60$ keV luminosity was $\\sim 2 \\times 10^{37}$ erg s$^{-1}$ at an assumed distance of 3 kpc. A re-analysis of the same data using the NPEX continuum have revised the value slightly to $E_{\\rm a} = 27.2 \\pm 0.3$ keV (Makishima et al.\\ 1999).   The object entered a bright outburst from 2004 November, becoming considerably more luminous than was observed with {\\it Ginga} (Swank et al.\\ 2004; Remillard 2004). Observations with {\\it INTEGRAL} (Kreykenbohm et al.\\ 2005) and {\\it RXTE} (Coburn et al.\\ 2004) have clearly reconfirmed the fundamental CRSF, and further revealed the second and third resonances. The fundamental resonance energy obtained with {\\it INTEGRAL} is $24.9 \\pm 0.1$ keV (Kreykenbohm et al. 2005) at a $3-50$ keV luminosity of $\\sim8\\times10^{37}$ erg s$^{-1}$. This means that a factor 4 increase in the luminosity (from the {\\it Ginga} outburst to the present one) is accompanied by a $\\sim 10\\%$ decrease in $E_{\\rm a1}$. As mentioned by Mihara et al.\\ (2004), the same effect was already visible between two pointings with {\\it Ginga}. For comparison, the resonance energy of 4U~0115+63 changes by a somewhat larger amount ($\\sim 20\\%$) across the same factor 4 luminosity range (Figure \\ref{fig8}). The archival {\\it RXTE} data of X0331+53 are currently being analyzed, and the results will be reported elsewhere (Nakajima et al.\\ , in preparation).   Now, we have two examples of luminosity dependent changes in the resonance energy. However, there is an apparent counter example, namely Her X-1. In a large amount of {\\it RXTE} data covering its ``main on'' phase, the $2-30$ keV luminosity (at 5.8 kpc) of Her X-1 varied by a factor of 2 over  $(1.8-3.2) \\times 10^{37}$ erg $s^{-1}$, but the CRSF stayed at $\\sim 40$ keV (Gruber et al.\\ 2001). One possible interpretation of these results on Her X-1 is that its luminosity swing was not large enough to reach the critical range ({\\it cf.} Figure~\\ref{fig8}) where $E_{\\rm a}$ starts changing significantly. Alternatively, the luminosity-related changes in the CRSF may depend on the object, for some reasons which are yet to be detailed.   In conclusion, our study using the {\\it RXTE} data has confirmed the inference made in Paper 1, that the cyclotron resonance energy of 4U~0115+63 increases as the X-ray luminosity decreases. While this provide a new tool with which we can diagnose the accretion column of strongly magnetized neutron stars, it remains yet to be confirmed whether the phenomenon is common among this type of objects."
    },
    "0601/astro-ph0601172_arXiv.txt": {
        "abstract": "{The role of kink instability in magnetically driven jets is explored through numerical one-dimensional steady relativistic MHD calculations. The instability is shown to have enough time to grow and influence the dynamics of Poynting-flux dominated jets. In the case of AGN jets, the flow becomes kinetic flux dominated at distances $\\simmore 1000 r_{\\rm{g}}$ because of the rapid dissipation of Poynting flux. When applied to GRB outflows, the model predicts more gradual Poynting dissipation and moderately magnetized flow at distances of $\\sim10^{16}$ cm where the deceleration of the ejecta due to interaction with the external medium is expected. The energy released by the instability can power the compact ``blazar zone'' emission and the prompt emission of GRB outflows with high radiative efficiencies. ",
        "introduction": "\\label{intro}  Relativistic collimated outflows have been extensively observed in active galactic nuclei (AGN) and X-ray binaries (XRB). Gamma-ray bursts (GRB) are also believed to be connected to ultrarelativistic and collimated outflows to overcome the ``compactness problem'' (e.g. Piran 1999) and to explain the achromatic afterglow breaks (Rhoads 1997; Sari et al. 1999). It is also believed that all these sources are powered by accretion of matter by a compact object.  The widely accepted mechanism for jet acceleration and collimation in the context of AGN and XRB jets is that of magnetic driving. According to this paradigm, magnetic fields anchored to a rotating object can launch an outflow. The rotating object can be a star (Weber \\& Davis 1967; Mestel 1968), a pulsar (Michel 1969; Goldreich \\& Julian 1970), an accretion disk (Bisnovatyi-Kogan \\& Ruzmaikin 1976; Blandford 1976; Lovelace 1976; Blandford \\& Payne 1982) or a rotating black hole (Blandford \\& Znajek 1977). The material is accelerated thermally up to the sonic point and centrifugally until the Alfv\\'en point, defined as the point where the flow speed equals the Alfv\\'en speed. After the Alfv\\'en point the inertia of mater does not allow corotation of the magnetic field. As a result, the magnetic field lines bend, developing a strong toroidal component.  Further out the flow passes through the fast magnetosonic point where most of the energy of the flow remains in the form of Poynting flux in the case of relativistic outflows (Michel 1969; Goldreich \\& Julian 1970; Sakurai 1985; Beskin et al. 1998). Further acceleration of the flow is not straightforward within ideal MHD. It can be shown, for example, that a radial flow is not accelerated after the fast point (e.g. Beskin 1998). This is a result of the fact that the magnetic pressure and tension terms of the Lorentz force almost cancel each other (Begelman \\& Li 1994).   A limited degree of acceleration of the flow  is possible if it has a decollimating shape (i.e. the magnetic field diverges faster than radial; Li et al. 1992; Begelman \\& Li 1994).  Magnetized jets suffer from a number of instabilities.  Interaction with the environment causes instabililty of the Kelvin-Helmholtz type and kink instability causes internal rearrangement of the field configuration. Here we focus   on kink instability, since it internally dissipates magnetic energy associated with the Poynting flux. As demonstrated elsewhere (Drenkhahn 2002; Drenkhahn \\& Spruit 2002; Spruit \\& Drenkhahn 2003) such internal energy dissipation directly leads to acceleration of the flow. Dissipation steepens the radial decrease of magnetic pressure, thereby lifting the cancellation between outward pressure force and inward magnetic tension, and allowing the magnetic pressure gradient to accelerate the flow.  \\subsection{``AC'' versus ``DC'' outflows}  While dissipation of magnetic energy can thus happen through kink instability in an initially axisymmetric (``DC'') flow, it can also happen more directly by reconnection in the outflow generated by a non-axisymmetric rotator (``AC'' flow). The two cases behave differently in terms of the acceleration profile, and the location and amount of radiation produced by the dissipation process.  A nonaxisymmetric rotator produces a ``striped'' outflow (as in the case of a pulsar wind) with reconnectable changes of direction of the field embedded in the flow, and energy release  independent of the opening angle of the jet. In the DC case, where energy release is instead mediated by an instability, the rate of energy release is limited by the time it takes an Alfv\\'en wave to travel across the width of the jet. This makes it a sensitive function of the jet opening angle.  The ``AC'' case has been studied in detail by Drenkhahn (2002), Drenkhahn and Spruit (2002), with application to Gamma-ray bursts. In the case of AGN and XRB, on the other hand, the collimated jet is arguably best understood if the field in the inner disk is of uniform polarity, resulting in an initially axisymmetric flow. Another difference is the lower bulk Lorentz factors in the AGN/XRB case, resulting in faster energy release (in units of the dynamical time of the central engine).  The purpose of this paper is to explore the consequences of  magnetic dissipation by internal instability in such axisymmetric (or DC) cases, and its observational signatures. We also apply the calculations to the GRB case, where we compare the results with the AC case studied  before.  \\subsection{Energy release and field decay by the instability}  We limit ourselves to a flow with constant opening angle. That is, we leave aside the collimation process. Kink instability is modeled by adding a sink term to the induction equation to account for the non-ideal MHD effects arising from it.  Linear stability theory of kink instability yields a growth time of the order $t_{\\rm k}=r\\theta /v_{\\rm A,\\phi}$ where $r$ is the radius of the jet, $\\theta$ its opening angle and $v_{\\rm A,\\phi}$ the Alfv\\'en speed based on the azimuthal component $B_\\phi$ of the field. This is independent of the poloidal field component, at least for the so-called internal kink modes (which do not disturb the outer boundary of the field, cf.\\  Bateman 1978 for details). For stability analysis and numerical simulations with astrophysical applications see, e.g. Begelman (1998); Appl et al. (2000); Lery et al. (2000); Ouyed et al. (2003); Nakamura \\& Meier (2004). Based on linear theory, we would predict that the poloidal field component can be ignored for the rate of energy release.  It is not clear, however, that the poloidal component can be ignored for the nonlinear development of the instability, which is what actually determines the energy release. As a way to explore the effect of possible nonlinear stabilization by a poloidal component we compare two cases in the calculations: one with energy release and field decay given by the Alfv\\'en time across the jet ( $t_{\\rm k}$ above), and one in which this rate is assumed to be reduced by the poloidal component. This mainly affects the early phases of the acceleration of the flow beyond the light cylinder.  We find that kink instability has time to grow  in the AGN and XRB cases,  dissipating energy in the toroidal component of the magnetic field  while accelerating the flow at the same time. The dissipation of magnetic energy is almost complete and fast in the case of AGN jets,  so that on parsec scales the flow has become kinetic energy dominated, in agreement with current interpretations of the observations (e.g. Sikora et al. 2005, where the possible effects of magnetic dissipation are also discussed briefly).  The DC model with kink instability also produces significant flow acceleration in the GRB case, but conversion of the Poynting flux is less effective than the AC model in this case.  The structure of the paper is as follows. In Sec.~2 we discuss MHD instabilities in jets and focus on the kink instability and its growth rate. The model is described in Sec.~3 including the assumptions, the dynamical equations and the parameters at the base of the flow. In Sec.~4, we apply the model to the case of AGN jets and GRBs, while the last two Sections present the discussion and conclusions. ",
        "conclusions": "The standard scenario for jet launching, acceleration  and collimation involves large scale magnetic fields anchored to a rotating object (e.g. Blandford \\& Payne 1982; Sakurai 1985). The flow passes through three critical points, i.e. the slow, the Alfv\\'en and the fast point. At the fast point the ratio of Poynting to matter energy flux is much larger than unity in the case of relativistic jets (Michel 1967; Camenzind 1986; Beskin et al. 1998) while further acceleration of the flow appears hard to achieve within ideal MHD except if the flow is decollimated (Li et al. 1992; Begelman \\& Li 1994).  In this work, we study how this picture is modified when one takes into account the fastest growing current driven instability, i.e. the $m=1$ mode kink instability. We have modeled the instability by modifying the induction equation to account for non-ideal MHD processes and solving the relativistic MHD equations in the case of a radial flow. The instability is driven by $B_\\phi$, dissipates Poynting flux and has been shown to be an efficient mechanism to accelerate the flow.  The key parameter of the model is the ratio $\\sigma_0$ of the Poynting to matter energy flux at the base of the flow. A large part of the AGN jet phenomenology can be understood in the context of this model for $\\sigma_0\\sim$ several. On sub-pc scales the flow is Poynting-flux dominated with $B_\\phi\\gg B_r$. The flow is shown to be accelerated fast and to become matter dominated already at $\\sim$pc scales, while it reaches terminal bulk factors of a few tens. The emission at the blazar zone can be a result of either internal shocks that take place in an unsteady flow, where fast shells catch up with slower ones, converting a small fraction of the bulk kinetic energy of the flow into radiation (Rees \\& M\\'esz\\'aros 1994; Spada et al. 2001), or direct manifestation of the energy released by the instability.  Within the same model, we propose that GRBs are a result of more Poynting flux dominated outflows with $\\sigma_0\\sim$100. For these values of $\\sigma_0$ the flow reaches terminal bulk Lorentz factors of the order of a few to several hundreds, while it remains moderately magnetized (i.e. $\\sigma_\\infty\\sim 1$) at the afterglow region region. Although there is evidence for magnetized ejecta from afterglow modeling (e.g. Kumar \\& Panaitescu 2003; Zhang \\& Kobayashi 2005), more results are anticipated from early afterglow follow-ups that can test the model.  In the internal shock scenario for the prompt GRB emission, the shells collide at typical distances of $10^{13}-10^{15}$ cm, where the flow is moderately Poynting-flux dominated.  On the other hand, internal shock and Poynting-flux models exclude each other somewhat. If a strong magnetic field is added to an internally-shocked outflow, the radiative efficiency is further reduced with respect to that expected from the collision of unmagnetized shells (e.g. Fan et al. 2004). At the same time, dissipation in a predominantly magnetic outflow by instability (DC model) or internal reconnection (AC model) can produce radiation naturally at very high efficiency (up to 50\\%)."
    },
    "0601/astro-ph0601344_arXiv.txt": {
        "abstract": "The kinematical properties of elliptical galaxies formed during the mergers of equal mass, stars+gas+dark matter spiral galaxies are compared to the observed low velocity dispersions found for planetary nebulae on the outskirts of ellipticals, which have been interpreted as pointing to a lack of dark matter in ellipticals (which poses a problem for the standard model of galaxy formation). We find that the velocity dispersion profiles of the stars in the simulated ellipticals match well the observed ones. The low outer stellar velocity dispersions are mainly caused by the radial orbits of the outermost stars, which, for a given binding energy must have low angular momentum to reach their large radial distances, usually driven out along tidal tails. ",
        "introduction": "There is a wide consensus that spiral galaxies must be embedded within dark matter halos, as there have been no other good explanations of the observed flat rotation curves of spiral galaxies, unless one resorts to modifying physics (e.g. MOND, see McGaugh in these proceedings). Moreover, dissipationless cosmological $N$-body simulations lead to structures, whose halos represent most spiral galaxies (Hayashi et el. 2005; Stoehr 2005).  If elliptical galaxies originate from major mergers of spiral galaxies (Toomre 1977; Mamon 1992; Baugh, Cole \\& Frenk 1995; Springel et al. 2001a), then they too should possess dark matter halos. Using planetary nebulae (PNe) as tracers of the dark matter at large radii, Romanowsky et al. (2003) found low velocity dispersions for their outermost PNe, which after some simple Jeans modeling and more sophisticated orbit modeling led them to conclude to a dearth of dark matter in ordinary elliptical galaxies. This result is not expected in the standard model of structure and galaxy formation. This has led us (Dekel et al. 2005) to analyze the final outputs of $N$-body simulations of spiral galaxies merging into ellipticals. ",
        "conclusions": ""
    },
    "0601/astro-ph0601469.txt": {
        "abstract": "\\baselineskip = 11pt \\leftskip = 0.65in \\rightskip = 0.65in \\parindent=1pc   {\\small  The study of planets beyond the solar system and the search for other habitable planets and life is just beginning. Ground-based (radial velocity and transits) and space-based surveys (transits and astrometry) will identify planets spanning a wide range of size and orbital location, from Earth-sized objects within 1 AU to giant planets beyond 5 AU, orbiting stars as near as a few parsec and as far as a kiloparsec. After this initial reconnaissance, the next generation of space observatories will directly detect photons from planets in the habitable zones of nearby stars. The synergistic combination of measurements of mass from astrometry and radial velocity, of radius and composition from transits, and the wealth of information from the direct detection of visible and mid-IR photons will create a rich field of comparative planetology. Information on proto-planetary and debris disks will complete our understanding of the evolution of habitable environments from the earliest stages of planet-formation through to the transport into the inner solar system of the volatiles necessary for life.  The suite of missions necessary to carry out the search for nearby, habitable planets and life requires a ``Great Observatories'' program for planet finding (SIM PlanetQuest, Terrestrial Planet Finder-Coronagraph, and Terrestrial Planet Finder-Interferometer/Darwin), analogous to the highly successful ``Great Observatories Program'' for astrophysics. With these new Great Observatories, plus the James Webb Space Telescope, we will extend planetology far beyond the solar system, and possibly even begin the new field of comparative evolutionary biology with the discovery of life itself in different astronomical settings. \\\\~\\\\~\\\\~} ",
        "introduction": "}  The search for habitable planets and life beyond Earth represents one of the oldest questions in natural philosophy, but one the youngest fields in astronomy. This new area of research derives its support among the scientific community and the general public from the fact that we are using 21$^{st}$ Century technology to address questions that were first raised by inquiring minds almost 2,500 years ago. Starting in 1995, radial velocity and, most recently, transit and microlensing studies have added more than 168 planets to the nine previously known in our own solar system. With steadily improving instrumentation, the mass limit continues to drop while the semi-major axis limit continues to grow: a 7.5 M$_\\oplus$ planet orbiting the M star GL 876 at 0.02 AU (Rivera \\etal 2005)  and a 4 M$_{Jup}$ planet orbiting 55 Cnc at 5.2 AU (Marcy \\etal 2002) define boundaries which are sure to be eclipsed by newer discoveries. While we do not yet have the tools to find true solar system analogues or an Earth in the habitable zone of its parent star, evidence continues to accumulate (Marcy \\etal 2005) that the number of potential Earths is large and that some could be detected nearby, if only we had the tools. In this article on space missions, we assess the prospects for the discovery and eventual characterization of planets of all sizes --- from gas-giants to habitable terrestrial analogues. Considerations of length necessarily make this discussion incomplete and we have omitted discussion of valuable techniques which do not by their nature, e.g. microlensing, lend themselves to follow-up observations relevant to physical characterization of planets and the search for life.  \\begin{table*}[tbp] \\tabletypesize{\\scriptsize} \\caption{Space Projects Presently Under Consideration \\label{projtable}} \\begin{tabular}{lccc} &Approx. Sensitivity &Planet & Approximate \\\\ & Planet Size, Orbit, Dist. &Yield$^1$ & Launch Date\\\\ \\hline \\multicolumn{4}{l}{\\it Survey for Distant Planets (Transits)} \\\\ COROT& \t$2 R_\\oplus$ at 0.05 AU (500 pc)&10s-100&2006\\\\ Kepler&\t$1 R_\\oplus$ at 1 AU (500 pc)&100s&2008\\\\ GAIA &  $10 R_\\oplus$ at $<$0.1 AU &4,000&2012\\\\ JWST & \t$2 R_\\oplus$ (100-500 pc) &250$^2$ &2013\\\\ \\hline \\multicolumn{4}{l}{\\it Determine Masses/Orbits (Astrometry)} \\\\ SIM& $ 3 M_\\oplus$ at 1 AU (10 pc)&250&2012\\\\ GAIA& $ 30 M_\\oplus$ at 1 AU ($<$200 pc)&10,000&2012\\\\ \\hline \\multicolumn{4}{l}{\\it Characterize Planets and Search for Life (Direct Detection)}\\\\ JWST&$ 1 R_{Jup}$ at $>30$ AU (50-150 pc) &250 young stars &2013\\\\ TPF-Coronagraph$^3$&$ 1 R_\\oplus$ at 1 AU (15 pc)&250&2018\\\\ TPF-Interferometer/Darwin$^3$&$1 R_\\oplus$ at 1 AU (15 pc)&250&2018\\\\ \\end{tabular} \\tablenotetext{1}{Yield is highly approximate and assumes roughly 1 planet orbiting each star surveyed.} \\tablenotetext{2}{Approximate number of follow-up of ground-based, COROT and Kepler transit targets.} \\tablenotetext{3}{Parameters for TPF-C and TPF-I/Darwin are still being developed.} \\end{table*}  Papers in these proceedings describe many results from planet-finding experiments. With the exception of a few (very exciting) HST and Spitzer observations, these are drawn primarily from ground-based observations. This chapter on space activities necessarily focuses on future activities. While the space environment is highly stable and offers low backgrounds and an unobscured spectral range, the technology needed to take full advantage of the environment will take many years to develop and the missions to exploit the technology and the environment will be expensive to construct, launch and operate. But when the missions described here are completed, they will revolutionize our conception of our place in the Universe.  Table~\\ref{projtable} identifies the major space-based projects presently under consideration grouped by observing technique: transit photometry, astrometry, and direct detection. Figure~\\ref{projfig} summarizes the discovery space of these projects as well as some ground-based activities in the Mass--Semi-Major Axis plane. COROT, Kepler and SIM will provide the many order-of-magnitude improvements in sensitivity and resolution relative to ground-based efforts needed to detect other Earths in the habitable zones of their parent stars using indirect techniques of transits, radial velocity and astrometry.   \\begin{figure*} \\epsscale{1.0} \\plotone{PlanetFindingSummary.eps} \\caption{\\small A wide variety of space- and ground-based capabilities are summarized in this figure showing the detectability of planets of varying sizes and orbital locations. Space-based techniques become critical as one goes after terrestrial planets in the habitable zone. For a detailed description of this figure, see Lawson \\etal  (2004).} \\label{projfig} \\end{figure*}    Once we have detected these planets via indirect means ($\\S$2 and $\\S$3), we argue below that we will need a number of missions to characterize these planets physically and to search for evidence of life in any atmospheres these planets may possess. Analogously to the improved knowledge gained about astrophysical phenomena from using all of NASA's four Great Observatories (the Hubble Space Telescope, the Compton Gamma Ray Observer, the Chandra X-ray Observatory, and the Spitzer infrared telescope) and ESA's Cornerstone missions (ISO and XMM), a comparable ``Great Observatories'' program for planet finding will yield an understanding of planets and the search for life that will greatly exceed the contributions of the individual missions.  As discussed in $\\S$2, the combination of transit photometry with COROT and Kepler, follow-up spectrophotometry from the James Webb Space Telescope (JWST) and follow-up radial velocity data gives unique information on planetary mass, radius, density, orbital location, and, in favorable cases, composition of the upper atmosphere. These data, available for large numbers of planets, will revolutionize our understanding of gas-giant and icy-planets. While less information will be available for smaller, rocky planets, a critical result of the transit surveys will be the frequency of Earth-sized planets in the habitable zone, $\\eta_\\oplus$ (Beichman 2000). Since the angular resolution and collecting area needed for the direct detection of nearby planets are directly related to the distance to the closest host stars, the value of $\\eta_\\oplus$ will determine the scale and cost of missions to find and characterize those planets.  Subsequent to the transit surveys, we will embark on the search for and the characterization of nearby planets, and the search for a variety of signposts of life. We will ultimately require three complementary datasets: masses via astrometry (SIM PlanetQuest, $\\S$3); optical photons (TPF-Coronagraph, $\\S$4); and mid-IR photons (TPF-Interferometer/Darwin, $\\S$5). JWST will play an important role in follow-up activities looking at ground-based, COROT and Kepler transits; making coronagraphic searches for hot, young Jupiters; and studying proto-planetary and debris disks. The synergy between the planet finding missions, with an emphasis on studies of nearby, terrestrial planets and the search for life, is addressed in $\\S$6. Studies of potential target stars are ongoing but will need to be intensified and their results collated to select the best targets ($\\S$7). The ordering of these missions will be the result of the optimization of a highly non-linear function incorporating technical readiness, cost, and political and scientific support on two or more continents ($\\S$8). ",
        "conclusions": "The exploration of extrasolar planetary systems is a rich and diverse field. It calls for measurements with many kinds of instruments, as well as theoretical studies and numerical modeling. To discover and characterize extrasolar planets that are habitable and to be sure beyond a reasonable doubt that we can detect life, we need to measure the statistical distribution of planet diameters, the masses of nearby planets, and the spectra at visible and infrared wavelengths. The missions that can carry out these measurements are COROT/Kepler, SIM, TPF-C, and TPF-I/Darwin. JWST will provide followup observations for the transit missions as well as make important measurements of protostellar and debris disks. Each of these missions is a vital element of the program. Not only does each mission by itself provide its own compelling science, but together these missions form a coherent approach that will advance our understanding better than any single one by itself.  The exciting scientific promise described in preceding sections will not happen cheaply or overnight. The Great Observatories program for astrophysics spanned more than a decade between the first and last launches --- HST in 1990 and Spitzer in 2003 --- and was still longer in gestation. While the transit missions COROT and Kepler are being readied for flight, the ``Great Observatories'' for planet finding will take a generation of scientific and political advocacy: SIM PlanetQuest has completed its technology development, has been endorsed repeatedly by the US science community, and awaits final NASA approval; TPF-C and TPF-I/Darwin are well along in their programs of technology development and will need strong endorsement by US and European scientific communities in coming years. Interest is TPF goals is growing in Japan as well (Tamura \\& Abe 2006). JWST is moving into its construction phase and will provide many observations useful to the planet finding endeavour.  While the study of extrasolar planets is a new field of research with a relatively small number of (young) practitioners, our field is growing rapidly. We will fare well in any assessment of the importance of our field to progress in astronomy and of the great interest our program holds for the general public. We will also fare well in any assessment of the value of these planet finding facilities for general astrophysical investigations: SIM has already allocated a significant amount of observing time to a broad suite of general astrophysics; TPF-C will provide wide-field optical imaging with unprecedented sensitivity and angular resolution to complement JWST\u0092s mid- and near-IR capabilities; TPF-I/Darwin will break new ground with milliarcsecond mid-IR imaging and micro-Jy sensitivity.  Near-term political considerations must not discourage us as we plan this new era of planetary exploration. From the funding agencies, we must demand the budgets needed to nurture young scientists and senior researchers, to prepare the difficult technologies of nulling, large space optics, and formation flying, and eventually to build these ``Great Observatories.'' With our colleagues, we must argue forcefully for a balanced program based on scientific priorities free from parochial considerations of individual facilities or institutions. 20$^{th}$ Century cosmologists expanded our conception of the Universe with the discovery of galaxies, the expanding universe, and dark matter. 21$^{st}$ Century planet finders will expand our conception of humanity's place in the Universe with the discoveries of other habitable worlds and possibly of life itself."
    },
    "0601/hep-th0601153_arXiv.txt": {
        "abstract": "We study the dynamics of Nambu--Goto strings with junctions at which three strings meet.  In particular, we exhibit one simple exact solution and examine the process of intercommuting of two straight strings, in which they exchange partners but become joined by a third string.  We show that there are important kinematical constraints on this process.  The exchange cannot occur if the strings meet with very large relative velocity.  This may have important implications for the evolution of cosmic superstring networks and non-abelian string networks. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601108_arXiv.txt": {
        "abstract": "The transformations taking place in late-type galaxies in the environment of rich clusters of galaxies at $z=0$ are reviewed. From the handful of late-type galaxies that inhabit local clusters, whether they were formed in-situ and survived as such, avoiding transformation or even destruction or if they are newcomers that recently infall from outside, we can learn an important lesson on the latest stages of galaxy evolution. We start by reviewing the observational scenario, covering the broadest possible stretch of the electromagnetic spectrum, from the gas tracers (radio and optical), the star formation tracers (UV and optical), the old star tracers (Near-IR) and the dust (Far-IR). Strong emphasis is given to the three nearby, well studied clusters Virgo, A1367 and Coma, representative of different evolutionary stages, from unrelaxed, spiral rich (Virgo) to relaxed, spiral poor clusters (Coma). We continue by providing a review of models of galaxy interactions relevant to clusters of galaxies. Prototypes of various mechanisms and processes are discussed and their typical time-scales are given in an Appendix.\\\\ Observations indicate the presence of healthy late-type galaxies falling into nearby clusters individually or belonging to massive groups. More rare are infalling galaxies belonging to compact groups where significant pre-processing might take place. Once entered the cluster, they loose their gas and quench their star formation activity, becoming anemics. Observations and theory agree in indicating that the interaction with the intergalactic medium is responsible for the gas depletion. This process, however, cannot be at the origin of the cluster lenticular galaxy population. Physical and statistical properties of S0 in nearby clusters and at higher redshift, indicate that they originate from spiral galaxies transformed by gravitational interactions. ",
        "introduction": "\\label{Introduction}  On scales $<100$ Mpc, the size of the largest known coherent matter aggregates,  the density of galaxies in the local universe spans from $\\sim$  0.2 $\\rho_0$ in voids to $\\sim$ 5 $\\rho_0$ in superclusters and filaments, $\\sim$ 100 $\\rho_0$ in the cores of rich clusters, up to $\\sim$ 1000 $\\rho_0$ in compact groups, where $\\rho_0$ is the average ``field'' density \\citep{GELH89}.\\\\ It has been known for decades \\citep{HUBH31} that galaxy Hubble type and local density are not independent quantities. In their analysis of nearby clusters, \\cite{DRE80} and \\cite{WHIG93} agreed that the fraction of spiral galaxies decreases from 80\\% in the ``field'', to 60\\% in the outskirts to virtually zero in the cores of rich clusters, compensated by an opposite increase of elliptical and S0 galaxies. Morphological segregation is perhaps the clearest signature of the environmental dependence of the processes that govern the formation and evolution of galaxies. Understanding the morphology segregation,  i.e. the shaping of the Hubble sequence in the various environments, has challenged astronomers for years: \"Nature or Nurture\"? The issue is still hotly debated, as it involves many controversial fields of modern astrophysics, including observations, simulations and theories.\\\\ With the advent of 10m class telescopes we hope that galaxy evolution will be shortly disclosed observationally, much as in a rewound movie taken along the ``fossil'' sequence, from today's fully evolved nearby galaxies to the early objects forming their first stars far away in space and time. Similarly we can imagine to shoot the movie of the morphology segregation by comparing frames of clusters taken at increasingly large $z$. Snapshots at $z=0$ would show the fully formed and evolved clusters with their final mix of E/S0/S. Frames taken between $z=0.2$ and 0.5 would show an unchanged fraction of E, with the action focused on increasingly active spirals (the Butcher-Oemler effect \\citep{BUTO78, BUTO84}) and a decreasing fraction of S0 \\citep{DRES97, DRES04b}. By $z=0.5-1$ the scene would be on clusters under formation, when large groups of galaxies coalesce. Effective \"pre-processing\" is taking place in these low velocity-dispersion environments, where minor merging and tidal interactions among galaxies are shaping up the S0s \\citep{MIHOS04, FUJI04}. None has yet observed elliptical galaxies under formation at $z=3$ and beyond. But there is general consent that they should have formed at these early epochs, very much from \"nature\" and little from \"nurture\", even before clusters existed \\citep{DRES04b, TREU04, NOLA04, FRAN04}. Maybe in a decade from now, when the JWST will provide these last photograms, the full movie of the shaping of the Hubble sequence will be on view, as many groups of researchers are focusing large efforts on it (the Morphs group \\citep{DRES99}, the CNOC1 group \\citep{YEE96}, the ACSCS team \\citep{POST96}, the EDisCS group \\citep{HALL04} and the team of van Dokkum and collaborators \\citep{VANDOK00}).\\\\ This is perhaps somewhat optimistic, however. Pretending to understand galaxy evolution really means shooting a movie whose individual photograms are composed of SEDs, obtained combining best quality data taken over a large stretch of the electromagnetic spectrum, each individual band providing information on the evolution of a particular galaxy component. Restframe UV fluxes and Balmer emission line strengths are the best current, massive star-formation tracers. Optical bands provide the morphology, NIR bands trace the oldest stars, MIR and FIR the dust (and extinction); the radio (continuum) tells about the nuclear activity, the supernova rate and the magnetic fields, and the 21 cm line the primordial gas reservoir and its dynamics.\\\\ The problem is that, as $z$ increases, characteristic features become unobservable or more difficult to observe. By $z=1$, for example, the H$\\alpha$ line ends up in the H band, a noisy place to carry out observations. The peak of the dust emission falls in the sub-mm bands where the sensitivity of telescopes will remain poor until ALMA. The 21 cm line drops to lower sensitivity and angular resolution bands where line correlators do not yet exist (we are not aware of a single HI emission measurement taken beyond $z$=0.15). Furthermore, the movie's linear resolution decreases with $z$, so that detailed morphological information quickly fades away. Most dramatic is the lost of sensitivity with increasing $z$. By $z<0.2$ dwarf galaxies run completely out of the scene. Meaningful, high-resolution SEDs in the whole luminosity range covered by galaxies are obtainable only at $z=0$, providing the boundary conditions to models and simulations of galaxy evolution.\\\\ The end point of the evolution is within the reach of modern telescopes. Significantly less evolved galaxies can only be observed, with lower resolution, if they are more luminous. Dwarf galaxies, Ims, BCDs and dEs, by far the most common types of galaxies in the universe, yet the less understood, can only be observed locally. \\\\ Since the rate of evolution of galaxies is found to correlate primarily with their dynamical mass (infrared luminosity) according to a remarkably anti-hierarchical scheme (\\cite{GAV03, GAVSC96, GAVPB96a, BOSG01, GAVB02c, 2002A&A...384..812S}) that has been lately named ``downsizing'' , distant surveys might undersample low-luminosity, slowly evolving systems. This is why studying nearby clusters still is a worthwhile activity. \\\\ Another pro of local clusters is that they are ideal laboratories for studying interaction mechanisms at high resolution. Prototypes of all relevant mechanisms can be found in nearby clusters. Abell 1367, for example, contains perhaps the best known example of currently ongoing ram-pressure stripping (CGCG 97-073), a rare example of a merging system in a cluster (UGC 6697) and even a local prototype of infalling group where pre-processing occurs at the present cosmological epoch (see Section \\ref{prototype}). Infall was much more effective, but unfortunately hard to observe, in the past ($z \\sim 0.5$), when it contributed shaping the S0 galaxies (pre-processing).\\\\ So far we insisted on the importance of observing nearby clusters. Among them we have focused the present review on three specific ones: Virgo, Abell 1367 and Coma that are among the best studied in the universe. The variety of their conditions, e.g. spiral fraction, X-ray luminosity, kinematic properties, makes them suitable ``laboratories'' for comparing clusters in different evolutionary stages: near formation (Virgo), still significantly accreting (A1367), near virialization (Coma).\\\\ Honestly speaking, the present review was conceived while we were assembling ``GOLDmine'', the galaxy database that we made available on the WEB (goldmine.mib.infn.it) \\citep{GAVB03a}. Two third of the galaxies contained in GOLDmine belong in fact to these three clusters and a significant fraction of them are among the best sampled SEDs. This does not mean that we haven't tried to review the work of other groups. Certainly we haven't done it as extensively as we should, and we apologize for that.\\\\ But why focusing the present review on late-type galaxies, that are less numerous, they evolve slower than the more intriguing S0s and Es? What can we learn from the handful late-type galaxies that still inhabit local clusters? There are several reasons for concentrating on them. One is that, due to the fragility of their stellar and gaseous disks, they are sensitive probes of their environment. Another is that, since their rate of evolution (star formation history) is mild and continuous, they provide a linear clock for evolution, as opposed to E whose action occurred suddenly in one brief and intense episode that occurred early in their life. Spirals are also interesting because significant infall of galaxies of this type has occurred and is still occurring (at decreasing rate due to the $\\Lambda$ dominated universe) onto clusters of galaxies. Will they sooner or later evolve into lenticulars? Can we use them as probes of the cluster environment? \\begin{figure*}[!htb] \\epsscale{0.7} \\plotone{fig01.ps} \\caption{The celestial distribution of galaxies brighter than 15.7 from the CGCG catalogue, \\cite{CGCG} in a 4x4 $\\rm deg^2$ box around the Coma cluster. The 146 early-type (E+S0+S0a) galaxies (empty symbols) clearly mark the cluster density enhancement, whereas the 49 late-type ($\\geq$Sa) objects (filled symbols, including 10 unclassified spirals) hardly trace the cluster. The circle of 1 degree radius is traced about  the X-ray center. The X-ray contours from XMM are superposed.} \\label{coma_map} \\end{figure*} To help better focusing the addressed issues, let's take a look at the Coma cluster whose bright galaxy distribution is shown in Figure \\ref{coma_map}. The early-type galaxies clearly mark the cluster density enhancement, whereas the late-type objects hardly trace the cluster at all, making \\cite{ABE65} to doubt their very membership to the Coma cluster. The latter are very ``healthy'' objects, actively star forming, as witnessed by their average H$\\alpha$ E.W. (equivalent width) of 35 \\AA~ (with peaks of 78 \\AA). This perhaps explains why Abell questioned their membership to Coma, because the majority of galaxies observed in clusters shows no signs of recent star formation activity. He didn't know that some of these objects are not as in ``good shape'' as it looks at first glance, something that was realized 20 years later. The pioneering work of \\cite{GIOH85}, continued by many other researchers, among them \\cite{GAV87, GAV89}, showed in fact that galaxies at projected distances $<$ 1 Mpc from the center of Coma suffer from significant neutral hydrogen deficiency (see Section \\ref{The atomic gas}), having lost from 30 \\% to 90 \\% of their original HI content. The first resolved map of their HI distribution was obtained even more recently by \\cite{BRAC00}. It shows that in most highly or moderately deficient objects the HI distribution is asymmetric or significantly displaced from the parent optical galaxy. This phenomenon occurs in objects projected within the X-ray emitting gas in the Coma cluster. For the interpretation of this observation, whether is due to ram-pressure stripping or to galaxy-galaxy or galaxy-cluster interactions, the reader is referred to the remaining of the discussion. However, generally speaking, a high HI deficiency represents perhaps the clearest signature of the interaction occurring to galaxies in the environment of rich clusters, in other words it gives a clear clue to their cluster membership. Abell was wrong! These are bona fide cluster members. What future is reserved to these galaxies? This is the type of issue addressed by the present review. The reader can find many excellent review papers in the literature that deal with the evolution of galaxies in the cluster environment \\citep{HAYGC84,DRES84,SAR86,DRES97,POGB01a,DRID03,VANGO03} each focused on a particular aspect or wavelength domain. The most recent achievements in cluster research are collected in the proceedings of the Carnegie Observatories Astrophysics Series, Vol. 3: Clusters of Galaxies: Probes of Cosmological Structure and Galaxy Evolution (2004). With the present review we wish, for the first time, to consider together all aspects. The price we pay is that we must narrow down the reviewed interval of $z$, in practice to $z$=0. We start by reviewing the observational scenario (Sections \\ref{constituents}, \\ref{obs_diff}), covering the broadest possible stretch of the electromagnetic spectrum, from the gas tracers (radio and optical), the star formation tracers (UV and optical), the old star tracers (Near-IR) and the dust (Far-IR). We continue by providing a review of models of galaxy interactions relevant to clusters of galaxies (Section \\ref{models}). We conclude with a discussion on what proposed model(s) better interprets various aspects of the observational scenario (Section \\ref{Discussion}), once again checking the models predictions on recent data on Virgo, Coma and A1367. ",
        "conclusions": "\\label{Discussion}  Conclusive evidence built up in the last 25 years that, beside the morphological segregation effect, late-type galaxies belonging to rich nearby clusters differ significantly from their ``field'' counterparts. Primarily they are found deficient in atomic gas with respect to isolated galaxies (see sect. \\ref{The atomic gas}). Large HI deficiency parameters provide the clearest signature of their cluster membership. Because of the lack of gas, their star formation is significantly quenched (see sect. \\ref{star formation}), despite the fact that the molecular gas component seems little perturbed (see sect. \\ref{The molecular gas}). Additional, but less clear-cut differences are: an excess of radio continuum activity, a higher metallicity and possibly a lower dust mass content of galaxies in high density environments. Kinematic properties instead (i.e. rotation curves), witnessing the dark matter distribution, seem little affected in clusters (see sect. \\ref{constituents} and \\ref{obs_diff}).  Which physical processes, among the ones reviewed in sect. \\ref{models}, are affecting the evolution of late-type galaxies falling on today's clusters? The gravitational perturbations seem unlikely. Tidal interactions among galaxies are uneffective simply because in today's high velocity dispersion clusters they are too short-lived to produce severe disturbances (sect. \\ref{models} and \\ref{timescales}). The gravitational interaction with the cluster potential might induce some nuclear gas infall, but it can hardly remove the outer disk HI and thus originate the cluster HI-deficient galaxy population (sect. \\ref{tidalcluster}). Gas stripping in HI-deficient galaxies seems too a recent event (some 10$^8$ yr) for being produced by galaxy harassment, whose time scale is $\\sim$ ten times longer (sect. \\ref{harassment}).\\\\ In today's universe the interactions with the IGM are likely to be dominant. Among interactions with the cluster IGM, evaporation is effective in hot Coma-like clusters, especially for small galaxies, unless they are shielded by relevant magnetic fields (sect. \\ref{evaporation}). Ram pressure seems however the most relevant process whenever the gas is not too hot (A1367, Virgo) and provided that the cluster crossing velocity is high (see sect. \\ref{Ram}). This condition seems satisfied since the orbits of spirals in nearby clusters are predominantly radial, as indicated by their non-Gaussian velocity distribution (see sect. \\ref{infall}). Dynamical simulations confirm this picture also for the young, unrelaxed Virgo cluster, whose IGM density is not particularly high. Gas stripping might be locally more effective if the IGM has a non-uniform, clumpy distribution. Considering that the effects of ram-pressure, thermal evaporation and viscous stripping are probably additive, on time scales comparable to a crossing time, any late-type galaxy infalling into a rich cluster is expected to loose most of its atomic gas reservoir, in particular the weakly gravitationally bound one located in the outer disk.\\\\ Observations corroborate the idea that galaxy - IGM interactions are the dominant phenomena affecting the evolution of late-type cluster galaxies in nearby clusters. The location of most severely HI-deficient late-type galaxies, with the smallest HI disks near the  peaks of the IGM gas density argues in favor of the hydrodynamic scenario (sect. \\ref{The atomic gas}). \\begin{figure}[!ht] \\epsscale{0.8} \\plotone{fig24.ps} \\caption{The correlation between the HI deficiency parameter and the square of the velocity deviation along the line of sight, divided by the angular distance from the center of cluster A in Virgo. Empty dots give averages in regularly spaced bins.} \\label{defagg30} \\end{figure} There is also evidence that spirals in the Virgo main cluster A show a weak correlation (see Fig. \\ref{defagg30}) between the HI deficiency parameter and the quantity $(1/\\theta) \\times \\Delta V_{||}^2$, proportional to the square of the velocity along the line of sight divided by the angular distance from the cluster center (a quantity that is expected to increase with the the ram-pressure $\\rho~V^2$), arguing in favor of the ram-pressure hypothesis (notice that $V^2$ is very poorly determined from the observable component of the velocity parallel to the line of sight). Certainly this evidence does not support the tidal interaction scenario where $HI-def$ should rather decrease with $\\Delta V_{||}$.\\\\ The truncation of the star forming disk in cluster HI-deficient spirals (see sect. \\ref{star formation}) can also be explained in the ram-pressure scenario because the star formation is progressively suppressed once the gas is removed from the outer disk, as in NGC 4569. The interaction with the cluster IGM might sometime increase, on short time scales and in specific configurations, the total star formation activity of galaxies, as observed (i.e. CGCG 97-073 in A1367) and predicted by models (sect. \\ref{Ram}), but on long timescales the galaxy global activity should however decrease.\\\\ On the contrary gravitational interactions are expected to trigger star formation, in particular in the nuclei (sect. \\ref{tidal}), while the observations have unambiguously shown that on average the star formation activity of giant late-type galaxies decreases towards the cluster center. There are also evidences that cluster galaxies did not undergo strong bursts of star formation, as their constant H$\\alpha$ to UV flux ratio implies, and that not all cluster galaxies with a truncated gas or young stellar disk have enhanced nuclear activity (gravitational interactions are expected first to trigger nuclear gas infall and only in the strongest cases to produce disk truncation) (see sect. \\ref{star formation}). \\\\ Several prototypical galaxies provide convincing cases that the interactions with the cluster IGM are predominant in the nearby universe. NGC 4569, the typical HI-deficient, anemic galaxy in the Virgo cluster has kinematic and spectro-photometric disk properties that mimic a recent (300-400 Myr) ram-pressure stripping event (see sect. \\ref{prototype}). We cannot exclude that some nuclear gas infall has been induced by the gravitational interaction with the potential of the cluster. Broadly speaking, the fact that most of the HI-deficient galaxies have truncated HI and H$\\alpha$  disks,  but not at longer wavelengths confirms that in most cases the interaction responsible for the gas depletion is relatively recent ($\\leq$ 5 10$^8$ yr), as in NGC 4569. The existence of several marginally HI-deficient cluster spirals with asymmetric HI-profiles provides further evidence that the galaxy-IGM interaction is a recent event. The long, ionized trails behind CGCG 97-073 and CGCG 97-079 impose even stronger time constraints, indicating that in a few cases the interaction is still undergoing (sect. \\ref{prototype} and sect.\\ref{timescales}).  The hydrodynamical galaxy - IGM interaction scenario poses some problems that might however only be apparent. It has been claimed that observing the threshold in the star formation activity at $\\sim$ one virial radius, i.e. at clustercentric distances where the IGM gas density is low \\citep{TREU03}, argues against the ram-pressure stripping scenario. This argument has been used in favor of other processes dominating the evolution of cluster gas-rich galaxies, such as pre-processing and starvation \\citep{BALN00, LEWB02, BOWB04, DRES04b, NICH04}. First of all this result is based on wide field surveys, such as SDSS and 2dF, where HI data are unavailable and the optical galaxy classification is based on light profiles and/or spectra, without direct plate inspection. The threshold in the mean star formation activity of late-type cluster galaxies might be contaminated by the morphology segregation effect.  Furthermore we have stressed that the evidence offered by CGCG 97-079 and 97-073 in A1367 that ram-pressure is active at least up to $\\sim$ 0.8 Mpc (0.4 virial radius) combined with other  examples showing that the effects of stripping might last until $\\sim$ 1 virial radius, reduces the necessity of invoking other processes such as pre-processing or starvation acting in the outskirt of clusters at $z$=0. Compact groups accreting on clusters, such as the Blue Infalling Group in A1367 (sect. \\ref{prototype}), where effective pre-processing might take place, are extremely rare at $z$=0. We have in fact evidence that galaxies infall as individuals (CGCG 97-073 and CGCG 97-079), or belonging to extended sub structures with velocity dispersion comparable to already formed clusters ($\\geq$ 500 km s$^{-1}$, sect. \\ref{infall}).  A unified picture that combines near and far, present and past can be drawn coherently if one takes into account that the environmental conditions and the physical properties of galaxies have evolved significantly during the cosmic time, making the ranking of the mechanism that prevailed in shaping the Hubble sequence to change with time. The number of low velocity dispersion infalling groups, where pre-processing is effective, increased with $z$; the density of the IGM and its temperature decreased with $z$. The gas fraction and the activity of star formation of galaxies, independently of their environmental conditions, increased significantly from $z$=0 at least up to $z$=1.  Given that the interaction with the cluster IGM cannot easily thicken the disk of a spiral galaxy and transform it into an S0 (sect. \\ref{Surface brightness})\\footnote{In the ram-pressure scenario bulges can form only under particular conditions, such as through the triggering of cycles of annealing after several crossings of the cluster \\citep{SCHS01} or edge-on stripping, when up to 50\\% of the gas can be re-accreted by the disk \\citep{VOLC01a} (see sect. \\ref{Ram})}, the formation of bulges is obtained straightforwardly from disk heating (perturbing or even stopping the rotation) during gravitational interactions with nearby companions and/or with the cluster potential (galaxy harassment) (sect. \\ref{tidal}, \\ref{tidalcluster}, \\ref{harassment}). The observed kinematical, structural and spectrophotometrical properties of bright lenticulars in nearby clusters are consistent with this interpretation. For low-luminosity lenticulars, however, as stated by \\cite{POGB04}, there are some evidences that the suppression of the star formation was due to their interaction with the IGM (sect. \\ref{Surface brightness}). These considerations have induced groups of researchers working on clusters at high $z$ to identify in gravitational pre-processing the most likely mechanism at the origin of the morphological segregation \\citep{DRES04b}, although others invoked smoother processes such as starvation \\citep{BALN00}. We refer the reader to several papers that appeared in the Carnegie Observatories Astrophysics Series, Vol. 3: Clusters of Galaxies: Probes of Cosmological Structure and Galaxy Evolution (2004) elaborating on this idea.  Several problems concerning the evolution of galaxies in clusters remain unsolved in this unified scenario. One deals for example with the origin of dwarf galaxies. Although the strong HI-deficiency of cluster dwarf irregulars expected in the hydrodynamical interaction scenario cannot be proved because of the lack of sensitive HI observations (sect. \\ref{The atomic gas}), there is evidence that their star formation activity does not decrease with clustercentric distances as for massive objects (sect. \\ref{star formation}), maybe due to continuous replenishment of gas-rich dwarfs from outside, or maybe because dwarf irregulars, after gas removal, quickly fade into dEs, thus they drop out of late-type surveys. A close link between dIs and dEs, first proposed by \\cite{LINF83}, is strengthened by the existence of rotationally supported dEs \\citep{VANZ04a,VANZ04b}, by the residual/past star formation activity of dEs observed in UV \\citep{BOSC05c} or in optical spectra (i.e. VCC1499, \\cite{GAVZ01, CONS03b}), by the existence of dEs not totally  devoid of gas \\citep{CONS03a}, by the shape of the faint end of the cluster luminosity functions \\citep{CONS02} and by the velocity distribution as broad as that of dwarf irregulars \\citep{CONG01}.  It seems fair to conclude that gravitational perturbations in infalling groups are at the origin of local lenticular galaxies. However the process dominating the evolution of gas-rich, late-type galaxies falling-in at the present epoch is the interaction with the hot and dense IGM. The fate of late-type galaxies will be slightly different from what it would have been a Gyr ago: they will become HI-deficient, anemic, rotationally-supported spirals. On long timescales (several crossing times), because of the lack of gas supply, they will lose angular momentum, somewhat heating their stellar disks. They will probably become disk dominated, quiescent systems (\\ref{Surface brightness}), but they will never make it to become real S0s.\\\\ If this evolutionary picture is correct, we still have to understand at which epoch and under what environmental conditions the hydrodynamical and the gravitational processes interchanged their prevailing role and, among other things, compare it with the inset of the Buchler-Oemler effect. These observational goals are within the reach of last generation ground based optical telescopes and of the HST. However we must wait another decade before neutral hydrogen measurements of galaxies at significant $z$ with SKA will shade the final light on the evolutionary processes that took place in clusters of galaxies.   \\newpage"
    },
    "0601/astro-ph0601614_arXiv.txt": {
        "abstract": "We investigate the quantity and composition of unburned material in the outer layers of three normal Type Ia supernovae (SNe Ia): 2000dn, 2002cr and 2004bw. Pristine matter from a white dwarf progenitor is expected to be a mixture of oxygen and carbon in approximately equal abundance.  Using near-infrared (NIR, $0.7-2.5$ \\mum) spectra, we find that oxygen is abundant while carbon is severely depleted with low upper limits in the outer third of the ejected mass. Strong features from the \\oi\\ line at $\\lambda_{rest} = 0.7773$ \\mum\\ are observed through a wide range of expansion velocities $\\approx 9 - 18 \\times 10^3$ \\kms.  This large velocity domain corresponds to a physical region of the supernova with a large radial depth.  We show that the ionization of C and O will be substantially the same in this region.  \\ci\\ lines in the NIR are expected to be $7-50$ times stronger than those from \\oi\\ but there is only marginal evidence of \\ci\\ in the spectra and none of \\ion{C}{2}.  We deduce that for these three normal SNe Ia, oxygen is more abundant than carbon by factors of $10^2 - 10^3$.  \\mg\\ is also detected in a velocity range similar to that of \\oi.  The presence of O and Mg combined with the absence of C indicates that for these SNe Ia, nuclear burning has reached all but the extreme outer layers; any unburned material must have expansion velocities greater than $18 \\times 10^3$ \\kms.  This result favors deflagration to detonation transition (DD) models over pure deflagration models for SNe Ia. ",
        "introduction": "There is general agreement that Type Ia supernovae (SNe Ia) result from the combustion of a degenerate carbon/oxygen (C/O) white dwarf (WD) star as first predicted by \\citet{hoyle60}.  The currently accepted scenario for SNe Ia is the explosion of a C/O WD near the Chandrasekhar mass ($M_{Ch}\\approx 1.44M_{\\odot}$). The explosion is triggered by compressional heating near the center of the WD.  A C/O WD is the final evolutionary stage for stars with a main sequence mass of $\\approx 3-8 M_{\\odot}$. Initially the progenitor has insufficient mass to produce a SN Ia.  To reach $M_{Ch}$, the WD must accrete additional matter through Roche-lobe overflow from a main sequence, red giant, or WD companion.   Quasi-steady shell burning of accreted material in conditions of low pressure and high temperature produces a mixture of C and O with a mass ratio near unity \\citep{dominguez01}.  If two WDs are entirely mixed in a merger scenario \\citep{webbink84}, the C to O ratio by mass will also be close to 1:1 \\citep{benz89}.  One of the keys to understanding the explosion is the propagation of the nuclear flame.  Two modes of burning can be distinguished: {\\em detonation}, with burning velocities slightly above the speed of sound, and {\\em deflagration}, with burning front velocities well below sound speed.  The two models for SNe Ia currently under discussion are Pure Deflagration \\citep{ivanova74, nomoto76} and Delayed Detonation Transition (DD) \\citep{khokhlov91}.  Despite some differences of the numerical treatment of nuclear burning fronts by different groups, both models share the same initial sequence: nuclear burning begins near the center of the WD and the burning front becomes Rayleigh-Taylor unstable.  These instabilities result in large scale mixing of burned and unburned matter.  During the early burning phase, energy is deposited into the WD causing the outer layers to expand with velocities close to the speed of sound \\citep{hk96}.  Classical deflagration models cease burning in this phase and leave a chemical structure characterized by a thick layer of unburned C and O surrounding inner regions of intermediate-mass and iron-peak elements.  Recent refinements with multi-dimensional deflagration models predict that, within the context of standard models, a substantial mass of unburned material is mixed toward the center of the SN and burning products are mixed outward toward the surface.  Primordial and processed elements intermingle throughout the radial chemical structure rather than forming distinct concentric layers \\citep{Gamezo03, Ropke05}.  Significant quantities of material remain unburned because the deflagration burning front is unable to catch up to the expansion of the WD near the surface \\citep{hk96, khokhlov01}.  Thus the outer layers in current deflagration models remain unburned by means of causality and large pockets of unburned C/O will be found at various depths because of the nature of the Rayleigh-Taylor instabilities.  These generic properties of pure deflagration models seem to be in contradiction to observations of so called ``Branch-normal'' SNe Ia \\citep{Branch93} that require vanishing amounts of unburned matter and radially layered chemical structures \\citep{Barbon90, pah95, Wheeler98, fisher98, Hamuy02, m03}.  DD models also begin with subsonic deflagration burning, but they make a transition to a supersonic detonation front.  For models of normal SNe Ia, the detonation eradicates most of the chemical inhomogeneities, producing a radially stratified chemical structure with very little unburned matter left at the surface.  We note that questions remain about how the deflagration to detonation transition occurs.  For recent reviews of SN Ia physical models see \\citet{Branch99}, \\citet{Hill00}, and \\citet{pah05}.  In this paper, we investigate the presence of unburned material in the outer layers of three normal SNe Ia by examining carbon and oxygen features in near-infrared (NIR, 0.7 - 2.4 \\mum) spectra.  The NIR is an excellent region for this search due to the presence of several strong lines from C and O ions.  We use observational evidence to make estimates for the relative abundance of O and C in these SNe Ia.  General methods are used for this analysis to produce upper limits on the abundance of these elements and avoid dependence on specific models or parameters.  Section 2 describes the observations including data acquisition and reduction as well as details of the supernovae.  Section 3 presents results with line identifications and measurements. Section 4 provides methods and calculations for C and O locations, ionizations, line strengths, and abundance ratios.  Section 5 provides a summary and our conclusions. ",
        "conclusions": "The amount of material remaining unburned after the explosion is an important discriminant between the leading models for SNe Ia.  Deflagration models predict that large quantities of unburned material will remain in the outer layers of the SN Ia after the explosion while Deflagration to Detonation Transition (DD) models predict that essentially all progenitor matter will experience nuclear burning.  There is agreement that within the context of standard models for stellar evolution, the composition of unburned material will be carbon and oxygen in approximately equal abundance.  Oxygen can be present in SNe Ia from either primordial material or as a product of carbon burning.  In unburned matter, C and O should both be detectable.  In C burning regions, C will have been consumed but Mg will be present as well as O.  The similarity of ionization potentials for carbon and oxygen, the small mean free path for UV photons, and features that cover several thousand \\kms\\ imply that carbon and oxygen will be in the same ionization state.  We have conducted a search for evidence of carbon and oxygen using three NIR spectra from normal SNe Ia.  A strong detection is made from \\oi\\ at 0.7773 \\mum\\ with a projected velocity range of $\\approx 9-18 \\times 10^3$ \\kms.  The line-forming region comprises about one third of the matter in SNe Ia and \\oi\\ extends to the outer layers.  We do not detect carbon at any velocity, in agreement with \\citet{Barbon90, pah95, fisher98, Stehle05}.  If any matter from these SNe Ia contains C and O in nearly equal abundance, it must have an expansion velocity in excess of 18,000 \\kms.  The upper limit on the mass of carbon beyond this velocity is: $3.6 \\times 10^{-2} M_{\\odot}$.  We use a DD velocity structure as a limiting case for these estimates because alternative models, such as deflagration models, produce even less mass above given expansion velocities.  Strong, broad lines from \\mg\\ are also found in all spectra in our sample.  We show that the \\mg\\ line-forming region covers a similar velocity range to \\oi, and hence a similar spatial region.  This result demonstrates that the domain of explosive carbon burning reaches the extreme outer layers of matter probed by these spectra.  We compare estimated line strengths and measured line depths for \\oi\\ and \\ci\\ to derive lower limits on the O to C abundance ratios for the SNe Ia in our sample.  We find that the oxygen abundance is greater than the carbon abundance by factors of $10^2-10^3$ at $\\approx$ 11,000 \\kms, and remains well above unity to velocities in excess of 18,000 \\kms.  This large imbalance and carbon depletion contradict pure deflagration models for standard progenitors that predict significant amounts of unburned C and O will remain after the explosion..  Patchiness of the carbon and oxygen regions is sometimes invoked as a possible explanation for the lack of detection of carbon in spectra from SNe Ia, but the outer layers of published deflagration models are not patchy.  The strong \\oi\\ absorption features in these spectra require a large covering factor and cannot be produced by thin clumps of processed material.  Our generalized treatment of the problem has limitations but it adds stability to our analysis and eliminates model dependencies.  The use of LTE departure coefficients introduces uncertainty in the estimated line strengths by factors of $2-3$.  The combined uncertainties in the measurement of expansion velocities are about $\\pm 750$ \\kms.  The estimated location of the local continuum affects measurements of line depths at less than two sigma noise levels, but is inconsequential at greater depths.   The sum of these uncertainties is not nearly large enough to change our results.  The presence of \\mg\\ in the same region as \\oi, combined with a lack of detection of \\ci\\ indicate that nuclear burning has taken place in at least the outer one third of the mass in these SNe Ia and extends to more than 18,000 \\kms.  We conclude that for these SNe Ia, carbon is absent at radial depths below a minimum velocity of 18,000 \\kms\\ because it has been explosively burnt to O and Mg.  This conclusion is in agreement with DD models for SNe Ia and contradicts pure deflagration models but it does not address the existence of C and O in SNe Ia at velocities below 9,000 \\kms or above 18,000 \\kms.  \\citet{Quimby05} address the extreme outer layers by showing that \\ion{Si}{2} is present in optical spectra of SNe Ia to velocities of $\\approx$ 24,000 \\kms.  Silicon in this region can only be produced from a C/O WD progenitor by explosive oxygen burning at temperatures greater than those required to burn carbon into oxygen and magnesium.  Non-standard stellar evolution processes may be able to enrich oxygen in the outer layers and reduce the amount of carbon.  This possible explanation for the lack of observed carbon is severely constrained by the presence of high temperature burning products in the same region.  Carbon abundance in the inner regions of the ejecta are constrained by \\citet{Kozma05} using nebular spectra to provide very low limits on the mass of unburned material remaining in the center of SNe Ia after the explosion.  The results of \\citeauthor{Quimby05} and \\citeauthor{Kozma05} combine nicely with the evidence we provide in this paper to demonstrate that nuclear burning must reach every part of the progenitor.  The lack of carbon detection in all regions of normal SNe Ia contradicts predictions of published 3-D Deflagration models that unburned C will be present at all velocities.  On the other hand, these observations are perfectly consistent with predictions of Deflagration to Detonation Transition models for SNe Ia.   \\noindent{\\bf Acknowledgments:}  We thank W. D. Li and M. Papenkova for providing light curve data for many of the SNe in our sample.  We thank the individuals at the IRTF for guidance and help with the observations.  In particular, Alan Tokanaga, John Rayner, Mike Cushing,  Bill Golisch,  Dave Griep, and Paul Sears have been most helpful.  We would also like to thank the TAC of the IRTF for support and instructive comments.  GHM would like to thank Mike Cushing for helpful comments.  This research is supported in part by NSF grant 0098644 and by NASA grant NAG5-7937.  CLG is supported through UK PPARC grant PPA/G/S/2003/00040."
    },
    "0601/astro-ph0601564_arXiv.txt": {
        "abstract": "We study Schwinger pair creation of charged particles due to the inhomogeneous electric field created by the thin electron layer at the surface of quark stars (the electrosphere). As suggested earlier, due to the low photon emissivity of the quark-gluon plasma and of the electrosphere, electron-positron pair emission could be the main observational signature of quark stars. To obtain the electron-positron pair creation rate we use the tunnelling approach. Explicit expressions for the fermion creation rate per unit time per unit volume are derived, which generalize the classical Schwinger result. The finite size effects in pair production, due to the presence of a boundary (the surface of the quark star), are also considered in the framework of a simple approach. It is shown that the boundary effects induce large quantitative and qualitative deviations of the particle production rate from what one deduces with the Schwinger formula and its generalization for the electric field of the electrosphere. The electron-positron pair emissivity and flux of the electrosphere of quark stars due to pair creation is considered, and the magnitude of the boundary effects for this parameters is estimated. Due to the inhomogeneity of the electric field distribution in the electrosphere and of the presence of the boundary effects, at high temperatures ($T\\geq T_{cr}\\approx 0.1 $ MeV) we find a lower electron-positron flux as previously estimated. The numerical value of the critical temperature $T_{cr}$ depends on the surface potential of the star. We briefly consider the effect of the magnetic field on the pair creation process and show that the magnetic field can enhance drastically the pair creation rate. ",
        "introduction": "The Schwinger mechanism of pair production \\citep{Sch51}, first proposed to study the production of electron-positron pairs in a strong and uniform electric field, has been applied to many problems in contemporary physics. Strong electromagnetic fields lead to two physically important phenomena: pair production and vacuum polarization. A strong electric field makes the quantum electrodynamic vacuum (QED) unstable, and consequently it decays by emitting a significant number of boson or fermion pairs \\citep{Gr85}.  For a spin $1/2$ particle Schwinger's predicted production rate per unit time and volume $w$ is given by \\citep{Sch51,So82} \\begin{equation}\\label{prod} w(E_0)=eE_0\\int \\frac{d^{2}k_{i}}{\\left( 2\\pi \\right) ^{2}}\\sum_{n=1}^{\\infty }% \\frac{1}{n}\\exp \\left[ -\\frac{\\pi n\\left( m_{e}^{2}+k_{i}^{2}\\right) }{% \\left| eE_0\\right| }\\right] , \\end{equation} where $m_{e}$ and $e$ are the electron mass and charge, respectively, $k_{i}$ is the transverse momentum and $E_0$ is the (constant) electric field. In the present paper we use units so that $\\hbar =c =k_B =1$. In these units, $e$ is equal to $\\alpha^{1/2}$ and $1$ MeV$=5.064\\times 10^{-3}$ fm$^{-1}$.   In most of the physical applications the proper time method introduced by \\citet{Sch51} has been used to calculate the pair production rate. The real part of the effective action leads to vacuum polarization and the imaginary part to pair production. Though that method is conceptually well defined and technically rigorous, it is generally difficult to apply it to concrete physical problems such as inhomogeneous electromagnetic fields. Schwinger's result was generalized to electric fields $E_3=E_3\\left(x_{\\pm }\\right)$, which depend upon either light cone coordinates $x_{\\pm }=x_3\\pm x_0$,  but not upon both, in \\citet{To00}. The form of the result is exactly the same as in the original formula given by Eq. (\\ref{prod}). The case of electric fields depending on both $x_+$ and $x_{-}$, $E_3=E_3\\left(x_+,x_{-}\\right)$ was considered in \\citet{Av03}.  An alternative approach to pair creation was initiated by \\citet{Ca79,Ca80}, who re-derived Schwinger's pair production rate by semi-classical tunnelling calculations. The boson and fermion pair production rate by strong static uniform or inhomogeneous electric fields was derived, in terms of instanton tunnelling through potential barriers in the space-dependent gauge, in \\citet{Kim02}.  The Schwinger mechanism for particle production in a strong and uniform electric field for an infinite system was generalized to the case were the strong field is confined between two plates separated by a finite distance by \\citet{Wa88}. The production rates, obtained by solving the Klein-Gordon and Dirac equations in a linear vector potential, can be expressed in an exact analytical form. The numerical evaluations of the production rates have shown large deviations from the Schwinger formula, thus indicating a large finite size effect in particle production. The explicit finite size corrections to the Schwinger formula for the rate of pair production for uniform electric fields confined to a bounded region were calculated, by using the Balian-Bloch expansion of the Green functions, by \\citet{Ma88,Ma89}.  One important astrophysical situation in which electron-positron pair creation could play an extremely important role is the case of the electrosphere of the quark stars \\citep{Us98a,Us98b,Us01,PaUs02}. Quark stars could be formed as a result of a hadron-quark phase transition at high densities and/or temperatures \\citep{It70,Bo71,Wi84}. If the hypothesis of the quark matter is true, then some neutron stars could actually be strange stars, built entirely of strange matter \\citep{Al86,Ha86}. For a review of strange star properties, see \\cite{Ch98}.  Several mechanisms have been proposed for the formation of quark stars. Quark stars are expected to form during the collapse of the core of a massive star after the supernova explosion, as a result of a first or second order phase transition, resulting in deconfined quark matter \\citep{Da}. The proto-neutron star core, or the neutron star core, is a favorable environment for the conversion of ordinary matter into strange quark matter \\citep{ChDa}. Another possibility is that some neutron stars in low-mass X-ray binaries can accrete sufficient mass to undergo a phase transition to become strange stars \\citep{Ch96}. This mechanism has also been proposed as a source of radiation emission for cosmological $\\gamma $-ray bursts \\citep{Ch96}. Some basic properties of strange stars like mass, radius, collapse and nucleation of quark matter in neutron stars cores have been also studied \\citep{Gl95,Gl97,ChHa,HaCh00,HaCh02,Ha04}.  The structure of a realistic strange star is very complicated, but its basic properties can be described as follows \\citep{Al86}. Beta-equilibrated strange quark - star matter consists of an approximately equal mixture of up $u$, down $d$ and strange $s$ quarks, with a slight deficit of the latter. The Fermi gas of $3A$ quarks constitutes a single color-singlet baryon with baryon number $A$. This structure of the quarks leads to a net positive charge inside the star. Since stars in their lowest energy state are supposed to be charge neutral, electrons must balance the net positive quark charge in strange matter stars. The electrons, being bounded to the quark matter by the electromagnetic interaction and not by the strong force, are able to move freely across the quark surface, but clearly cannot move to infinity because of the electrostatic attraction of quarks. For hot stars the electron distribution could extend up to $\\sim 10^{3}$ fm above the quark surface \\citep{Ke95,ChHa03,Us05}. The electron distribution at the surface of the quark star is called electrosphere. The effect of the color-flavor locked phase of strange matter on the electric field in the electrosphere was discussed by \\citet{Us04}.  Photon emissivity is the basic parameter for determining macroscopic properties of stellar type objects. \\citet{Al86} have shown that, because of very high plasma frequency near the strange matter edge, photon emissivity of strange matter is very low. For temperatures $T<<E_{p} $, where $E_{p}\\approx 23$ MeV is the characteristic transverse plasmon cutoff energy, the equilibrium photon emissivity of strange matter is negligibly small, as compared to the blackbody one. The spectrum of equilibrium photons is very hard, with $\\omega >20$ MeV \\citep{Chm91}.  The bremsstrahlung emissivity of the quark matter and of the electrosphere have been estimated recently in \\citet{ChHa03,Ja04,HaCh05}. By taking into account the effect of the interference of amplitudes of nearby interactions in a dense media (the Landau-Pomeranchuk-Migdal effect) and the absorption of the radiation in the external electron layer, the emissivity of the quark matter could be six orders of magnitude lower than the equilibrium blackbody radiation \\citep{ChHa03}. The presence of the electric field strongly influences the radiation spectrum emitted by the electrosphere and strongly modifies the radiation suppression pattern \\citep{HaCh05}. All the radiation properties of the electrons in the electrosphere essentially depend on the value of the electric potential at the quark star surface.  The Coulomb barrier at the quark surface of a hot strange star may also be a powerful source of $e^{+}e^{-}$ pairs, which are created in the extremely strong electric field of the barrier. In the case of the energy loss via the production of $e^{-}e^{+}$ pairs, the energy flux is given by $F_{\\pm }=\\varepsilon _{\\pm }\\dot{n}$, with $\\varepsilon _{\\pm }=m_{e}+T$ and \\begin{equation}\\label{usov} \\dot{n }=\\left( \\frac{9T}{2\\pi \\varepsilon _{F}^{2}}\\right) \\sqrt{\\frac{\\alpha }{\\pi }}\\exp \\left( -2m_{e}/T\\right) n_{e}T^{2}J\\left( \\xi \\right) \\Delta r_E, \\end{equation} where $\\varepsilon _{F}=\\left( \\pi ^{2}n_{e}\\right) ^{1/3}$ is the Fermi energy of the electrons, $\\xi =2\\sqrt{\\alpha /\\pi }\\left( \\varepsilon _{F}/T\\right) $, \\begin{equation}\\label{j} J\\left( \\xi \\right) =\\frac{1}{3}\\frac{ \\xi ^{3}\\ln \\left( 1+2\\xi ^{-1}\\right)} {\\left( 1+0.074\\xi \\right) ^{3}}+\\frac{ \\pi ^{5}}{6}\\frac{ \\xi ^{4}}{\\left( 13.9+\\xi \\right) ^{4}}, \\end{equation} and $\\Delta r_E$ is the thickness of the region with the strong electric field \\citep{Us98a,Us01}.  At surface temperatures of around $% 10^{11}$ K, the luminosity of the outflowing plasma may be of the order $% \\sim 10^{51}$ ergs$^{-1}$ \\citep{Us98a,Us98b}. Moreover, as shown by \\citet{PaUs02}, for about one day for normal quark matter and for up to a hundred years for superconducting quark matter, the thermal luminosity from the star surface, due to both photon emission and $e^{+}e^{-}$ pair production may be orders of magnitude higher than the Eddington limit. Hence, electron-positron pair creation due to the electric field of the electrosphere could be one of the main observational signatures of quark stars.  It is the purpose of the present paper to consider the Schwinger process of pair creation in the electrosphere of the quark stars, by systematically taking into account the main physical characteristics of the environment in which pair production takes place. The electric field in the electrosphere is not a constant, as assumed in the Schwinger model, but it is a rapidly decreasing function of the distance $z$ from the quark star surface. Therefore, in order to realistically describe the pair production process one must consider electron-positron creation in an inhomogeneous electric field. To study the pair production process we adopt the tunnelling approach and the fermion production rate in the electric field of the electrosphere is derived. The production rate is a local quantity, it depends on the distance to the quark star's surface and it is a rapidly decreasing function of $z$. The emissivity and energy flux due to pair creation at the quark star surface is also considered. An important factor which can reduce significantly the pair production rate is the presence of a boundary (the quark star surface). For electric fields perpendicular to the boundary there is a significant reduction in the magnitude of the pair production rate, which is also associated with an important qualitative change in the process, which becomes a periodic function of the distance to the quark star surface.  The present paper is organized as follows. In Section 2 we review the basic properties of the electrosphere of quark stars. The electron-positron rate production in the electric field of the electrosphere is obtained in Section 3. The boundary effects are considered in Section 4. In Section 5 we calculate the electron-positron pair flux of the electrosphere and compare our results with those obtained by \\citet{Us98a,Us98b,Us01}. A brief summary of our results is given in Section 6. ",
        "conclusions": "In the present paper we have re-considered the electron-positron pair emission from the electrosphere of quark stars, as originally proposed by \\citet{Us98a,Us98b}, by pointing out the important role the boundary effects and the inhomogeneity in the distribution of the electric field may play in the pair creation process. At zero temperature, there are no available free energy states in the electron plasma at the strange star's surface. Therefore, at low temperatures $T\\leq T_{cr}\\approx 0.1$ MeV (corresponding to a quark star surface electric potential of $V_q=5$ MeV), the pair production mechanism by the strong electric field of the electrosphere is severely limited by the quantum effects and the exclusion principle specific to the Fermi-Dirac statistics. At high temperatures $T\\geq T_{cr}\\approx 0.1$ MeV, the pair creation rate is controlled by the electric field $E$ and not by the temperature, because such a process is essentially a quantum process. Once the number of available electron states becomes higher than the Schwinger pair production rate, electron-positron pairs can be freely created by the electric field at the surface of strange stars. This happens at a critical temperature $T_{cr}$, which strongly depends on the electrostatic properties of the quark star surface. The critical temperature increases with the increase of the electrostatic potential $V_q$. At high enough temperatures, the pair creation process is almost independent of the temperature and is controlled exclusively by the electric field. On the other hand, the actual thermalized pair creation rate, which, from an astrophysical and observational point of view is the most relevant quantity,  depends on the temperature through the thermalized time scale.The number density of pairs can be accurately evaluated by using the Schwinger formalism and by taking into account the inhomogeneities in the electric field distribution. However, the boundary effects also induce large quantitative and qualitative deviations of the particle production rate from what one deduces with the Schwinger formula and its generalization for the inhomogeneous electric field of the electrosphere.  Due to all these effects, we estimate that at high temperatures the energy flux due to $e^{-}-e^{+}$ pairs production could be lower than in the initial proposal of \\citet{Us98a,Us98b}. However, this flux could still be the main observational signature of a quark star. On the other hand, the presence of a strong magnetic field at the quark star surface may significantly enhance the electron-positron flux.  The possible astrophysical and observational implications of the direct pair production effect will be considered in a future publication."
    },
    "0601/astro-ph0601278_arXiv.txt": {
        "abstract": "A new suite of three dimensional radiative, gravitational hydrodynamical models is used to show that gas giant planets are unlikely to form by the disk instability mechanism at distances of $\\sim$ 100 AU to $\\sim$ 200 AU from young stars. A similar result seems to hold for the core accretion mechanism. These results appear to be consistent with the paucity of detections of gas giant planets on wide orbits by infrared imaging surveys, and also imply that if the object orbiting GQ Lupus is a gas giant planet, it most likely did not form at a separation of $\\sim 100$ AU. Instead, a wide planet around GQ Lup must have undergone a close encounter with a third body that tossed the planet outward to its present distance from its protostar. If it exists, the third body may be detectable by NASA's {\\it Space Interferometry Mission}. ",
        "introduction": "Because radial velocity planet-finding surveys are most sensitive to planets with relatively short-period orbits, the search for gas giant planets on wide orbits (i.e., separations much greater than $\\sim 10$ AU) has been undertaken primarily by direct imaging surveys at infrared wavelengths. These surveys have largely turned up empty-handed: McCarthy \\& Zuckerman (2004) found no evidence for any planets with masses of 5 to 10 $M_J$ (Jupiter mass) at distances of 75 to 300 AU from $\\sim$ 100 stars. Similarly, Masciadri et al. (2005) found nothing in their search of 28 stars for similar-mass planets with separations of at least $\\sim 36$ to $\\sim 65$ AU, as did Lowrance et al. (2005) in their search of 45 nearby stars.  However, recently two ground-based surveys have detected very low mass companions with wide separations. Using the NACO adaptive optics systems on the Very Large Telescope, Chauvin et al. (2004) found a $\\sim 5 M_J$ companion to the $\\sim 25 M_J$ brown dwarf 2M1207, with a separation of $\\sim 60$ AU. Neuhauser et al. (2005) used the same NACO system to detect an object orbiting $\\sim 100$ AU from the $\\sim$ 1 Myr-old classical T Tauri star GQ Lup. The mass of this object appears to lie between $\\sim 1 M_J$ and $\\sim 42 M_J$, with a good chance that its mass is low enough ($< 13 M_J$) for it to be classified as a planet rather than as a brown dwarf.  Core accretion (Mizuno 1980), the generally accepted mechanism of gas giant planet formation, encounters difficulties in forming gas giants at distances much greater than $\\sim 5$ AU (Pollack et al. 1996; Inaba, Wetherill, \\& Ikoma 2003) because of the decreasing surface density of solids available to make the solid cores and the consequent increase in the formation time scale. Forming gas giant planets at distances of $\\sim 100$ AU by core accretion appears to be extremely unlikely.  Observations of the circumstellar disk orbiting the Herbig Ae star AB Aurigae show clear evidence for spiral structures at both near-infrared (Fukagawa et al. 2004) and millimeter-wavelengths (Corder, Eisner, \\& Sargent 2005). Four trailing spiral arms can be resolved, at distances of 200 AU to 450 AU from the central $\\sim$ 4 Myr-old star with a mass of 2.8 $M_\\odot$. These observations are perhaps the first direct evidence for large scale gravitational instability in a circumstellar disk, and suggest that the competing mechanism for gas giant planet formation, disk gravitational instability (Cameron 1978; Boss 1997, 2003), might be able to operate at distances of $\\sim 100$ AU or more. This paper examines the latter possibility, by extending the same disk instability models that suggest gas giant planet formation is possible at distances of $\\sim 10$ AU, to examine the case of much larger disks, $\\sim 100$ AU in size. ",
        "conclusions": "These models have shown that the disk instability mechanism has great difficulty forming gas giant planets {\\it in situ} at distances of $\\sim$ 100 AU to $\\sim$ 200 AU from solar-mass protostars. The models presented here show no clump formation, even in an initially marginally gravitiationally unstable disk. Given that the core accretion mechanism is even more hampered in forming gas giant planets at $\\sim$ 100 AU distances (Pollack et al. 1996; Inaba, Wetherill \\& Ikomoa 2003), it seems clear that gas giant planets do not appear to be able to form on such wide orbits, consistent with the failure of most direct imaging surveys to detect massive gas giant planets (McCarthy \\& Zuckerman 2004; Masciadri et al. 2005; Lowrance et al. 2005).  On the other hand, the detection of the wide, very low mass objects around 2M1207 by Chauvin et al. (2004) and around GQ Lup by Neuhauser et al. (2005) demands an explanation. The 2M1207 system would seem to be best explained as a binary system composed of a brown dwarf and a sub-brown dwarf, with a mass ratio $q = 0.2$, similar to that of many binary star systems. Binary and multiple star systems are generally believed to have formed from the collapse and fragmentation of dense molecular clouds, a process that precedes the planet formation processes considered here.  If the very low mass companion to GQ Lup turns out to have a mass close to the upper limit of $\\sim 42 M_J$, then it also most likely formed as a brown dwarf companion to GQ Lup through collapse and fragmentation into a binary system with $q \\le 0.06$ (assuming a mass of $\\sim 0.7 M_\\odot$ for the T Tauri star GQ Lup). However, if the object has a much lower mass, implying that it did not form by fragmentation, but through the planet formation process in a protoplanetary disk, it is hard to see how such an object could have formed {\\it in situ} at $\\sim$ 100 AU by either core accretion or disk instability. Instead, one would be forced to consider a scenario where the object formed much closer to GQ Lup, perhaps within $\\sim$ 30 AU (e.g., Boss 2003), and then was flung outward on a highly eccentric orbit to $\\sim$ 100 AU by a close encounter with another body in the GQ Lup system (cf., Debes \\& Sigurdsson 2006). Outward scattering of planets to wide orbits is a likely possibility in crowded systems of giant planets (Adams \\& Laughlin 2003).  NASA's {\\it Space Interferometry Mission (SIM)} should be able to detect such a third body orbiting relatively close to GQ Lup, if it exists, settling the question of the origin of the putative planet's wide orbit. GQ Lup has a V magnitude of 14.4 but is highly reddened, allowing {\\it SIM} to perform narrow angle astrometry on it in the R band. Because the third body would be left on a tighter orbit around GQ Lup as a result of the close encounter with the planet now on a wide orbit, the third body would be expected to be on a relatively short period orbit that could be detected during {\\it SIM's} nominal mission lifetime of 5 yrs. At GQ Lup's distance of 140 pc, {\\it SIM} would be able to detect a third body with a mass as low as $\\sim 1 M_J$ orbiting with a semi-major axis of 10 AU or less. As the third body is likely to be significantly more massive than the wide planet, {\\it SIM} should be able to detect this hypothetical object.  I thank Alycia Weinberger for motivating these calculations, Charles Beichman for advice about {\\it SIM's} capabilities, Fred Adams for his excellent comments on the manuscript, and Sandy Keiser for her extraordinary computer systems expertise. This research was supported in part by NASA Planetary Geology and Geophysics grant NNG05GH30G and by NASA Astrobiology Institute grant NCC2-1056. The calculations were performed on the Carnegie Alpha Cluster, the purchase of which was partially supported by NSF Major Research Instrumentation grant MRI-9976645."
    },
    "0601/astro-ph0601715_arXiv.txt": {
        "abstract": "We propose to extend the well-known MUSCL-Hancock scheme for Euler equations to the induction  equation modeling the magnetic field  evolution in kinematic dynamo problems.  The scheme is based on an integral form of the underlying conservation law  which, in our formulation, results in a ``finite-surface'' scheme for the induction equation.  This naturally leads to  the  well-known  ``constrained transport''  method, with additional    continuity  requirement    on    the   magnetic    field representation.   The  second ingredient in the    MUSCL scheme is the predictor step  that ensures second order accuracy   both in space and time.  We explore   specific constraints  that  the  mathematical properties of the induction   equations place on this predictor  step, showing that three possible variants can be considered. We  show  that the  most aggressive formulations (referred to as C-MUSCL and U-MUSCL) reach  the same  level  of accuracy  as the other one (referred to as Runge-Kutta), at a  lower computational cost. More interestingly, these two schemes  are  compatible  with  the  Adaptive Mesh   Refinement  (AMR) framework. It has been implemented in the AMR code RAMSES. It offers a novel  and efficient implementation of  a  second order scheme for the induction equation.  We have tested it by solving two kinematic dynamo problems in   the  low diffusion   limit.     The construction  of this scheme for the  induction equation constitutes a step towards solving the full MHD  set of equations using an extension of our current methodology. ",
        "introduction": "The  extension  of  Godunov-type  conservative  schemes  for Euler equations of fluid dynamics  \\citep{Toro99,Bouchut05} to the system of ideal magnetohydrodynamics (MHD)   has been   a matter of    intensive research, starting from the early 90's. The great variety of different MHD  implementations of the  original  Godunov method, especially in a multidimensional setting, has left  several unexplored paths opened in designing MHD conservative methods.  The most natural approach in adapting finite-volume schemes to the MHD equations is to  define the magnetic field  component at the center of each cell, where  the  traditional hydrodynamical variables  are  also defined. One then takes advantage of decades of experience in the development  of stable and  accurate shock-capturing schemes.  In this case,   the solenoidality  constraint  $\\vnabla\\cdot\\B=0$  has to be enforced using either      a ``divergence cleaning''    step (see   for example \\citeauthor{Brackbill80},          \\citeyear{Brackbill80}          and \\citeauthor{Ryu98},  \\citeyear{Ryu98}),  or various  reformulations of the   MHD   equations        including  additional    divergence-waves \\citep{Powell99} or divergence-damping terms \\citep{Dedner02}. A novel cell-centered  MHD    scheme   has    been   recently   developed   by \\cite{Crockett05} that  combines most of  these ideas  into one single algorithm.  An   alternative approach is   to  use the Constrained Transport  (CT) algorithm for the induction equation, as suggested in the late 60's by \\cite{Yee66},   and later revisited  by  \\cite{EvansHawley88}. In this description, the  magnetic field is defined  at the cell  faces, while other hydrodynamical variables  are defined at  the cell center.  This is often called a ``staggered mesh'' discretization.   As we will see in this paper, CT provides a natural expression of the induction equation in conservative form.  Combining CT  with the Godunov framework to design high-order,  stable schemes is  therefore  a very attractive solution. This  combined approach was  first explored in  the context of the MHD equations   by   \\cite{Balsara99}.    This   method   directly    uses face-centered Godunov  fluxes and averages these on  the  cell edges to estimate the  Electro-Motive Force   (EMF).  \\cite{Toth00}  proposed  an interesting    cell-centered  alternative to   this scheme. More recently, \\cite{Londrillo00, Londrillo04} have revisited the  problem   and shown   that  the  proper   way of  defining the edge-centered EMF is to solve a 2D  Riemann problem at the cell edges. They  have applied this  idea  to design high-order, Runge-Kutta,  ENO schemes.  Finally, \\cite{Gardiner05} have extended Balsara and Spicer scheme to design a more stable and more robust way  of computing the EMF.  The implementation of these  various schemes within the Adaptive  Mesh Refinement framework is another challenging issue. It introduces two main  new  technical   difficulties:  first, proper   fluxes  and  EMF corrections  between different levels  of refinement must be accounted for. Second, when   refining  or de-refining   cells,  divergence-free preserving  interpolation   and     prolongation operators  must    be designed. Both  of these  issues  have recently been discussed in the framework   of the CT  algorithm by  several authors \\citep{Balsara01, Toth02, Li04}.  The purpose of this article is  to present a  novel algorithm based on a high-order   Godunov  implementation of    the  CT algorithm within  a tree-based     Adaptive Mesh  Refinement    (AMR)  code called  RAMSES \\citep{Teyssier02}.  As opposed to   the grid-based  (or  patch-based) original     AMR   designed  introduced   by   \\cite{Berger84}   and \\cite{Berger89}, tree-based  AMR  trigger local grid  refinements on a cell by cell  basis. In this  way, the grid  follows  more closely the geometrical features of  the computed flow, at the  cost of a greater algorithm's complexity. Nevertheless, such tree-based AMR schemes have been implemented  with success by  various authors in the framework of astrophysics and      fluid dynamics  \\citep{Kravtsov97,   Khokhlov98, Teyssier02,  Popinet03} but not yet  in  the MHD  context. On  the other hand, patch-based AMR    algorithms  have been developed  by   several authors              in                 recent                   years \\citep{Balsara01,Kleimann04,Powell99,Samtaney04,Ziegler99}   and  used for MHD applications. The main requirement that tree-based AMR usually place on  the  underlying solver  is the  compactness of the computational stencil: any high  order scheme with a stencil extending to two points, or less,  in each direction  can  easily be coupled  to   an ``octree'' data structure \\citep{Khokhlov98}.  In this paper, our goal is  to solve the  induction equation using the MUSCL  scheme, originally presented  by  \\cite{VanLeer77},  and widely used in the literature for the Euler equations.  This very simple method is second order  accurate in time and space  and has a  compact stencil:  only  2 neighboring cells  in each  direction  (and for each dimension) are necessary  to update the central  cell  solution to the next time step.  This compactness property is of particular importance for our  tree  based AMR approach. It   is also useful for an  efficient parallelization  relying  on domain decomposition.   To our knowledge, this is the first implementation of the MUSCL scheme combined with the Constrained  Transport algorithm that  solves the  induction equation. The key ingredient that ensures second order accuracy is the so-called ``predictor step'', in which the solution  is first advanced by half a time step.  We will consider  a few different computational strategies for this predictor step and discuss their respective merits.  Finally, we will present our overall tree-based AMR scheme.  This paper  is limited to the  induction equation. We  intend to apply the  same approach to  the full  MHD equations  in a future paper. Nevertheless,  it is  interesting  to determine  if  such a  numerical approach can  be applied to  kinematic dynamo problems, for  which the induction equation  alone applies.  The induction  equation is linear, but it can yield  remarkably rich magnetic instabilities corresponding to   exponential   field   growth   and  referred   to   as   ``dynamo instabilities''.   The  description  of  these instabilities,  and  the conditions  under which  they  occur, constitute  an  active field  of research,   with  important  consequences   in  astrophysics   and  in geophysics, since  they account for  the origin of magnetic  fields in the Earth,  planets, stars  and even galaxies.   We will  restrict our attention here to well known dynamo flows and use them to investigate the numerical properties of our scheme.   An important problem in dynamo theory  is related to a subclass of dynamo flows, known as ``fast  dynamos'' which yield exponential field growth with finite growth rates in  the limit of vanishing resistivity.  This is  of  particular importance  for  astrophysical applications.   Fast dynamos,  when  investigated   with  small,  but  finite,  resistivity yield eigenmodes  that  are  very localized  in  space,  and are therefore ideal candidates for an investigation using the AMR scheme.  Dynamo problems have traditionally been studied using spectral methods \\citep{Galloway86,Benchmark01}.  Some recent models have been produced using  finite   differences    \\citep{Archontis03}, finite     volumes \\citep{Harder05} or  finite elements \\citep{Matsui05}.  However, all of these methods rely on  explicit  physical diffusion to  ensure numerical stability.   The  interest of using  CT   within the Godunov framework together with  an AMR approach is twofold.   First,  fast dynamo modes have a very localized  spatial structure (scaling as $Rm^{-1/2}$ where $Rm$  is the  magnetic  Reynolds number).   Adapting the computational grid to the typical geometry  of  these modes  therefore appears as  a very  natural strategy to  minimize  computational cost.   Second, the Godunov  methodology, using   the  CT scheme,  introduces the  minimal amount of numerical dissipation needed to ensure stability. This is an important property when using an AMR approach, for which cells of very different sizes   coexist.  This last  property of  the scheme is then mandatory to allow the use of a coarse grid in regions barely affected by the physical diffusion.  We will  present several tests that demonstrate the efficiency of our tree-based AMR Godunov CT scheme for solving complex dynamo problems: we will first reproduce a simple advection problem of a magnetic loop and then validate the approach on two well known dynamo flows: the Ponomarenko dynamo and a fast ABC dynamo. ",
        "conclusions": "We  have shown that the  Constrained Transport approach for preserving the solenoidal character of the  magnetic field could be combined with a Godunov  method, provided a two-dimensional   Riemann solver can be used. We have  further shown how this  could be combined with  a MUSCL high order scheme. We   considered  three schemes for  the  predictive step, each with its own merits. For a uniform velocity field, these CT schemes are strictly equivalent to well known finite volume schemes on the staggered grid.  This important result provides additional support to the advection properties of the CT framework.  We have  implemented this strategy  on a kinematic dynamo problem, for which only  the induction  equation  needs   to be considered. We have shown that the Godunov framework  allows an efficient AMR treatment of fast dynamos,  by ensuring the  numerical stability  of the scheme  in regions  solved with a coarse  grid (for which the effects of the physical diffusion are vanishing).  The approach introduced here  clearly needs to  be adapted to the full set  of  MHD  equations, for   which  solving the Riemann problem is no longer a trivial task. This important step  raises several  additional difficulties and is the object of a forthcoming paper (\\cite{Fromang06})."
    },
    "0601/astro-ph0601523_arXiv.txt": {
        "abstract": "We analyze V, I and H band HST images and two seasons of $R$-band monitoring data for the gravitationally lensed quasar SDSS0924+0219.  We clearly see that image D is a point-source image of the quasar at the center of its host galaxy.  We can easily track the host galaxy of the quasar close to image D because microlensing has provided a natural coronograph that suppresses the flux of the quasar image by roughly an order of magnitude. We observe low amplitude, uncorrelated variability between the four quasar images due to microlensing, but no correlated variations that could be used to measure a time delay.  Monte Carlo models of the microlensing variability provide estimates of the mean stellar mass in the lens galaxy ($ 0.02 {\\rm M_\\sun} \\lesssim \\avgm \\lesssim 1.0 {\\rm M_{\\sun}}$), the accretion disk size (the disk temperature is $5 \\times10^4$~K at $3.0 \\times 10^{14}\\:{\\rm cm} \\lesssim r_s \\lesssim 1.4 \\times 10^{15} \\: {\\rm cm}$), and the black hole mass ($2.0 \\times 10^7 {\\rm M_{\\sun}} \\lesssim M_{BH}\\, \\eta_{0.1}^{-1/2} \\, (L/L_{E})^{1/2} \\lesssim 3.3 \\times 10^8 {\\rm M_{\\sun}}$), all at 68\\% confidence. The black hole mass estimate based on microlensing is consistent with an estimate of $M_{BH} =(7.3 \\pm 2.4 )\\times 10^7 {\\rm M_{\\sun}}$ from the \\ion{Mg}{2} emission line width. If we extrapolate the best-fitting light curve models into the future, we expect the the flux of images A and B to remain relatively stable and images C and D to brighten.  In particular, we estimate that image D has a roughly 12\\% probability of brightening by a factor of two during the next year and a 45\\% probability of brightening by an order of magnitude over the next decade. ",
        "introduction": "In the standard Cold Dark Matter (CDM) galaxy model, early-type galaxies are composite objects. Their mass is dominated by an extended dark matter halo that surrounds the luminous stars of the visible galaxy; any remaining baryons are left as hot gas \\citep{White1978}.  The halos grow by mergers with other halos, with a small fraction of the accreted smaller halos surviving as satellite halos (``sub-structure\") orbiting in the larger halos, but the mass fraction in these satellites is uncertain (\\citealt{Moore1999}, \\citealt{Klypin1999}, \\citealt{Gao2004}, \\citealt{Taylor2005} and  \\citealt{Zentner2005} ). There is increasing evidence from time delay measurements (e.g. \\citealt{Kochaneketal.2005}) and stellar  dynamical observations (e.g. locally, \\citealt{Romanowsky2003}, and in lenses \\citealt{Treu2002}, \\citealt{Treu2006}), that the density structure of some early-type galaxies on scales of 1--2$R_e$ is heterogeneous, but this needs to be changed from a qualitative assessment to something more quantitative. Thus the detailed balance between stars, dark matter and substructure (luminous or dark) remains a matter of debate.  One approach to addressing these problems is to monitor variability in gravitational lenses. Variability in lenses arises from two sources. Correlated variability between the images due to fluctuations in the source flux allows the measurement of the time delays $\\Delta t$ between the quasar images, which constrain the combination  $\\Delta t \\propto (1-\\kbar)/H_0$ of the surface density near the lensed images $\\kbar=\\langle\\Sigma\\rangle/\\Sigma_c$ and the Hubble constant $H_0$ to lowest order \\citep{Kochanek2002}.\\footnote{The dimensionless surface density $\\kappa$ is the surface density $\\Sigma$ divided by the critical surface density $\\Sigma_c\\equiv c^2 D_{\\rm OS}/4\\pi G D_{\\rm OL}D_{\\rm LS}$, where $D_{\\rm OL}$, $D_{\\rm OS}$ and $D_{\\rm LS}$ are the angular diameter distances between the Observer, Lens and Source.}  By measuring the surface density near the lensed images, we can strongly constrain the radial mass profile of the lens. Uncorrelated variability between the images is a signature of microlensing by stars in the lens galaxy.  Microlensing can constrain the mass distribution because the statistics of the variations depend on the fraction of the local density in stars \\citep{Schechter&Wambsganss2002}. In addition to providing an estimate of the surface density in stars near the lensed images, microlensing can also be used to estimate the mean stellar mass in the lens and to determine the structure of quasar accretion disks.  Understanding microlensing is also required to improve estimates of the substructure mass fraction in radio--quiet lenses, because most other quasar source components are affected by both substructure and microlensing. In order to use time variability in lenses to probe these astrophysical problems, we have undertaken a program to monitor roughly 25 lenses in several optical and near-IR bands.  The first results of the program and a general description of our procedures are presented in \\citet{Kochaneketal.2005}.  In this paper we study the four-image $z_s=1.52$ quasar lens SDSS 0924+0219 \\citep{Inada2003} using $V$, $I$ and $H$-band Hubble Space Telescope ({\\it HST}) observations and the results from two seasons of monitoring the system in the $R$-band. The lens galaxy is a fairly isolated $z_l=0.393$ \\citep{Ofek2005} early-type galaxy.  The most remarkable feature of this lens is that it shows a spectacular flux ratio anomaly between the A and D images.  These two images are merging at a fold caustic with a flux difference of nearly 3 magnitudes when they should have approximately equal fluxes by symmetry \\citep{Keeton2005}. \\citet{Keeton2006} recently measured the flux ratios in the Ly$\\alpha$ line and the adjacent continuum, finding that the anomaly is weaker in the emission line but still present.  This indicates that the anomaly is partly due to microlensing, since the expected size difference between the broad line region and the optical continuum emission region should matter for microlensing but be irrelevant for substructure.  Unfortunately, the continued existence of the anomaly in the emission line means either that the broad line region is not large enough to eliminate the effects of microlensing or that the flux ratio anomaly in SDSS0924+0219 is due the combined effects of microlensing and substructures.  We present the HST data in \\S\\ref{sec:observations} as well as a series of mass models for the system consisting of the observed stellar distribution embedded in a standard dark matter halo.  In \\S\\ref{sec:models} we present the light curves and model them using the Monte Carlo methods of \\citet{Kochanek2004}.  The analysis allows us to estimate the mean stellar mass in the lens galaxy, the size of the quasar accretion disk and the mass of the black hole powering the quasar.  In \\S\\ref{sec:expectations} we present our predictions for the expected variability of this source over the next decade. Finally, we summarize our findings in \\S\\ref{sec:conclusions}. All calculations in this paper assume a flat $\\Lambda$CDM cosmology with $h=0.7$, $\\Omega_{M}=0.3$ and $\\Omega_{\\Lambda}=0.7$. ",
        "conclusions": "\\label{sec:conclusions}  During the course of our monitoring campaign we have observed uncorrelated variability in the four images of SDSS 0924+0219, evidence that microlensing is occurring in this system. Furthermore, our models demonstrate that microlensing is a viable explanation for the system's anomalous flux ratios. This study does not rule out the alternative hypothesis that dark matter substructure contributes to the anomaly, but it does firmly establish the presence of microlensing variability and the ability of microlensing to explain the anomaly. As we continue to monitor SDSS 0924+0219, we expect to eventually measure the time delay, thereby restricting the range of permissible halo models, and to steadily reduce the uncertainties in the estimated mean stellar mass, accretion disk structure and black hole mass.  At some point over the next few years, we should also see dramatic changes in the fluxes of the merging images.  We can also begin to compare microlensing estimates of the structure of quasar accretion disks.  In our original study \\citep{Kochanek2004}, we modeled the significantly more luminous, but very similar redshift, quasar Q2237+0305 ($M_V = -25.8 \\pm 0.5$ versus $M_V = -21.7 \\pm 0.7$ after correcting for magnification).  As we would expect from accretion disk theory, the microlensing analyses indicate that the more luminous quasar has a significantly larger scale ($r_s \\simeq 4.1 \\times 10^{15} \\: {\\rm cm}$ versus $r_s \\simeq 6.9 \\times 10^{14} \\: {\\rm cm}$) and black hole mass ($M_{BH} \\simeq 1.1 \\times 10^9 \\, {\\rm M_{\\sun}}$  versus $1.3 \\times 10^8 \\, {\\rm M_{\\sun}}$).  The next step is to combine the microlensing analyses of many lenses to explore these correlations in detail and to use the wavelength dependence of the microlensing variability to study the structure of individual disks.  This next step should be possible very shortly."
    },
    "0601/astro-ph0601009_arXiv.txt": {
        "abstract": "Numerical simulations of the intergalactic medium have shown that at the present epoch a significant fraction ($40-50\\%$) of the baryonic component should be found in the ($T\\sim 10^6$K) Warm-Hot Intergalactic Medium (WHIM) - with several recent observational lines of evidence indicating the validity of the prediction. We here recompute the evolution of the WHIM with the following major improvements: (1) galactic superwind feedback processes from galaxy/star formation are explicitly included; (2) major metal species (O V to O IX) are computed explicitly in a non-Equilibrium way; (3) mass and spatial dynamic ranges are larger by a factor of 8 and 2, respectively, than in our previous simulations. We find: (1) non-equilibrium calculations produce significantly different results from ionization equilibrium calculations. (2) The abundance of O VI absorption lines based on non-equilibrium simulations with galactic superwinds is in remarkably good agreement with latest observations, implying the validity of our model, while the predicted abundances for O VII and O VIII absorption lines appear to be lower than observed but the observational errorbars are currently very large. The expected abundances for O VI (as well as \\lya), O VII and O VIII absorption systems are in the range $50-100$ per unit redshift at $EW=1$km/s decreasing to $10-20$ per unit redshift at $EW=10$km/s. The number of O VI absorption lines with $EW>100$km/s  is very small, while there are about $1-3$ lines per unit redshift for O VII and O VIII absorption lines at $EW=100$km/s. (3) Emission lines, primarily O VI and \\lya in the UV and O VII and O VIII in the soft X-rays are potentially observable by future missions and they provide complimentary probes of the WHIM in often different domains of the temperature-density-metallicity phase space as well as in different spatial locations; however, sometimes they overlap spatially and in phase space, making joint analyses very useful. The number of emission lines per unit redshift that may be detectable by planned UV and soft X-ray missions are in the order of $0.1-1$. ",
        "introduction": "Cosmological hydrodynamic simulations have strongly suggested that most of the previously ``missing\" baryons may be in a gaseous phase in the temperature range $10^5-10^7$K and at moderate overdensity (Cen \\& Ostriker 1999, hereafter \"CO\"; Dav\\'e 2001), called the warm-hot intergalactic medium (WHIM), with the primary heating process being hydrodynamic shocks from the formation of large-scale structure at scales currently becoming nonlinear. The reality of the WHIM has now been quite convincingly confirmed by a number of observations from HST, FUSE, Chandra and XMM-Newton (Tripp, Savage, \\& Jenkins 2000; Tripp \\& Savage 2000; Oegerle \\etal 2000; Scharf \\etal 2000; Tittley \\& Henriksen 2001; Savage \\etal 2002; Fang \\etal 2002; Nicastro \\etal 2002; Mathur, Weinberg, \\& Chen 2002; Kaastra \\etal 2003; Finoguenov, Briel, \\& Henry 2003; Sembach \\etal 2004; Nicastro \\etal 2005).  In addition to shock heating, feedback processes following star formation in galaxies can heat gas to the same WHIM temperature range. What is lacking theoretically is a satisfactory understanding of the known feedback processes on the WHIM and how the WHIM may be used to understand and calibrate the feedback processes. Another unsettled issue is how the predicted results will change if one has a more accurate, non-equilibrium calculation of the major metal species, such as O VI, O VII and O VIII, since time scales for ionization and recombination are not widely separated from the Hubble time scale. In a companion paper (Cen \\& Ostriker 2005) we have studied the effects of galactic superwinds on the IGM, and in particular on WHIM. Additional improvements include significantly larger dynamic ranges of the simulation, a WMAP normalized cosmological model and an improved radiative transfer treatment. The purpose of this paper is to present additional effects due to non-equilibrium calculations on major observable metal species, such as O VI, O VII and O VIII. Our current work significantly extends previous theoretical works by our group and others (CO; Dav\\'e 2001; Cen \\etal 2001; Fang \\etal 2002; Chen \\etal 2003; Furlanetto \\etal 2004,2005a,b; Yoshikawa \\etal 2003; Ohashi \\etal 2004; Suto \\etal 2004; Fang \\et al. 2005). The outline of this paper is as follows: the simulation details are given in \\S 2; in \\S 3 we give detailed results and discussion and conclusions are presented in \\S 4. ",
        "conclusions": "In Cen \\& Ostriker (2005) we show that our significantly improved simulations confirm previous conclusions based on earlier simulations: {\\it nearly one half of all baryons at the present epoch should be found in the WHIM} - a filamentary network in the temperature range of $10^5-10^7$~K. There, we also presented significant effects due to feedback from star formation.  Here we investigate detailed ionization distributions of oxygen, computed explicitly in a non-Equilibrium way, and present results related to important UV and soft X-ray lines and how their observations may shed light on WHIM and galaxy formation processes. Our finding are: (1) non-equilibrium calculations produce significantly different results from equilibrium calculations, and the differences are complex and difficult to characterize simply. (2) The abundance of O VI absorption lines based on non-equilibrium simulations with galactic superwinds is in remarkably good agreement with latest observations, implying the validity of our model, while the predicted abundances for O VII and O VIII absorption lines appear to be lower than observed but the observational errorbars are currently very large. The expected abundances for O VI (as well as \\lya), O VII and O VIII lines are in the range $50-100$ per unit redshift at $EW=1$km/s decreasing to $10-20$ per unit redshift at $EW=10$km/s. The number of O VI absorption lines with $EW>100$km/s  is very small, while there are about $1-3$ lines per unit redshift for O VII and O VIII absorption lines at $EW=100$km/s. (3) Emission lines, primarily O VI and \\lya in the UV and O VII and O VIII in the soft X-rays are potentially observable by future missions and they provide complimentary probes of the WHIM in often different domains in the temperature-density-metallicity phase space as well as in different spatial locations; however, sometimes they overlap spatially and in phase space, making joint analyses very useful. The number of emission lines per unit redshift that may be detectable by planned UV and soft X-ray missions are in the order of $0.1-1$.   We expect that future missions in UV and X-ray should be able to provide much more accurate characterization of the WHIM through two complimentary approaches: major UV and soft X-ray absorption lines and emission lines. The former will be able to trace out bulk of the mass as well as volume occupied by the WHIM, whereas the latter will be able to probe higher density regions. Together they may allow us to construct a coherent picture of the evolution of the IGM. In addition, detailed useful information may be obtained by investigating relations between galaxies and properties of the WHIM, keeping in mind that GSW play an essential role exchanging mass, metals and energy between them. We stress that proper comparisons between observations and simulations in the vicinity of galaxies where GSW effects originate and are strongest may provide one of the most useful ways to probe star formation and its all important feedback processes."
    },
    "0601/astro-ph0601379_arXiv.txt": {
        "abstract": "Accretion flows with pressure gradients permit the existence of standing waves which may be responsible for observed quasi-periodic oscillations (QPO's) in X-ray binaries.  We present a comprehensive treatment of the linear modes of a hydrodynamic, non-self-gravitating, polytropic slender torus, with arbitrary specific angular momentum distribution, orbiting in an arbitrary axisymmetric spacetime with reflection symmetry.  We discuss the physical nature of the modes, present general analytic expressions and illustrations for those which are low order, and show that they can be excited in numerical simulations of relativistic tori.  The mode oscillation spectrum simplifies dramatically for near Keplerian angular momentum distributions, which appear to be generic in global simulations of the magnetorotational instability.  We discuss our results in light of observations of high frequency QPO's, and point out the existence of a new pair of modes which can be in an approximate 3:2 ratio for arbitrary black hole spins and angular momentum distributions, provided the torus is radiation pressure dominated.  This mode pair consists of the axisymmetric vertical epicyclic mode and the lowest order axisymmetric breathing mode. ",
        "introduction": "Observations of quasiperiodic oscillations (QPO's) in the X-ray light curves of black hole and neutron star X-ray binaries have motivated the theory of relativistic ``diskoseismology'', the study of hydrodynamic oscillation modes of geometrically thin accretion discs (see \\citealt{wag99} and \\citealt{kat01} for reviews).  Standing waves with discrete frequencies can exist in such discs because the radial profiles of general relativistic test particle oscillation frequencies can create finite regions where modes are trapped. Because the disc is considered to be geometrically thin, radial pressure gradients are usually negligible in the equilibrium structure of the accretion flow, which is presumed to consist of fluid rotating on nearly circular geodesics.  However, radial pressure gradients themselves can permit the existence of discrete modes because they produce a disc, or torus, of finite extent. At least in the case of black hole X-ray binaries, QPO's are only observed in the ``hard'' and ``steep power law'' spectral/variability states (e.g. \\citealt{mcc05}), states in which the flow is not solely composed of a standard geometrically thin accretion disc.  It is therefore conceivable that radial pressure gradients may be important in trapping modes, and various groups have recently explored the possibility that accretion tori may explain observed QPO's, e.g. \\citet{gia04} for low frequency QPO's, \\citet{rez03a,klu04,lee04} for high frequency QPO's.  Geometrically thick tori may form in stellar collapse, and modes of dense tori around black holes have also been suggested as a detectable source of gravitational waves (e.g. \\citealt{zan03}).  A complete analysis of the spectrum of modes in tori has not yet been done except in the limiting case of constant specific angular momentum, slender tori in a Newtonian point mass gravitational potential \\citep{bla85}.  In this paper we greatly generalize that early work and present a rather complete analysis of the oscillation modes of non-self-gravitating, polytropic, hydrodynamic tori with small cross-section and arbitrary specific angular momentum distributions.  The torus equilibria are assumed to be stationary and axisymmetric, and orbit in an arbitrary background spacetime which is itself axisymmetric and possesses reflection symmetry. Although limited to slender tori, our work here could be numerically extended to thicker tori by tracking all the mode frequencies.  Armed with a complete understanding of all the mode frequencies and their physical properties, it should be possible to better identify what modes might be responsible for observed QPO's.  An analysis of torus oscillation modes is also useful as a code check for global numerical simulations of accretion flows.  In part because of their finite spatial resolution, these simulations generally use tori rather than geometrically thin discs as initial conditions (e.g. \\citealt{dev03,gam03}).  It is noteworthy that global simulations of non-radiative, magnetized accretion flows often produce small, albeit highly variable, hot inner tori near the innermost stable circular orbit \\citep{haw02}.  This paper is organized as follows.  In section 2 we derive the basic equations describing the equilibria and linear perturbations of relativistic slender tori.  We then make some general remarks about the types of modes and their classification in section 3.  In section 4 we present complete solutions for all the modes of relativistic fluid tori in four special cases where this can be done analytically.  Then in section 5 we show how the low order modes of the general torus can be derived exactly, and also present numerical solutions for higher order modes.  In section 6 we compare these analytic results with numerical simulations of oscillating hydrodynamic tori, and we finally summarize our conclusions in section 7.  Readers who are primarily interested in applications to QPO's should focus on sections 5 and 7.  There we identify a robust new candidate for the modes comprising the observed 3:2 ratio of high frequency QPO's in black hole X-ray binaries. We apply our analytic results for these mode frequencies to GRO~J1655-40 in Figure 9.  We remind the reader that slender hydrodynamic tori are violently unstable to the Papaloizou-Pringle instability \\citep{pap84}.  Perhaps even more important, in the presence of a weak magnetic field, they are unstable to the development of magnetohydrodynamical turbulence due to the magnetorotational instability \\citep{bal98}.  Whether or not the hydrodynamic modes we discuss here can survive these instabilities or even exist is a problem which still needs to be investigated, and we discuss this briefly in section 7. ",
        "conclusions": "We have presented a rather complete discussion of the oscillation modes of non-self-gravitating, relativistic, polytropic tori in an arbitrary axisymmetric spacetime with reflection symmetry, in the limit where the tori are very slender.  Non-slender torus oscillation modes can be explored numerically, and some work on this has already been done by \\citet{zan05} and \\citet{rub05b}.  It would also be possible to examine the behavior of modes for thicker tori analytically using perturbative methods, generalizing the technique employed by \\citet{bla85}.  All the slender tori considered in this paper are hydrodynamically unstable to the nonaxisymmetric, Papaloizou-Pringle instability (PPI) \\citep{pap84}. They would also develop magnetohydrodynamical turbulence in the presence of a weak magnetic field because of the MRI \\citep{bal98}.  Almost all simulations examining oscillations of tori, including those presented in this paper, have been axisymmetric and neglected magnetic fields, so both of these instabilities have not been allowed to develop.  A major outstanding issue is therefore which, if any, of the hydrodynamic torus modes survive these instabilities.  A preliminary investigation of this has been done by \\citet{fra05}, but more work needs to be done.  At least in radially extended tori, the PPI can be suppressed by dynamical accretion through the torus inner edge \\citep{bla87,dev02}. Moreover, the MRI appears to dominate the PPI when magnetic fields are included, and generally leads to a near-Keplerian distribution of specific angular momentum in the flow \\citep{haw02}.  If modes which are essentially hydrodynamic in nature can exist in the presence of MRI turbulence, then the near Keplerian limit is therefore probably of most interest.  The turbulence itself may then determine the mode amplitudes, but all of these questions remain to be investigated.  A systematic understanding of all the types of modes which may be present, as we have done in this paper, should aid such investigations.  Figure \\ref{fig:sigmavskappa} shows that all mode families converge to the epicyclic frequencies or vertical sound waves in the Keplerian limit. If physical discs are indeed nearly Keplerian, this greatly simplifies seismology as only a few distinct frequencies exist. But how close are real discs to the Keplerian limit? Equation (8) can be rewritten in terms of the angular momentum gradient for Keplerian motion, $\\partial l_K/\\partial r$, by differentiating equation (6) and substituting into equation (8). The result is exactly the same as for $\\kappa$ in equation (16), with $\\partial l/\\partial r$, the angular momentum gradient for the fluid torus, replaced by $\\partial l_K/\\partial r$.  The difference of these two quantities becomes \\begin{equation} \\omega_r^2-\\kappa_0^2 = \\frac{E_0^2}{l_0A_0^2} \\left( \\frac{g^{tt}_{,r}-l g^{t\\phi}_{,r}}{g_{rr}} \\right)_0 \\frac{\\partial}{\\partial r}\\left( l_K - l \\right), \\label{eq:diff} \\end{equation} which manifestly goes to zero in the Keplerian limit.  \\citet{dev03} present plots of $l_K$ and $l$ (averaged against density over spherical shells) against radius for MRI driven accretion onto Kerr black holes. In the region of the ``inner torus\" (see their Figure 3), they find $l$ has a similar slope as $l_K$, but is smaller by $\\sim 1-10\\%$, the difference going to zero at the marginally stable orbit. Using equation \\ref{eq:diff}, this implies $(\\omega_r^2- \\kappa_0^2)/\\omega_r^2 \\sim 0.01-0.1$ in our notation. Hence Figure 2 shows that low order modes should be quite accurately described by the Keplerian limit, while the deviation from the Keplerian limit increases with number of nodes. For emission from an optically thick region, flux variations from the photosphere become smaller as more angular nodes are included due to averaging of hot and cold spots. Hence low order modes, which are more nearly in the Keplerian limit, are also likely more observable.  In view of this, the lowest order breathing mode/vertical acoustic wave and the vertical epicyclic mode may be an attractive candidate for a pair of modes which have a near 3:2 commensurability for radiation pressure-dominated ($n=3$) tori, as this persists in the Keplerian limit.  This is in contrast to the plus mode and radial epicyclic mode pair, whose frequency ratio approaches unity in the Keplerian limit.  Identifying the vertical epicyclic mode with the lower of the two observed frequencies also allows the black holes to have any spin parameter, in contrast to the case where the lower of the two frequencies is identified as the radial epicyclic mode.  This is illustrated in Figure \\ref{fig:nuepi1655} for the case of GRO J1655-40.  This source has a mass of $\\sim 6.3$~M$_\\odot$ \\citep{oro03}, and models which identify the 300~Hz QPO as the axisymmetric radial epicyclic mode require $a/M\\gta0.9$ \\citep{rez03a,tor05}.  The same is true for diskoseismology models which identify the 300~Hz QPO with an axisymmetric ``$g$-mode'', which has a frequency below the radial epicyclic frequency within the mode trapping region (e.g. \\citealt{wag01}).  This is in conflict with recent continuum spectral fitting results which suggest a lower black hole spin \\citep{sha05}.  If the 300~Hz QPO is instead the axisymmetric vertical epicyclic mode, any spin is possible.  Many other mode pairs would probably also work, especially if we consider nonaxisymmetric ($m\\ne0$) modes.  However, one still needs to explain why the QPO's occur at the actual observed frequencies, i.e. what singles out a preferred radius for the torus in this source.  More work needs to be done to investigate how torus oscillation modes might manifest themselves in X-ray light curves and power spectra.  Some work has been done based on tracking photon geodesics from the torus to the observer \\citep{bur04,sch05}, but additional work needs to be done on the photon emission mechanism of the tori themselves.  The difficulties that some diagnostics have in identifying torus oscillation modes in numerical simulations, even when they are present in the initial conditions as discussed in section 6 above, should serve as a warning to numericists searching for QPO's in their simulations. The X-ray observations generally have nothing to do with these diagnostics.  It remains unclear whether tori are appropriate models for the flow geometry for the ``steep power law state'' \\citep{mcc05} where high frequency QPO's are seen in black hole X-ray binaries.  \\citet{kub04} suggest that the steep power law state consists of a standard geometrically thin disc sandwiched between a magnetized corona which has sucked up much of the accretion power out of the disc.  Oscillation modes that are analogous to the torus modes with vertical reflection symmetry {\\it might} exist in such a magnetized corona, because the vertical velocities vanish in the equatorial plane, where the high inertia disc is located.  We hope that our comprehensive hydrodynamic mode analysis may act as a foundation for investigation of these and the other questions noted above.  \\begin{figure} \\includegraphics[width=160mm]{fig9.eps} \\caption{Vertical (upper set of curves) and radial (lower set of curves) epicyclic mode frequencies as a function of Boyer-Lindquist coordinate radius for a black hole of mass 6.3~M$_\\odot$, appropriate for GRO J1655-40, and various black hole spins $a/M$ as labelled.  The horizontal dashed lines indicate the frequencies of the observed high frequency QPO's in this source. Models which identify the 300 Hz QPO as the axisymmetric radial epicyclic mode require $a/M\\gta0.9$.  Dotted lines show the lowest order breathing mode frequency for the same black hole spins and both constant specific angular momentum tori and tori in the Keplerian limit.  Models which identify the 300 Hz QPO with the vertical epicyclic mode and the 450 Hz QPO with the breathing mode can work for any spin. \\label{fig:nuepi1655}} \\end{figure}"
    },
    "0601/astro-ph0601186_arXiv.txt": {
        "abstract": "{} { We examined which exo-systems contain moons that may be detected in transit. } { We numerically modeled transit light curves of Earth-like and giant planets that cointain moons with 0.005--0.4 Earth-mass. The orbital parameters were randomly selected, but the entire system fulfilled Hill-stability. } { We conclude that the timing effect is caused by two scenarios: the motion of the planet and the moon around the barycenter. Which one dominates depends on the parameters of the system. Already planned missions (Kepler, COROT) may be able to detect the moon in transiting extrasolar Earth-Moon-like systems with a 20\\% probability. From our sample of 500 free-designed systems, 8 could be detected with the photometric accuracy of 0.1 mmag and a 1 minute sampling, and one contains a stony planet. With ten times better accuracy, 51 detections are expected. All such systems orbit far from the central star, with the orbital periods at least 200 and 10 days for the planet and the moon, while they contain K- and M-dwarf stars. Finally we estimate that a few number of real detections can be expected by the end of the COROT and the Kepler missions. } {} ",
        "introduction": "The photometric transit detection of exoplanets has become an effective tool in the past 6 years. The known systems (9 transits up to now) are summarized in the Schneider catalogue (2005). In the future a few programs are planned to achieve measurements that are as accurate as ca. 0.01--0.1 mmag, allowing the detection of a large sample of Earth-like planets while they are in service (e.g. {\\sc corot} mission, Auvergne et al., 2003, Kepler mission, Borucki et al., 2003, Basri et al., 2005). The power of photometric detection is expected to keep increasing in the future.  Sartoretti  \\& Schneider (1999, SS99), Deeg (2002, D02), and Doyle \\& Deeg (2003) argue that a larger moon than the Earth that is orbiting around an exoplanet causes measurable photometric effects. Although there are only negative ground-based observations (Brown et al., 2001, Charbonneau et al., 2005), it is expected that the planned space missions will be able to discover the exomoons. ",
        "conclusions": "What is the suggested observational strategy in order to find exomoons? Although the required photometric accuracy (0.1--0.15 mmag) seems to be about the highest quality we can expect nowadays, short sampling intervals should be used whenever possible. This will help by increasing the number of systems where the timing effect is less, e.g. when the moon is smaller, closer to the planet, or the mass ratio is smaller. The strategy of the planned missions is concordant with this.  We note that the observing strategy of Kepler (Basri et al., 2005) may allow the detection of exomoons, but may also lead to the rejection of those observations. The transit pipeline of Kepler includes at least 3 transits with timings that are consistent with a periodically revolving planet, within the observational errors. Any inconsistent transits will be rejected. If an exomoon has an observable timing effect such that timings occur {\\it discordantly} under the assumption of strict monoperiodicity, Kepler may not recognize the light variation as a transit. Therefore we suggest including some timing variation in the acceptance criteria.  The possible spectroscopic confirmation of a moon could be either the observation of the perturbations in the radial velocity of the star or the detection of the moon within the Rossiter-McLaughlin effect, which e.g. Bouchy et al. (2005) observed for a transiting planet. Unfortunately both mean only slight variations in the fine structures, which is why there is really little hope for their success.  Based on the presented calculations, one may estimate the magnitude of the expected real detections with Kepler. The total number of Earth-sized planets to be discovered is a few hundred during the entire mission. If only 5 percent of them have a moon similar to our Moon, and only every fifth moon are really found, we should get a few positively detected exomoons. If we take the {\\sc corot} mission and the forthcoming missions into account, we may have a few dozen known exomoons by the end of the following decade."
    },
    "0601/astro-ph0601465_arXiv.txt": {
        "abstract": "Late-time $BVRIJHK$ photometry of the peculiar Type Ic SN 2002ap, taken between 2002 June 12 and 2003 August 29 with the MAGNUM telescope, is presented. The light curve decline rate is derived in each band and the color evolution is studied through comparison with nebular spectra and with SN 1998bw. Using the photometry, the $OIR$ bolometric light curve is built, extending from before light maximum to day 580 after explosion. The light curve has a late-time shape strikingly similar to that of the hypernova SN 1998bw. The decline rate changes from 0.018 mag/day between day 130 and 230 to 0.014 mag/day between day 270 and 580. To reproduce the late-time light curve, a dense core must be added to the 1-D hypernova model that best fits the early-time observations, bringing the ejecta mass from 2.5 $M_\\sun$ to 3 $M_\\sun$ without much change in the kinetic energy, which is $4\\times 10^{51}$ ergs. This is similar to the case of other hypernovae and suggests asymmetry. A large $H$-band bump developed in the spectral energy distribution after $\\sim$ day 300, probably caused by strong [Si I] 1.646$\\mu$m and 1.608$\\mu$m emissions. The near-infrared flux contribution increased simultaneously from $<30\\%$ to $>50\\%$ at day 580. The near-infrared light curves were compared with those of other Type Ib/c supernovae, among which SN 1983I seems similar to SN 2002ap both in the near-infrared and in the optical. ",
        "introduction": "There is growing interest in Type Ib/c supernovae (SNe) following the discovery of a subclass that display very broad-lined spectra, indicating the existence of ejecta expanding with velocity $>0.1c$. Models of the optical spectra and the light curves (LCs) of this high-velocity subclass, also called hypernovae (HNe, e.g., \\citealt{nom04}), have concluded that these are the explosions of massive C+O stars, producing up to 10 times the kinetic energy of normal core-collapse SNe (e.g., \\citealt{iwa98,iwa00,maz03,den05}). It is unclear why some HNe are apparently associated with GRBs, i.e., SN 1998bw and GRB 980425, \\citep{gal98}, SN 2003dh and GRB 030329 \\citep{sta03}, and SN 2003lw and GRB 031203 \\citep{mal03}, while others are not.  The Type Ic SN 2002ap, discovered in M74 \\citep{nak02}, is one of the nearest SNe in recent years. Its host galaxy has a distance modulus of only $29.5_{-0.2}^{+0.1}$ (\\citealt{sha96,soh96}; see also \\citealt{hen05} for the latest review). \\citet{tak02} measured the total extinction toward the SN with a high-resolution spectrum and concluded that $E(B-V)=0.09\\pm 0.01$. We adopt these values in this paper.  Early-time optical spectroscopy and photometry was published by \\citet{maz02}, \\citet{gal02}, \\citet{kin02}, \\citet{fol03}, \\citet{pad03a}, \\citet{vin04}, and others. Most authors emphasized the spectral similarity to the broad-lined SN 1998bw and SN 1997ef. The early-time light curve (LC), was much broader than that of the normal Type Ic SN 1994I, but was significantly fainter than that of SN 1998bw and peaked earlier.  \\citet{yos03} (hereafter Paper I) presented early-time near-infrared (NIR) and optical photometry, obtained with the MAGNUM telescope. Using the $UBVRIJHK$ photometry, they built a well-sampled $OIR$ bolometric LC for SN 2002ap. NIR photometry has also been reported by \\citet{nis02} and \\citet{mat02}. \\citet{ger04} described the evolution of NIR spectra, which were dominated by lines of intermediate-mass elements (see also \\citealt{dan02}).  \\citet{maz02} modelled the early-time optical spectra and bolometric LC using 1-D codes. They concluded that SN 2002ap was the energetic explosion of a C+O star, having evolved from a star of 20-25 $M_\\sun$ on the main sequence, and that it ejected $\\sim$ 2.5 $M_\\sun$ material with a kinetic energy of $\\sim$ $4\\times10^{51}$ ergs, including $\\sim$ 0.07 $M_\\sun$ $^{56}$Ni. This places SN 2002ap at the low-mass low-energy end of hypernovae. They also constrained the explosion date to MJD 52300.0 $\\pm$ 0.5 days. We use the explosion date as the reference point for any SN epoch in this paper. Based on observed limits in pre-discovery images, \\citet{sma02} argued that the progenitor was either a single W-R star of main-sequence mass $<40 M_\\sun$ or more likely part of an interacting binary.  SN 2002ap was also monitored in the radio \\citep{ber02,sod05}, X-rays \\citep{sut03,sor04}, and UV \\citep{sor02}. \\citet{ber02} estimated quite a small energy in any relativistic material from the faintness of the radio. They addressed the material with velocities larger than the shock velocity ($\\sim0.3c$). Therefore, their results do not contradict the high energy obtained in the 1-D model for the bulk SN ejecta by \\citet{maz02}, whose ejecta velocity distribution has a cut-off at 70,000 km~s$^{-1}$. As a possible confirmation of this result, the shock velocity derived by \\citet{bjo04} from models of the radio and X-ray observations was also $\\sim$ 70,000 km~s$^{-1}$.  SN 2002ap was not apparently associated with a GRB \\citep{hur02}. However, \\citet{kaw02} discovered a polarized spectrum component that resembled the total flux spectrum but redshifted by $z\\sim0.3$ (see also \\citealt{leo02,wan03}). They proposed that this component is due to scattering from jet material. \\citet{tot03} further suggested that the jet may be too baryon-contaminated to trigger a GRB and that it was expanding almost freely and hence was radio quiet. \\citet{tom04} reproduced a bump in the LC of the XRF 030723 afterglow with a SN 2002ap-like LC.  SN 2002ap was monitored into the late nebular phase. \\citet{fol03} published optical photometry covering from 2002 June to December and spectroscopy continuing to 2003 February. The spectra show unusually strong and sharp-peaked [\\ion{O}{1}] and \\ion{Mg}{1}] emission lines (see also \\citealt{leo02}). Late-time spectroscopy was also obtained by other observers (Y. Qiu; K. Kawabata; S. B. Pandey; private communications), and late-time photometry by \\citet{vin04} and \\citet{dor03}.  It is also important to have intensive late-time NIR observations of Type Ib/c SNe. As demonstrated in the case of SN 1998bw \\citep{pat01,sol02}, the NIR contribution to the total flux is quite large, in particular at late times. Late-time NIR observations may also reveal molecule and dust formation, if any exists (e.g., \\citealt{ger04}). However, few Type Ib/c SNe have been observed at late phases in the NIR. For SN 2002ap, the only late-time NIR observation published so far is a spectrum taken in 2002 August \\citep{ger04}.  In this paper, we report on $BVRIJHK$ imaging photometry obtained for SN 2002ap between 2002 June 12 and 2003 August 29, completing the 1.5 years follow-up project with the MAGNUM telescope (see also Paper I). We describe the observations and data reduction and present our multicolor LCs in \\S 2. An $OIR$ bolometric LC is then constructed using the photometry. In \\S 3, we discuss the bolometric LC, analyze the evolution of the spectral energy distribution (SED), compare the late-time NIR LCs with those of other Type Ib/c SNe, and present 1-D models for the bolometric LC. ",
        "conclusions": "\\subsection{Bolometric Light Curve}  We built the $OIR$ bolometric LC integrating the $BVRIJHK$ broadband flux (see Paper I for details). A reddening of $E(B-V)=0.09\\pm0.01$ was corrected for \\citep{tak02}, and a distance modulus of 29.5 was adopted (but see \\citealt{hen05}).  The LC is shown in Figure 4, extending that in Paper I to $\\sim 580$ days since explosion. It consists of three independent data sets, which agree with one another. Most of the data ({\\em filled circles}), tabulated in Table 3 ({\\em late time}) and Paper I ({\\em early time}), were made from the MAGNUM photometry. We interpolated and extrapolated the observations to estimate the the $K$ magnitudes between day 280 and 340 and the $J$ and $I$ magnitudes at day 520. For epochs $<11$ days, the observations of \\citet{mat02}, \\citet{gal02}, and \\citet{coo02} were used for interpolation. Bolometric magnitudes at day 407 and 578 ({\\em triangles}) were converted from $H$ magnitudes using approximate bolometric corrections. The other two data sets were based on the \\citet{fol03} ({\\em open circles}) and \\citet{pad03a,pad03b} ({\\em crosses}) optical photometry, respectively, which was combined with the \\citet{nis02} and our $JHK$ photometry.  The bolometric LC at late time has a shape similar to that of SN 1998bw (Figure 4; {\\em squares}), another hypernova, although the latter was much brighter. The SN 1998bw LC in Figure 4 is shifted down by 1.75 mag to match the peak magnitude of SN 2002ap. Both LCs show a similar late-time slowing-down. For SN 2002ap, the decline rate changed from $\\sim 0.018$ mag/day between day 130 and 230 to $\\sim 0.014$ mag/day between day 270 and 580. For SN 1998bw, the decline rate was $\\sim 0.019$ mag/day between day 70 and 220 and $\\sim 0.014$ mag/day between day 320 and 540. For comparison, the late-time bolometric LC of another hypernova, SN 1997ef, followed the decay rate of $^{56}$Co, i.e., $\\sim 0.01$ mag/day \\citep{maz04}. On the other hand, SN 2002ap reached its light maximum $\\sim 5$ days earlier than SN 1998bw did, although their peak widths are actually comparable (Figure 4; {\\em inset}).  \\subsection{Spectral Energy Distribution}  We show the SED of SN 2002ap at 5 typical late epochs in Figure 5, where the monochromatic fluxes converted from the MAGNUM $BVRIJHK$ photometry are connected using spline-fitting curves. The observed upper limits of the $K$-band flux at day 393 and 520 are marked by {\\em arrows}. Zero flux was assumed at both the the blue edge of the $B$ band and the red edge of the $K$ band. We used such SEDs to obtain the $OIR$ bolometric magnitudes.  The SEDs are dominated by the flux in the $R$ and $I$ bands before about day 340-390, which was mainly due to strong [\\ion{O}{1}], \\ion{Ca}{2}], and \\ion{Ca}{2} emissions \\citep{fol03}. The $B$ bump at the three intermediate epochs reflects the contribution of \\ion{Mg}{1}] $\\lambda$4571. After about day 340-390, the $V$ flux, attributed to [\\ion{Fe}{2}] multiplets, increased relative to that in other optical bands, while the $I$-band-dominated \\ion{Ca}{2} emissions died away.  The NIR flux contribution rose rapidly after $\\sim$ day 300 (Figure 6). This can been seen in Figure 5 as a big $H$ bump in the SED and also in Figure 2 as a significant flattening of the $H$ and $J$ LCs. To explain this, we examined the synthesized NIR spectra of the models that reproduce the late-time optical spectra of SN 2002ap (P. A. Mazzali et al. 2005, in preparation). The most likely candidates are the strong [\\ion{Si}{1}] 1.646 $\\mu$m and 1.608 $\\mu$m lines, while the 1.099 $\\mu$m line may account for the concurrent $J$-band flux increase. The lines were already developed in the day $\\sim$ 200 NIR spectrum of \\citet{ger04}, although they were still not as strong as a 1.5 $\\mu$m feature, possibly \\ion{Mg}{1} emission. We note in passing that the appearance of the CO first-overtone band in that spectrum suggested by those authors seems to coincide with the $K$-band LC steepening seen in Figure 2.  The late-time $H$-band spectra of SN 1983N, taken by \\citet{gra86} and of very low resolution, showed a strong 1.65 $\\mu$m feature, which was explained by those authors as [\\ion{Fe}{2}] lines but was soon, together with an adjacent feature, re-identified as [\\ion{Si}{1}] lines by \\citet{oli87}. These [\\ion{Si}{1}] NIR features can also be found in the synthetic late-time spectra of Type Ib SNe computed by \\citet{fra89}.  \\subsection{Late-time NIR Light Curves of Type Ib/c Supernovae}  In Figure 7 we compare the late-time $JHK$ LCs of SN 2002ap with those of other Type Ib/c SNe (see Table 4). SN 1998bw and SN 1984L were 1.5--2.5 mag brighter than all other SNe, indicating a large mass of $^{56}$Ni ejected. The NIR LCs of SN 1984L were relatively slow, as was its late-time optical LC which was modelled using very massive ejecta but with normal kinetic energy \\citep{swa91,bar93}. The few points of SN 1998bw,when shifted down by 1.8 mag ($J$), 1.5 mag ($H$), and 1.7 mag ($K$), respectively, fall on the LCs of SN 2002ap. The LCs of SN 1983N have later and broader peaks than those of SN 2002ap, but they become similar between in the $J$-band and $H$-band day $\\sim$ 200 and 350. On the other hand, the $JHK$ data of SN 1983I and 1982R seem to bridge the LC gap in SN 2002ap data between day 50 and day 140 (solar conjunction).  Figure 7 contradicts the picture that SNe 1983I, 1983N, and 1984L have similar $JHK$ LCs \\citep{eli85}. Type Ib/c SNe were first established as a new class by those authors, who showed that these SNe have similar $JHK$ LCs which are different from those of Type Ia SNe [and also by \\citet{whe85} through an analysis of the optical spectra]. They artificially shifted the data of SN 1983I and SN 1984L in both phase and brightness in order to match the NIR LCs of SN 1983N. Their SN 1983I data were assigned epochs 20 days later than ours, and SN 1984L 14 days earlier. Our epoch estimates are more reasonable because we based them on optical LCs and spectra which are better observed and understood than the NIR data.  \\subsection{Was SN 1983I similar to SN 2002ap?}  Inspecting by eye the only published spectrum of SN 1983I \\citep{whe87}, we noticed that it resembles those of SN 2002ap between 2002 February 16 and 22 \\citep{kin02}. The similarity can be seen not only in the overall shape and strongest line features, but also in relatively weak ones, like those near $4900$\\AA, $5900$\\AA, and $6200$\\AA. The spectrum was taken on 1983 May 17, 5 days after discovery. We estimate the epoch of that spectrum as $22\\pm3$ days since explosion, through this comparison, and hence derive an explosion date of 1983 April 25.  The LCs of SN 1983I seem to match those of SN 2002ap not only in the NIR (Figure 7) but also in the optical (Figure 8), if our estimates of the explosion date and the distance modulus are correct. We adopt 31.0 as the distance modulus to the host galaxy of SN 1983I, NGC 4051 in the Ursa Major cluster \\citep{tsv85}. This value is also close to the average of the distance modulus of the Ursa Major cluster (31.4, \\citealt{tul00}), and that corresponding to its radial velocity corrected for the Local Group infall onto Virgo (30.5, LEDA \\footnote{http://leda.univ-lyon1.fr/}; \\citealt{pat97})  Therefore, SN 1983I could be a precedent of the unusual SN 2002ap, which lies at the low-mass low-energy end of hypernovae \\citep{maz02}. However, its data are insufficient for us to make a solid conclusion.  \\subsection{Modelling the Late-time Light Curve}  The late-time bolometric LC of SN 2002ap declines more slowly than that of the 1-dimensional (1-D) model of \\citet{maz02} that best reproduced the early-time spectra and LC (see Figure 9; {\\em dotted line}). This is expected, since the LC follows that of SN 1998bw closely in terms of the decline rate, as shown in Figure 4. The best-fitting 1-D model for the early-time spectra and LC of SN 1998bw also fails to explain its late-time spectrum and LC \\citep{nak01,maz01}. To mimic the outcome of 2-dimensional jet-induced explosion calculations \\citep{mae02}, which applies to SN 1998bw and may also apply to SN 2002ap, \\citet{mae03} introduced a dense core. That core absorbs $\\gamma$-rays efficiently at late time. Using this structure, the slow LC decline of hypernovae was reproduced by those authors using a Monte-Carlo radiative transfer code but with simplified physics.  We tested the dense-core scenario on SN 2002ap using our sophisticated 1-D SN radiation hydrodynamical and $\\gamma$-ray transfer code \\citep{iwa00,maz02,den05}. The best-fitting model ({\\em solid line}) has $\\sim 0.6 M_\\sun$ ejecta below 3,000 km~s$^{-1}$, compared with 0.1 $M_\\sun$ in the model without a dense core, and $\\sim 0.01 M_\\sun$ low-velocity $^{56}$Ni. Compared with \\citet{maz02}, who modelled only the early-time observations, the total ejecta mass has increased from 2.5 $M_\\sun$ to 3 $M_\\sun$, but with little change in the total $^{56}$Ni mass (0.08 $M_\\sun$) and the kinetic energy ($4\\times 10^{51}$ ergs). A dense ejecta core is also required in 1-D late-time spectrum models (P. A. Mazzali et al. 2006, in preparation) in order to explain the observed sharp line cores of [\\ion{O}{1}] and \\ion{Mg}{1}] \\citep{leo02,fol03}.  A dense core cannot be formed in 1-D explosion simulations for hypernovae \\citep{nak01}, but it is a natural product of 2-D jet-induced explosions \\citep{mae02}. This strongly indicates asymmetry in the SN 2002ap explosion, an intrinsic feature also shared by other hypernovae. Evidence of asymmetry in SN 2002ap can also be found in spectropolarimetry \\citep{kaw02,leo02,wan03}, which shows an intrinsic continuum polarization varying around 0.5\\%, corresponding to an asphericity of $\\sim$ 10\\% for the bulk of the ejecta."
    },
    "0601/hep-ph0601016_arXiv.txt": {
        "abstract": "We review the main proposals of particle physics for the composition of the cold dark matter in the universe. Strong axion contribution to cold dark matter is not favored if the Peccei-Quinn field emerges with non-zero value at the end of inflation and the inflationary scale is superheavy since, under these circumstances, it leads to unacceptably large isocurvature perturbations. The lightest neutralino is the most popular candidate constituent of cold dark matter. Its relic abundance in the constrained minimal supersymmetric standard model can be reduced to acceptable values by pole annihilation of neutralinos or neutralino-stau coannihilation. Axinos can also contribute to cold dark matter provided that the reheat temperature is adequately low. Gravitinos can constitute the cold dark matter only in limited regions of the parameter space. We present a supersymmetric grand unified model leading to violation of Yukawa unification and, thus, allowing an acceptable $b$-quark mass within the constrained minimal supersymmetric standard model with $\\mu>0$. The model possesses a wide range of parameters consistent with the data on the cold dark matter abundance as well as other phenomenological constraints. Also, it leads to a new version of shifted hybrid inflation. ",
        "introduction": "\\label{sec:intro}  The recent measurements of the Wilkinson microwave anisotropy probe (WMAP) satellite \\cite{wmap} on the cosmic microwave background radiation (CMBR) have shown that the matter abundance in the universe is $\\Omega_mh^2=0.135^{+0.008}_{-0.009}$, where $\\Omega_i=\\rho_i/\\rho_c$ with $\\rho_i$ being the energy density of the $i$-th species and $\\rho_c$ the critical energy density of the universe and $h$ is the present value of the Hubble parameter in units of $100~{\\rm km}~{\\rm sec}^{-1}~{\\rm Mpc}^{-1}$. The baryon abundance is also found by these measurements to be $\\Omega_bh^2=0.0224\\pm 0.0009$. Consequently, the cold dark matter (CDM) abundance in the universe is $\\Omega_{\\rm CDM}h^2= 0.1126^{+0.00805}_{-0.00904}$. The $95\\%$ confidence level (c.l.) range of this quantity is then, roughly, $0.095 \\stackrel{<}{_{\\sim}}\\Omega_{\\rm CDM}h^2 \\stackrel{<}{_{\\sim}}0.13$. Taking $h\\simeq 0.72$, which is its best-fit value from the Hubble space telescope (HST) \\cite{hst}, and assuming that the total energy density of the universe is very close to its critical energy density (i.e. $\\Omega_{\\rm tot} \\simeq 1$), as implied by WMAP, we conclude that about $22\\%$ of the energy density of the present universe consists of CDM.  \\par The question then is, what the nature, origin, and composition of this important component of our universe is. Particle physics provides us with a number of candidate particles out of which CDM can be made. These particles appear naturally in various particle physics frameworks for reasons completely independent from CDM considerations and are, certainly, not invented for the sole purpose of explaining the presence of CDM in the universe.  \\par The basic properties that such a candidate particle must satisfy are the following: ($i$) it must be stable or very long-lived, which can be achieved by an appropriate symmetry, ($ii$) it should be electrically and color neutral, as implied by astrophysical constraints on exotic relics (like anomalous nuclei), but can be interacting weakly, and ($iii$) it has to be non-relativistic, which is usually guaranteed by assuming that it is adequately massive, although even very light particles such as axions can be non-relativistic for different reasons. So, what we need as constituent of CDM is a weakly interacting massive particle. There are several possibilities, but we will concentrate here on the major particle physics candidates which are the axion, the lightest neutralino, the axino, and the gravitino (for other candidates, see e.g. Ref.~\\cite{other}). Note that the last three particles exist only in supersymmetric (SUSY) theories.  \\par In Sec.~\\ref{sec:axion}, we examine the possibility that the axions are constituents of CDM. Sec.~\\ref{sec:mssm} is devoted to outlining the salient features of the minimal supersymmetric standard model (MSSM), which will be used as a basic frame for discussing SUSY CDM. In Sec.~\\ref{sec:relic}, we summarize the calculation of the relic abundance of the lightest neutralino, which is normally the lightest supersymmetric particle (LSP), and investigate the circumstances under which it can account for the CDM in the universe. In Secs.~\\ref{sec:axino} and \\ref{sec:gravitino}, we discuss, respectively, axinos and gravitinos as constituents of CDM. In Sec.~\\ref{sec:quasi}, we present a SUSY grand unified theory (GUT) model which solves the bottom-quark mass problem by naturally and modestly violating the exact unification of the third generation Yukawa couplings. We study the parameter space of the model which is allowed by neutralino dark matter considerations as well as some other phenomenological constraints. Finally, in Sec.~\\ref{sec:concl}, we summarize our conclusions. ",
        "conclusions": "\\label{sec:concl}  \\par We showed that particle physics provides us with a number of candidate particles out of which the CDM of the universe can be made. These particles are not invented solely for explaining the CDM, but they are naturally there in various particle physics models. We discussed in some detail the major candidates which are the axion, the lightest neutralino, the axino, and the gravitino. The last three particles exist only in SUSY theories and can be stable provided that they are the LSP.  \\par The axion is a pseudo Nambu-Goldstone boson associated with the spontaneous breaking of a PQ symmetry. This is a global anomalous ${\\rm U}(1)$ symmetry invoked to solve the strong $CP$ problem. It is, actually, the most natural solution to this problem which is available at present. The axions are extremely light particles and are generated at the QCD phase transition carrying zero momentum. We argued that these particles can easily provide the CDM in the universe. However, if the PQ field emerges with non-zero value at the end of inflation, they lead to isocurvature perturbations, which, for superheavy inflationary scales, are too strong to be compatible with the recent results of the WMAP satellite on the CMBR anisotropies.  \\par The most popular CDM candidate is, certainly, the lightest neutralino which is present in all SUSY models and can be the LSP for a wide range of parameters. We considered it within the simplest SUSY framework which is the MSSM whose salient properties were summarized. We used exclusively the constrained version of MSSM which is known as CMSSM and is based on universal boundary conditions. In this case, the lightest neutralino is an almost pure bino, whereas the NLSP is the lightest stau. We sketched the calculation of the neutralino relic abundance in the universe paying particular attention not only to the neutralino pair annihilations, but to the neutralino-stau coannihilations too. It is very important for the accuracy of the calculation to treat poles and final-state thresholds properly and include the one-loop QCD corrections to the Higgs boson decay widths and the fermion masses. We find that two effects help us reduce the neutralino relic abundance and satisfy the WMAP constraint on CDM: the resonantly enhanced neutralino pair annihilation via an $A$-pole exchange in the s-channel, which appears in the large $\\tan\\beta$ regime, and the strong neutralino-stau coannihilation, which is achieved when these particles are almost degenerate in mass.  \\par The axino, which is the SUSY partner of the axion, can also be the LSP in many cases since its mass is a strongly model-dependent parameter in the CMSSM. It is produced thermally by 2-body scattering or decay processes in the thermal bath, or non-thermally by the decay of sparticles which are already frozen out of thermal equilibrium. For small axino masses, TP is more important yielding a very narrow favored region in the parameter space. For larger axino masses, however, NTP is more efficient and the favored region in the parameter space becomes considerably wider. One finds that, in the case of the CMSSM, almost any point on the $m_0-M_{1/2}$ plane can be allowed by axino CDM considerations. The required reheat temperatures though are quite small ($\\stackrel{<}{_{\\sim}}{\\rm few}\\times 100~{\\rm GeV}$).  \\par The mass of the gravitino is a practically free parameter in the CMSSM. So, the gravitino can easily be the LSP and, in principle, contribute to the CDM of the universe. It is produced thermally by 2-body scattering processes in the thermal bath as well as non-thermally by the decay of the NLSP, which can be either the neutralino or the stau. In contrast to the axino case, however, the NLSP can now have quite a long lifetime. The electromagnetic showers resulting from the NLSP decay can destroy the successful predictions of BBN. So, we obtain strong constraints which allow only very limited regions of the parameter space of the CMSSM. As it turns out, NTP in these regions is not efficient enough to account for CDM. We can, however, make these regions cosmologically favored by raising $T_{\\rm r}$ to enhance TP of gravitinos.  \\par We studied the CMSSM with $\\mu>0$ and $A_0=0$ applying a YQUC which originates from a PS SUSY GUT model. This condition yields an adequate deviation from YU which allows an acceptable $m_b(M_Z)$. We, also, imposed the constraints from the CDM in the universe, $b\\rightarrow s\\gamma$, $\\delta\\alpha_\\mu$ and $m_h$. We found that there exists a wide and natural range of CMSSM parameters which is consistent with all the above constraints. The parameter $\\tan\\beta$ ranges between about 58 and 59 and the asymptotic splitting between the bottom (or tau) and the top Yukawa coupling constants varies in the range $25-29\\%$ for central values of $m_b(M_Z)$ and $\\alpha_s(M_Z)$. The predicted LSP mass can be as low as about $176~{\\rm GeV}$. Moreover, the model resolves the $\\mu$ problem of MSSM, predicts stable proton, generates the baryon asymmetry of the universe via primordial leptogenesis, and gives rise to a new version of shifted hybrid inflation which is based solely on renormalizable interactions."
    },
    "0601/astro-ph0601629_arXiv.txt": {
        "abstract": "We derive the spectral energy distribution (SED) of the nucleus of the Seyfert galaxy \\object{NGC~4565}. Despite its classification as a Seyfert~2, the nuclear source is substantially unabsorbed. The absorption we find from Chandra data ($N_H=2.5 \\times 10^{21}$ cm$^{-2}$) is consistent with that produced by material in the galactic disk of the host galaxy. HST images show a nuclear unresolved source in all of the available observations, from the near-IR H band to the optical U band. The SED is completely different from that of Seyfert galaxies and QSO, as it appears basically ``flat'' in the IR-optical region, with a small drop-off in the U-band. The location of the object in diagnostic planes for low luminosity AGNs excludes a jet origin for the optical nucleus, and its extremely low Eddington ratio $L_o/L_{Edd}$ indicates that the radiation we observe is most likely produced in a radiatively inefficient accretion flow (RIAF). This would make NGC~4565 the first AGN in which an ADAF-like process is identified in the optical. We find that the relatively high [OIII] flux observed from the ground cannot be all produced in the nucleus. Therefore, an extended NLR must exist in this object. This may be interpreted in the framework of two different scenarios: i) the radiation from ADAFs is sufficient to give rise to high ionization emission-line regions through photoionization, or ii) the nuclear source has recently ``turned-off'', switching from a high-efficiency accretion regime to the present low-efficiency state. ",
        "introduction": "Low  luminosity active  galactic  nuclei (LLAGN)  are  believed to  be powered  by accretion of  matter onto  the central  supermassive black hole, similarly  to powerful  AGN. In a  large fraction of  LLAGN, the central  black hole  is  as  massive as  in  powerful distant  quasars ($M_{BH}  \\sim 10^{8}- 10^{9}  M_\\sun$), thus  their very  low nuclear luminosity  implies  that accretion  occurs  with  very low  radiative efficiency      \\citep[or     at      very     low      rates;     see e.g.][]{ho04coz,papllagn}. If so, the physics of the accretion process may be different from  the ``standard'' optically thick, geometrically thin  accretion disks.   Starting from  the ``ion-supported  tori'' of \\citet{rees82}, a number of  theoretical models have been developed to describe  such radiatively  inefficient  accretion flows  \\citep[RIAF, e.g.   advection-dominated accretion flows,  ADAF, advection-dominated inflow-outflow solutions, ADIOS, convection-dominated accretion flows, CDAF][]{narayanyi94,quataert99,abramowicz02}. But  because of the very low radiation they emit at  all wavelengths, these objects (if they at all  exist) are  very difficult  to  be observed.   Recently, the  AGN nature of optical nuclear components seen in HST images of a sample of LLAGN have been  unambiguously established.  \\citet{Maoz05} have shown that among  a sample of 17 LLAGN,  15 of them show  variability over a timescale  of  a  few  months, which  demonstrate  their  non--stellar origin. However, it  is still unclear whether the  radiation is from a jet or from the accretion flow. LLAGN have also been found to lie on  the  so-called  ``fundamental   plane  of  black  hole  activity'' \\citep{merloni03,falcke04}, which attempts  to unify the emission from all  sources around  black holes,  over a  large range  of  masses and luminosities,  from  Galactic sources  to  powerful  quasars. But  the origin of such  a ``fundamental plane'' and its  relationship with the origin   of   the   radiation    is   still   a   matter   of   debate \\cite[e.g][]{bregman05,koerding06}.  RIAF  models  have  been  applied  to  several  sources  belonging  to different classes,  such as low-luminosity  radio galaxies, ``normal'' ellipticals,       the       Galactic       center       Sagittarius~A \\citep[e.g.][]{quataert99,dimatteo00,dimatteo03}.   However,  for most of these  objects, the models  cannot be properly  constrained, mostly because the nuclear radiation is  swamped by other processes.  This is particularly  problematic in  the optical  band, which  appears  to be crucial to  fix the models, where  the stellar emission  from the host galaxy is substantial.  It  is indeed in the IR-to-UV region that different  accretion disk  models  are expected  to  show the  largest difference  in  spectral  shape.   RIAFs  should lack  both  the  \"big blue-bump\" and  the IR (reprocessed) bump,  which instead characterize optically thick,  geometrically thin  accretion disk emission  and the surrounding  heated  dust.   For  example, in  low  luminosity  radio galaxies  non-thermal  emission from  the  jet  dominates the  optical nuclear  radiation  \\citep{pap1}, while  the  Galactic  center is  not visible in the optical because it is hidden by a large amount of dust.  Therefore, neither of the above mentioned classes of objects appear to be suitable  laboratories to  test models of  low-efficiency accretion through  the analysis  of their  overall SED.   Recently,  among LLAGN which  show very  low Eddington  ratios $L_{bol}/L_{Edd}  << 10^{-3}$, \\citet{papllagn} have found that a  class of LLAGN, mainly composed by LINERs  and  low-luminosity  Seyfert~1  galaxies, show  faint  optical unresolved  nuclei in  HST images  that may  be interpreted  as direct radiation from  a very low  efficiency accretion flow.  In  fact, when the radio-optical properties of  LLAGN are considered, Seyfert, LINERs and low  luminosity radio galaxies separate into  different regions of diagnostic planes,  according to the  properties of their  nuclei.  If this  interpretation  is correct,  we  now  have  a powerful  tool  to identify the  nature of the  nuclear radiation (i.e.   jet-dominated or accretion-dominated).   The best  possibility  of detecting  radiation from an ADAF-like  process would then be to  study unobscured Seyferts of lowest luminosity, as well as  a sub-class of LINERs.  In all other objects other radiation processes dominate.  In  this   paper  we  further  test  the   picture  outlined  in \\citet{papllagn} by studying  the nuclear spectral energy distribution of a galaxy,  \\object{NGC~4565}, that seems to be  a perfect candidate for hosting  a RIAF around  the central supermassive black  hole.  The object is part of the ``Palomar sample'' of LLAGN \\citep{ho97}, and it is included in both the \\citet{merloni03} and \\citet{falcke04} samples that  were  used to  define  the  ``fundamental  plane of  black  hole activity''. It  is worth  mentioning that NGC~4565  does not  show any significant peculiarity  in that plane.  NGC~4565 is  a nearby (d=9.7 Mpc)  LLAGN  classified  as  a  Seyfert~1.9 because  of  the  possible presence  of a  faint,  relatively  broad (FWHM  =  1750 km  s$^{-1}$) H$\\alpha$ line.   However, as \\citet{ho97broad} have  pointed out, the detection of a broad component is highly uncertain.  As we show in the following,  although it  is  a Type  2  Seyfert, this  object is  only moderately  absorbed, and  the  nuclear radiation  is  visible in  the optical spectral region.  NGC~4565  may thus represent the first clear example of low-luminosity accretion  onto a supermassive black hole in the optical band.  In  Section~\\ref{obs}  we describe  the  {\\it Chandra}  and {\\it  HST} observations, the  data analysis procedures and  flux measurements, in Section~\\ref{discussion}  we  present  the  results,  we  derive  the spectral energy  distribution and  we discuss its  interpretation.  In Section~\\ref{conclusions}  we give  a summary of  our findings  and we draw conclusions.  \\begin{figure*} \\epsscale{1} \\plotone{f1.ps} \\caption{HST WFPC2 ``true-color'' RGB image. The R-channel corresponds to the F814W filter, the B-channel is the F450W and the G-channel is the average of the two filters. Direction to the North is indicated by the arrow. \\label{galaxy}} \\end{figure*} ",
        "conclusions": "\\label{conclusions}   We  have  derived  the  spectral  energy distribution  of  a  peculiar low-luminosity   Seyfert~2   galaxy   which,  despite   its   spectral classification,  basically   shows  no  evidence   for  local  nuclear absorption. The SED is peculiar, as it is almost flat in a $\\log \\nu - log (\\nu  F\\nu)$ representation, with  no sign of  both a UV  bump and thermally  reprocessed IR emission.   The very  low luminosity  of the source associated with a relatively high central black hole mass imply an  extremely small value  of the  Eddington ratio  ($L_o/L_{Edd} \\sim 10^{-6}$). This, together with the position occupied by this object on diagnostic planes  for low  luminosity AGN, represents  clear evidence for  a low  radiative  efficiency  accretion process  at  work in  the innermost  regions  of  NGC~4565.   The direct  detection  of  optical emission  from such  radiative inefficient  processes  is particularly important  for  providing  constraints  to  ADAF  models  or  similar. NGC~4565 is therefore a perfect candidate for studying RIAFs, and more observations aimed at  achieving a complete coverage of  the SED, from the radio-mm to the X-ray bands,  should be taken in order to test the models.  As part of a ``search for RIAFs in LLAGN'', it would also be extremely useful  to derive the SED of objects  that are located in the same region of the diagnostic planes as NGC~4565.  The fact  that the  [OIII] emission line  flux is substantial  in this object implies that  an extended narrow line region,  similar to other Seyfert galaxies, is still  present in NGC~4565. A possible intriguing scenario  is  that the  active  nucleus  has recently  ``turned-off'', switching  from a  high  efficiency, standard,  accretion  disk, to  a radiative inefficient accretion process.  However, since the EW of the [OIII]5007 emission  line is  rather small, with  the present  data we cannot rule out  that the amount of ionizing photons  from the RIAF is sufficient to produce the observed [OIII] flux."
    },
    "0601/astro-ph0601303_arXiv.txt": {
        "abstract": "We present four improved empirical relationships useful for estimating the central black hole mass in nearby AGNs and distant luminous quasars alike using either optical or UV single-epoch spectroscopy. These mass-scaling relationships between line widths and luminosity are based on recently improved empirical relationships between the broad-line region size and luminosities in various energy bands and are calibrated to the improved mass measurements of nearby AGNs based on emission-line reverberation mapping. The mass-scaling relationship based on the \\hb\\ line luminosity allows mass estimates for low-redshift sources with strong contamination of the optical continuum luminosity by stellar or non-thermal emission, while that based on the \\civ\\,$\\lambda$1549 line dispersion allows mass estimates in cases where only the line dispersion (as opposed to the FWHM) can be reliably determined. We estimate that the absolute uncertainties in masses given by these mass-scaling relationships are typically around a factor of 4. We include in an Appendix mass estimates for all the Bright Quasar Survey (PG) quasars for which direct reverberation-based mass measurements are not available. ",
        "introduction": "A problem of current interest is determination of the mass function of the central black holes in active galactic nuclei (AGNs) and quasars over the history of the Universe in order to determine how these black holes evolve with time. Unfortunately, measurement of black hole masses by direct methods such as modeling of stellar or gas dynamics requires high spatial resolution and is thus limited to relatively nearby galaxies. Moreover, the brightness of the AGN itself makes it extremely difficult to observe suitable stellar absorption lines for dynamical studies within the black hole radius of influence, and the complex gas dynamics in AGNs frustrate attempts to disentangle the emission-line kinematics. Megamaser dynamics have been used successfully to measure the black hole mass in NGC 4258 (Miyoshi et al.\\ 1995), but this required particular fortunate circumstances that do not seem to be generally realized. Thus, for AGNs and quasars, the most promising method for measuring black hole masses is reverberation mapping of the broad emission lines. The advantages of this technique are that (a) it does not depend on angular resolution and (b) it yields straightforward empirical relationships that provide effective secondary indicators that can be used to estimate the masses of large numbers of AGNs and quasars based on single observations. The disadvantages of the technique are that (a) the accuracy of reverberation-based black hole masses are fundamentally limited by our lack of knowledge of the detailed structure and kinematics of the BLR and (b) it is observationally demanding.  Reverberation-based black hole masses are computed from the virial equation \\begin{equation} \\label{eq:virial} M_{\\rm BH} = \\frac{f\\,R \\Delta V^2}{G}, \\end{equation} where $R$ is the size of the region as estimated by the mean emission-line lag $\\tau$ (time delay relative to continuum variations), i.e.,  $R = c\\tau$, $\\Delta V$ is the emission-line width (preferably the width of the {\\em variable} part of the emission line), and $f$ is a scale factor of order unity that depends on the structure and geometry of the BLR. Two lines of evidence suggest that these masses have some validity: \\begin{enumerate} \\item In AGNs for which time delays have been measured for multiple lines, there appears to be a virial relationship between time delay and line width, i.e., $\\tau \\propto \\Delta V^{-2}$ (Peterson \\& Wandel 1999, 2000; Onken \\& Peterson 2002; Kollatschny 2003). \\item In reverberation-mapped AGNs for which host galaxy bulge velocity dispersions $\\sigma_*$ are available, the reverberation-based masses $M_{\\rm BH}$ are consistent with the $\\Msigma$ relationship seen in quiescent galaxies (Gebhardt et al.\\ 2000; Ferrarese et al.\\ 2001; Onken et al.\\ 2004; Nelson et al.\\ 2004). \\end{enumerate}  Reverberation mapping also shows that there is a simple relationship between the size of the BLR and the continuum luminosity $L$ of the AGN of the form $R \\propto L^\\gamma$ (Kaspi et al.\\ 2000; 2005). This is an important result, not only because it constrains the physics of the BLR, but also because it provides a secondary method of estimating the black hole masses by using $L^\\gamma$ as a surrogate for $R$ in eq.\\ (\\ref{eq:virial}). Since a single spectrum of an object in principle yields both $L$ and a line width $\\Delta V$, we have a powerful tool for estimating the masses of large populations of quasars. Wandel, Peterson, \\& Malkan (1999) carried out some preliminary tests of this method using the \\hb\\ emission line. Vestergaard (2002; hereafter Paper I) used the \\civ\\,$\\lambda1549$ emission line to probe much higher redshifts, up to $z \\approx 6$. There have been several other extensions of this methodology. McLure \\& Jarvis (2002) used the Mg\\,{\\sc ii}\\,$\\lambda2798$ emission line in a similar study. Wu et al.\\ (2004) suggested that recombination-line luminosities should be used since they are a better measure of the ionizing continuum that drives the line variations than the longer-wavelength continuum, which may be contaminated by hard-to-quantify jet emission or host-galaxy starlight, depending on the wavelength at which the continuum is measured.  Since the original papers appeared, there have been a number of significant developments that have led us to decide to revisit the mass-scaling relationships based on the \\hb\\ and \\civ\\ emission lines. Specifically, \\begin{enumerate} \\item The reverberation-mapping database that provides the fundamental calibration for the mass-scaling relationships has been completely reanalyzed (Peterson et al.\\ 2004). Of particular relevance here is that some inadequate or poor data were identified and removed from the database. \\item The reverberation-based masses have now been empirically scaled to the quiescent galaxy black hole mass scale through use of the \\Msigma\\ relationship (Onken et al.\\ 2004). The zeropoint of the AGN \\Msigma\\ relationship was adjusted to that of quiescent galaxies by determining a mean value for the scale factor $f$. \\item The radius--luminosity ($R-L$) relationship between the broad-line region size and continuum luminosity has been updated based on the reanalyzed reverberation data (Kaspi et al.\\ 2005), and new {\\it HST} imaging of reverberation-mapped AGNs enable us to correct for host galaxy contamination of the optical continuum luminosity measured from spectra (Bentz \\et 2006). \\item Additional spectra of reverberation-mapped AGNs have become available in the public domain, making it possible to improve the calibration and better quantify the uncertainties of mass estimates based on single-epoch spectra. \\end{enumerate} In contrast to Paper I, we have also adopted the current benchmark cosmology with $H_0$ = 70 ${\\rm km~ s^{-1} Mpc^{-1}}$, $\\Omega_{\\Lambda} = 0.7$, and $\\Omega_m = 0.3$; in Paper I, we used  $H_0$ = 75 ${\\rm km~ s^{-1} Mpc^{-1}}$, $q_0 = 0.5$, and $\\Lambda = 0$ to effect more direct comparisons of quasar luminosities with those in previous work.  In the following,  we describe the data and the spectral measurements (\\S\\ref{data}), perform the calibration of the single-epoch unscaled mass estimates (\\S\\ref{Mcal}) and determine the inherent statistical uncertainties of these relationships (\\S\\ref{errors}). In \\S\\ref{discussion}, we briefly discuss (1) the improvements in the updated mass-scaling relationships and (2) the appropriateness of using the \\civ\\ emission line for mass estimates. Our main results are summarized in \\S\\ref{summary}.  Also, there has been a controversy over the use of the \\civ\\ emission line in particular (\\eg Baskin \\& Laor 2005) and we wish to address that issue as well. In Appendix A, we scrutinize the data used by Baskin \\& Laor that led them to challenge the validity of \\civ-based mass estimates and we conclude that a more suitable selection of data largely removes the problems that they identified. Finally, in Appendix B, we provide a complete list of estimated masses for all those quasars from the Bright Quasar Survey (the Palomar--Green or ``PG'' quasars; Schmidt \\& Green 1983) for which reverberation-based masses are not available. ",
        "conclusions": "}  We conclude this section with a summary of the mass-scaling relationships for obtaining black hole masses estimates from single-epoch spectra. \\begin{enumerate} \\item {\\bf FWHM(\\mbox{\\boldmath \\Hbeta}) and \\mbox{\\boldmath $L_{\\lambda}$(5100\\,\\AA)}.} For the optical continuum luminosity and FWHM of the \\Hbeta\\ broad component, \\begin{equation} \\log \\,M_{\\rm BH} (\\rm H\\beta) = \\log \\,\\left[ \\left(\\frac{\\rm FWHM(H\\beta)}{1000~km~s^{-1}} \\right)^2 ~ \\left( \\frac{\\lambda \\it L_{\\lambda} {\\rm (5100\\,\\AA)}}{10^{44} \\rm erg~s^{ -1}}\\right)^{0.50} \\right]\t+ (6.91 \\pm 0.02). \\label{logMopt_L51.eq} \\end{equation} The sample standard deviation of the weighted average zeropoint offset, which shows the intrinsic scatter in the sample, is $\\pm 0.43$ dex. This value is more representative of the uncertainty in the zero-point than is the formal error.  \\item {\\bf FWHM(\\mbox{\\boldmath \\Hbeta}) and \\mbox{\\boldmath $L(\\Hbeta)$}.} For the \\Hbeta\\ broad-component luminosity and FWHM, \\begin{equation} \\log \\,M_{\\rm BH} (\\rm H\\beta) = \\log \\,\\left[ \\left( \\frac{\\rm FWHM(H\\beta)}{1000~km~s^{-1}} \\right)^2 ~ \\left( \\frac{\\it L(\\rm H\\beta)}{10^{42} \\rm erg~s^{- 1}}\\right)^{0.63} \\right]\t+ (6.67 \\pm 0.03). \\label{logMopt_LHb.eq} \\end{equation} The sample standard deviation of the weighted average zeropoint offset is $\\pm 0.43$ dex.  \\item {\\bf FWHM(\\mbox{\\boldmath \\civ}) and \\mbox{\\boldmath $L_{\\lambda}$(1350\\,\\AA)}.} For the ultraviolet continuum luminosity and the FWHM of the \\civ\\ line, \\begin{equation} \\log \\,M_{\\rm BH} \\mbox{(\\rm \\civ)} = \\log \\,\\left[ \\left(\\rm \\frac{FWHM\\mbox{(\\rm \\civ)}}{1000~km~s^{-1}} \\right)^2 ~ \\left( \\frac{\\lambda L_{\\lambda} (1350\\,{\\rm \\AA})}{10^{44} \\rm erg~s^{- 1}}\\right)^{0.53} \\right]\t+ (6.66 \\pm 0.01). \\label{logMuv_fw.eq} \\end{equation} The sample standard deviation of the weighted average zeropoint offset is $\\pm0.36$ dex.  \\item {\\bf \\mbox{\\boldmath $\\sigma_l$(\\mbox{\\boldmath \\civ})} and \\mbox{\\boldmath $L_{\\lambda}$(1350\\,\\AA)}.} For the ultraviolet continuum luminosity and the dispersion of the \\civ\\ emission line, \\begin{equation} \\log \\,M_{\\rm BH} \\mbox{(\\rm \\civ)} = \\log \\,\\left[ \\left(\\frac{\\sigma_l\\mbox{(\\rm \\civ)}} {\\rm 1000\\,~km~s^{-1}} \\right)^2 ~ \\left( \\frac{\\lambda L_{\\lambda} (1350\\,{\\rm\\AA})}{10^{44} \\rm erg~s^{- 1}}\\right)^{0.53} \\right]\t+ (6.73 \\pm 0.01). \\label{logMuv_sigma.eq} \\end{equation} The sample standard deviation of the weighted average zeropoint offset is $\\pm0.33$ dex. \\end{enumerate}  As noted earlier, the $L_{\\lambda}$(1450\\,\\AA) luminosity can be straightforwardly be substituted for $L_{\\lambda}$(1350\\,\\AA) without error or penalty in precision.  }  In Paper~I, we found that the predominant source of scatter in the mass-scaling relationships is traceable to the relatively large scatter in the $R - L$ relationship (Kaspi \\et 2000). The recent improvements in both the $R - L$ relationships (Kaspi \\et 2005) and in the reverberation-based mass measurements (Peterson \\et 2004) have correspondingly decreased the scatter in the mass estimates, at least for the estimates based on UV data. The 1$\\sigma$ scatter is now only a factor of about 2, compared to the factor of 3.2 found in Paper~I. Implicit in this estimate is that the reverberation-mapped AGNs are reasonably representative of the AGN and quasar population as a whole.  It appears that no improvement has been achieved for the optical relationship based on FWHM(\\Hbeta) and $L_{\\lambda}$(5100\\,\\AA). However, this relationship was, in fact, {\\it not calibrated} in Paper~I; the single-epoch estimates based on the relationships quoted by Kaspi \\et (2000) were merely confirmed to be consistent with the reverberation-based masses. In the current calibration, these optical mass estimates now distribute more evenly around zero offset (Fig.~\\ref{MoptCal-problty.fig}, left panel) compared to those of Paper~I (cf.\\ Fig.\\ 6a in Paper~I).  There are a number of advantages to the use of the \\civ\\ line width as a mass indicator (eq.\\ \\ref{logMuv_fw.eq}), as argued in Paper~I and by Warner, Hamann, \\& Dietrich (2003) and Vestergaard (2004). It has been argued, however, that \\civ\\ may be an inappropriate choice for this application on the grounds that the dynamics of the \\civ-emitting gas may not be determined primarily by gravitation; in particular, it has been suggested (1) that the line profiles might be affected by obscuration within the line-emitting region, and (2) that there are reasons to believe that a significant fraction of the \\civ\\ line arises in an outflowing wind. We address these two issues in turn.  Richards \\et (2002) used a large sample of SDSS quasars to investigate the observed blueshift of the peak of the \\civ\\ profile relative to the low-ionization lines, which is a well-known phenomenon (\\eg Gaskell 1982; Wilkes 1984; Tytler \\& Fan 1992). They suggest that the apparent blueshift may actually be a lack of emission in the red profile wing, suggestive of occultation or obscuration, i.e., a non-gravitational alteration of the profile. Richards \\et divide their database into four groups based on the magnitude of the \\civ\\ blueshifts (their Fig.~5 and Table~2) and form composite spectra for each of these subsets. They find that the most blueshifted profile typically has FWHM(\\civ) = $1600 \\pm 300\\,\\kms$, which is only 15\\% broader than the \\civ\\ profile from the least blueshifted composite. There is thus a bias in the sense that the subset with the most strongly affected profiles will have masses overestimated by $\\sim$30\\%. However, this effect is considerably smaller than the typical uncertainties even in the reverberation-based masses, which as noted earlier are uncertain by a factor of  $\\sim2.9$ (Onken \\et 2004).  It is sometimes stated (\\eg Dunlop 2004; Bachev \\et 2004; Shemmer \\et 2004; Baskin \\& Laor 2005) that the \\civ\\ profile is unsuitable for estimating AGN black hole masses from single-epoch spectra because the component of \\civ\\ that arises in an outflowing wind, as is the case with other higher-ionization lines, is much more significant than in, say, lower ionization lines such as \\mgii\\ or \\Hbeta. The \\civ\\ profile of some narrow-line Seyfert~1 galaxies (NLS1s) may contain a significant contribution from a wind, and indeed Vestergaard (2004) cautioned against using \\civ\\ to estimate the masses of black holes in this type of object. It should be pointed out, however, that not all NLS1s exhibit this behavior. But there does seem to be a tendency for higher-luminosity NLS1s to exhibit more highly asymmetric profiles (Leighly 2000), and there is thus a concern that the masses of luminous quasars may be overestimated by failing to account for a strong wind contribution. Vestergaard (2004) examined in detail the \\civ\\ line asymmetries of high-luminosity quasars in the range $1.5 \\leq z \\leq 5$ and found that none of the profiles resembled in any way the triangular shape seen in the spectra of some luminous NLS1 galaxies, like I\\,Zw\\,1 (e.g., Vestergaard \\& Wilkes 2001). Also, such profiles are not seen by Richards et al.\\ (2002). With a conservative selection of objects with possible blue wing asymmetries, Vestergaard (2004) determined the typical mass and luminosity of the quasars with and without these asymmetries and found insignificant differences between the two subsets.  Baskin \\& Laor (2005) find that the \\civ\\ equivalent width EW(\\civ) correlates with the Eddington luminosity ratio computed from the \\Hbeta-based masses, \\lol(\\Hbeta). Since they do not find a similarly strong correlation between EW(\\civ) and \\lol(\\civ) for their sample of BQS sources, they argue that the \\civ\\ profile may yield biased black hole mass estimates.  As described in Appendix~\\ref{Laor.app}, we have undertaken a reanalysis of the Baskin \\& Laor sample in which we removed (a) the NLS1s, (b) low-quality {\\em IUE} spectra of PG quasars, (c) objects with strong absorption near the \\civ\\ emission-line peak, and (d) measurements based on what we consider to be a dubious assumption about a narrow-line component of the \\civ\\ emission line, and we find that the problems cited by Baskin \\& Laor are much less severe. The remaining \\civ\\ FWHM values scatter within $\\pm0.2$\\,dex of FWHM(\\Hbeta), thereby showing reasonable consistency. This result is also consistent with that of Warner \\et (2003), who find, on average, reasonable agreement between the mass estimates based on \\Hbeta\\ and those based on \\civ.  Indeed, the most compelling reason to have some confidence in the \\civ-based mass estimates is because in the handful of objects in which reverberation results are available for multiple emission lines, the virial products are consistent for {\\em all} the measured emission lines, including \\civ\\ (Peterson \\& Wandel 1999, 2000; Onken \\& Peterson 2002)."
    },
    "0601/astro-ph0601135_arXiv.txt": {
        "abstract": "{\\it Spitzer} mid-infrared images of the dusty warped disk in the galaxy Centaurus\\,A show a parallelogram-shaped structure.  We successfully model the observed mid-infrared morphology by integrating the light from an emitting, thin, and warped disk,  similar to that inferred from previous kinematic studies.   The models with the best match to the morphology lack dust emission within the inner 0.1 to 0.8\\,kpc, suggesting that energetic processes near the nucleus have disturbed the inner molecular disk, creating a gap in the molecular gas distribution. ",
        "introduction": "Centaurus\\,A (NGC\\,5128) is the nearest of all the giant radio galaxies. Because of the disk of gas and dust in its central regions,  Centaurus\\,A is suspected to be the product of a merger of a small gas rich  spiral galaxy with a larger elliptical galaxy \\citep{baade}.   Numerical simulations of such mergers  predict large shell-like features \\citep{hernquist88} that have been observed in Centaurus\\,A over a large range of radii \\citep{malin83,peng02}.  Some contain atomic and molecular gas, implying that they were formed from material stripped from a galaxy relatively recently, fewer than a few galactic rotation periods ago \\citep{schiminovich94,charmandaris}. Tidal features associated with this debris have also been identified, supporting the relatively short  (few times $10^8$ years) estimated timescale since the merger took place  \\citep{peng02}.  In its central regions, NGC\\,5128 exhibits a well recognized, optically-dark band of absorption across its nucleus. This dusty disk was first modeled as a transient  warped structure by \\citet{tubbs80}. \\citet{bland86}, \\citet{bland87}, \\citet{quillenCO}, and \\citet{nicholson92} found that the kinematics of the ionized and molecular gas are well modeled by a warped disk composed of a series of inclined  connected rings undergoing circular motion  (as also explored for other galaxies with peculiar morphology by \\citealt{steiman92}). The model explored by Quillen, Graham \\& Frogel (1993) modified the kinematic model by \\citet{quillenCO} to fit the morphology of the absorptive, dusty disk seen in near-infrared images  and proposed a timescale of about $200$ million  years since the core of an infalling spiral galaxy reached and merged with the elliptical galaxy nucleus. An initially flat disk, misaligned with the galaxy principal symmetry axis, becomes increasingly corrugated as a function of time. The short timescale estimated since the merger in NGC\\,5128 is approximately consistent with the timescale suggested by the presence of tidal debris and by the shell-like features containing atomic hydrogen \\citep{schiminovich94,peng02} and molecular gas \\citep{charmandaris}. An alternative model accounting for the warped disk is the polar ring model by \\citet{sparke96},  consistent with the polar orbit of the disk implied by the galaxy isophotes in the outer galaxy \\citep{malin83,haynes83,peng02} but requiring a longer timescale to  account for the twist of the warp.  More recent observations of the central region include submillimeter imaging with the Submillimeter Common User Bolometer Array (SCUBA) and mid-IR imaging with the Infrared Camera (ISOCAM) on the {\\it Infrared Space Observatory (ISO)} satellite \\citep{mirabel99,leeuw02}.   At these wavelengths, the dusty disk is seen in  emission rather than absorption. At 100--200\\,pc from the nucleus the galaxy contains a molecular circumnuclear  disk that has been studied in molecular line emission \\citep{israel90,israel91} and resolved in Pa\\,$\\alpha$ emission \\citep{schreier96,marconi01}.  For a recent summary of the wealth of observational studies carried out on  this peculiar and active nearby galaxy, see the comprehensive review by \\citet{israel}. Based on the discussion by \\citet{israel},  we adopt a distance to Centaurus\\,A of $3.4$\\,Mpc.  At this distance $1\\arcm$ on the sky corresponds to $\\sim 1$\\,kpc.  In this manuscript we focus on the geometry of the dusty disk in Centaurus\\,A as seen from {\\it Spitzer\\/} Infrared Array Camera (IRAC) images.  These images resolve the structure of the disk more clearly than previous observations and also show the disk out to larger radii.  They provide us with an opportunity to study the geometry of the warped disk in much higher detail than previously possible. Observations are described in \\S\\,2. Our geometric model is described in \\S\\,3.  A discussion and summary follow in \\S\\,4. ",
        "conclusions": "By integrating the light through an emitting, optically thin, dusty, and warped disk, we have successfully matched the morphology of Centaurus\\,A seen in mid-infrared {IRAC} images.  We confirm previous proposals that the disk morphology is well explained by a warped disk \\citep{bland86,bland87,quillen,nicholson92,leeuw02} rather than a barred one \\citep{mirabel99}.  The disk is nearly edge-on with respect to the viewer, but tilts so that folds appear above and below the galaxy equator.  There is a fold south of the nucleus at a radius $r\\sim 60\\arcs$ responsible for high extinction seen in near-infrared images, and a fold north of the nucleus at larger radii responsible for the northern edge of the dust lane seen in optical and images.  Extinction features seen in the near-infrared extinction map correspond to folds in the disk that are located nearer the observer and so absorb more background stellar light from the galaxy.  In the mid infrared, however, folds on both the near and opposite side of the galaxy correspond to bright emission features. The inner folds account for the parallelogram shape, while the outer folds account for the northern edge of the dust lane seen in optical images and a fainter oval of emission seen outside the parallelogram in the IRAC images.  The disk geometry we use to match the mid-infrared morphology is similar to that found previously by \\citet{quillen},  and is also similar to that required to fit the CO and H\\,$\\alpha$ velocity field \\citep{bland87,quillenCO,nicholson92}. The geometric warp models by \\citet{quillen,nicholson92} have been predictive; they provide good matches to the mid-infrared morphology.  Some differences in the precession angle exist in the model at $r \\gtrsim 100\\arcs$ compared to previous work. Previous CO, H\\,$\\alpha$ spectra and near-infrared imaging were not sensitive enough to provide tight kinematic constraints on the outer disk. The {\\it Spitzer\\/} data have allowed us to extend the model past $r\\sim 100\\arcs$ compared to previous models and see closer in to the nucleus where the extinction is high at shorter wavelengths.  Better constraints on the disk geometry could be achieved in the future by fitting observations at more than one wavelength.  For example, modern high resolution kinematic observations that are fit along with the {\\it Spitzer\\/} data using the same model would allow much stronger constraints on the disk geometry and gas and dust distribution than the slight modification to previous geometric models that we have used here.  The warp disk model suggests that there is a gap in the dusty disk at $6\\arcs \\lesssim r\\lesssim 50\\arcs$. A gap exists in same region in the gas distribution.  \\citet{marconi01} saw a deficit of Pa\\,$\\alpha$ emission near the nucleus at intermediate velocities.   \\citet{nicholson92} suggested that the discrepancy between the measured linearly rising rotation curve within a radius of $60\\arcs$ might be explained by a hole in the ionized gas distribution. It is not easy to determine if there is a gap in the dust distribution since the infrared surface brightness depends on the inclination of the disk and the disk is highly corrugated.   If the disk twists to lower inclinations (closer to face-on) at small radii, it would have a lower surface brightness near the nucleus.  We have explored models with smoothly varying radial surface brightness distributions, and smoothly varying precession and inclination angles.  We find that only models with a deficit of dust interior to $\\sim 50\\arcs$ resemble the mid-infrared images.  We exclude models that have more than $\\sim 1/5$ the surface density within $r=50\\arcs$ as at that radius.  We conclude that there is a gap in the gas and dust distribution between 0.1 and 0.8\\,kpc from the nucleus.  The inner radius of the gap we have taken from studies of the circumnuclear disk \\citep{marconi01,israel91}, and the outer radius is estimated from our model.  It is interesting that only the region between the circumnuclear disk (100--200\\,pc) and $\\sim 0.8$\\,kpc has been depleted of gas and dust.  While we find no evidence for a large scale stellar bar in Centaurus\\,A, a dense gas disk could have exhibited dynamical instabilities depositing gas into the circumnuclear disk (e.g., \\citealt{shlosman}).  Energetic star formation or activity associated with the black hole can deplete and evacuate the central region of a galaxy (e.g., \\citealt{springel05,veilleux05}), though it is not clear how a circumnuclear disk would be protected from or reformed following this activity.  Nearby galaxies exhibit evacuated central regions or circumnuclear rings.  For example, the Circinus galaxy has a $\\sim 500$\\,pc radius molecular ring \\citep{curran98} and contains a Seyfert nucleus.  M82 contains a 1\\,kpc radius circumnuclear molecular ring;  however, a previous epoch of star formation has occurred within this ring (e.g., \\citealt{forster03}). Future observational studies may differentiate between the role of star formation, dynamical instabilities and nuclear activity in disrupting the gas and dusty disk in Centaurus\\,A.  The best matching geometric warp model requires a change in the slope of the precession angle at a radius of about $r\\sim 100\\arcs$.  Two previous models account for this twist.  \\citet{quillen} used a model in which the galaxy was prolate in its outer region and the galaxy ellipticity abruptly dropped to zero within $r \\sim 80\\arcs$.  However, this abrupt drop is not consistent with the isophote shapes at $3.6\\,\\mu$m.  A prolate model (with the advantage of a relatively short timescale) might account for the disk geometry if modified to include the self-gravity of the disk.  The polar ring model by \\citet{sparke96} naturally accounts for the reversal, but requires a longer timescale to operate.  These dynamical models could be updated to use a more accurate mass distribution and rotation curve. Improvements to these models may lead to better understanding of the galaxy merger that created Centaurus\\,A's peculiar morphology, as well as the merger's role in feeding the active galactic nucleus."
    },
    "0601/astro-ph0601245_arXiv.txt": {
        "abstract": "We present \\chandra and \\xmm observations of 12 bright  ($\\rm f(2-10 keV) > 10^{-13}$ \\funits) sources from the \\asca SHEEP (Search for the High Energy Extragalactic  Population) survey.  Most of these have been either not observed or not detected previously  with the  {\\it ROSAT} mission and therefore they constitute a sample biased towards   hard sources.  The \\chandra observations are important in  locating with  accuracy the optical counterpart of the X-ray sources. Optical spectroscopic observations show that  our sample is associated with both narrow-line (NL) (six objects),   and Broad-Line (BL) AGN (five objects) with one source  remaining unidentified.  Our sources cover the redshift range  0.04 to 1.29 spanning luminosities from $10^{42}$ to $10^{45}$ \\lunits (2-10 keV). The NL sources have preferentially  lower redshift (and luminosity) compared with the BL ones. This can be most easily explained in a model where the  NL AGN are intrinsically less luminous than the BL ones in line with the results of Steffen et al. The X-ray spectral fittings show a roughly equal number  of obscured ($N_H>10^{22}$ \\cunits) and unobscured  ($N_H<10^{22}$ \\cunits) sources. There is a clear tendency for obscured sources to  be associated with NL  AGN and unobscured sources with BL ones. However, there is a  marked exception with the highest obscuring column observed at a BL AGN  at a redshift of z=0.5. ",
        "introduction": "\\chandra surveys resolved 80 per cent of the X-ray background  at hard energies (i.e. 2-10 keV) Brandt et al. (2002),  Giacconi et al. (2002), Alexander et al. (2003), shedding ample light on its origin.  Optical identifications of the X-ray sources are showing  that these are mainly AGN. The largest fraction of these are absorbed in X-rays, presenting an average spectrum with a slope of $\\Gamma=1.4$ (Tozzi et al. 2001) similar to the slope of the X-ray background. Large area bright surveys performed  with XMM-Newton in the 2-10 keV band are complementary   as they help us to explore the bright, nearby counterparts  of the sources detected in the deep Chandra fields. One remarkable finding from these bright surveys  is the scarcity of obscured AGN ($\\rm N_H>10^{22}$ \\cunits)  at bright fluxes (Piconcelli et al. 2002, Georgantopoulos et al. 2004, Perola et al. 2004). This comes in contrast to the predictions of the population  synthesis models (e.g. Comastri et al. 1995, Gilli et al. 2001).  Selection in the hardest band allowed by imaging surveys  (5-10 keV) facilitates the selection of the most obscured  candidates, as this bandpass is relatively unbiased to absorption.  Indeed columns as high as a few times $10^{23}$ \\cunits ~ do not   practically affect X-ray energies above 5 keV. The HELLAS survey (Fiore et al. 1999)  pioneered these studies using the MECS detector on board {\\it BeppoSAX}.  The HELLAS survey covered  an area of 85 $\\rm deg^2$ detecting about 150 sources  down to a flux limit  of $5\\times 10^{-14}$ \\funits in the 5-10 keV band.  The SHEEP survey further explored this energy bandpass with {\\it ASCA} GIS. It comprises of 69 objects serendipitously  detected in a non-contiguous area of 39 $\\rm deg^2$ down to a flux limit of  $\\sim10^{-13}$\\funits in the 5-10 keV band (Nandra et al. 2003).  35 of these objects were detected by {\\it ROSAT}, in either pointed observations or in the All Sky Survey, with 13 of them  having secure optical counterparts in optical  catalogues. For the remaining 34 sources,  which were either not observed or not detected by {\\it ROSAT}, it is very difficult  to locate the optical counterpart as the {\\it ASCA}  error box is large (typically 60 arcsec rms). We have thus obtained \\chandra snapshot observations  in order to  extract reliable positions ($<$1 arcsec) of the X-ray source that will enable us to search for their optical counterparts.  Preliminary results from  these  observations are presented in Nandra et al.~(2004).  Here, we report our results from X-ray and optical observations of 12 out of the 34 \\chandra sources which have a $2-10$ keV flux larger than $10^{-13}$ \\funits (estimated assuming a Photon index $\\Gamma=1.9$). We have obtained optical spectroscopy for 11 of these objects. Furthermore, eight of these have been also observed by \\xmm, providing  good quality spectra. Consequently, for these bright sources, we can derive their X-ray spectral parameters (i.e. \\nh\\ and $\\Gamma$) using proper spectral fits with good photon statistics (instead of hardness ratios only). Moreover,  Nandra et al.~(2004) noticed that a fraction of the \\chandra sources are much fainter  (in some cases more than an order  of magnitude) than the original {\\it ASCA} source. Although variability could play some role,  there is clearly some ambiguity about the reality  of these sources. Watanabe et al. (2002) who performed  \\chandra pointings of hard {\\it ASCA} sources also report the same problem. By choosing to study initially only the bright sources, we minimize the possibility that a source is spurious. The optical counterparts for three of the sources in the present SHEEP subsample have already been presented in Nandra et al. (2004).We include them in the present work as they satisfy the flux limit we have set, and we also present the results from \\xmm observations of them. ",
        "conclusions": "Our survey provides a glimpse of bright, nearby examples  of the hard AGN population which produces an  appreciable fraction of the X-ray background. The subsample of the SHEEP survey presented here  contains both Type I (BL) and Type II (NL) AGN, according to the optical classification.  Although the number of objects that we study is small, we find that there is a considerable difference in their redshifts, with NL objects preferentially found at redshifts lower than those of BL AGN.  The cross-correlation of {\\it GINGA} and {\\it HEAO-1} X-ray background maps above 2 keV with nearby galaxy catalogues, demonstrates that an appreciable fraction  of the hard X-ray emission may arise at low redshift, $z\\leq0.1$ (Jahoda et al. 1991, Lahav et al. 1993).  Our NL AGN clearly provide examples of this  nearby AGN population.    The distance separation between NL and BL AGN  is ubiquitous in hard X-ray surveys with the redshift peak  of the populations depending on the surveys's flux limit.  Our results are consistent with the hypothesis that NL AGN are, on average, less luminous than $10^{44}$ \\lunits.  Although our statistics is clearly small, the same  effect i.e. the NL AGN  occupying preferentially the low luminosities, can be also witnessed in the \\xmm hard 4.5-7.5 keV sample of Caccianiga et al. (2004)  which contains 28 sources. Five out of the seven NL AGN in Caccianiga et al. (2004) have 2-10 keV intrinsic luminosities lower than $10^{44}$ \\lunits.  Obviously, larger surveys are needed in order to confirm unambiguously that  NL AGN are intrinsically less luminous than the BL AGN. For example, this could be the case in a scenario where highly luminous AGN (QSO) come preferentially in a ``Broad Line flavour\" in contrast to the low luminosity  AGN (Seyferts) which  appear as both NL or BL according to the viewing angle in respect to the obscuring torus. Steffen et al. (2003) find that the fraction of NL AGN decreases at high redshift, $z>1$. Treister et al. (2004) observe the same effect but raise the criticism that this could be caused by incompleteness in the spectroscopic identification. Indeed luminous NL AGN at high redshift cannot be easily detected in the optical owing to large amounts of redenning (e.g. Fiore et al. 2003). The results emerging from our own survey as well as from that of Caccianiga et al. (2004), which has practically complete spectroscopic identification,  lend support to the findings of Steffen et al. (2003). We note that the density of NL AGN in Steffen et al. (2003) decreases only at $z>1$. A direct comparison of the redshift distribution of our sample with that of Steffen et al. is not straightforward as the flux limit of the latter sample is much deeper. Nevertheless, if the NL AGN are intrinsically fainter than the BL AGN, the effect is that the NL AGN are observed at lower redshifts compared with the BL ones in any given survey. Similarly, recent surveys (e.g. Barger et al. 2004) show evidence of ``cosmic downsizing'' where the low luminosity AGN have been formed at a later epoch compared with the bright QSOs. In particular, the peak of the redshift distribution of the low luminosity AGN peaks around z=0.7. Thus given that the NL AGN are low luminositiy sources this provides an additional reason why these are observed at relatively low redshifts.   The column density distribution is almost equally divided  between high and low column densities (table \\ref{chandrafits}). In particular the fraction of sources with $N_H>10^{22}$ \\cunits is $58\\pm22$ per cent. It has to be mentioned that we are dealing with a biased  'hard' sample  since we have primarily observed with \\chandra, the  sources for which no {\\it ROSAT} positions are available (ie those either not detected or not  observed by {\\it ROSAT}) and hence the above fraction should be only viewed as an upper limit. Therefore, our sample cannot provide a fair comparison with the X-ray background population synthesis models of e.g. Comastri et al. (2001). We note that Caccianiga et al. (2004) find a fraction of $26\\pm10 $ for the fraction of obscured AGN, inconsistent with the predictions of Comastri et al. (2001). We should be able to provide stronger constraints, when all the objects in the SHEEP survey have been observed by \\chandra.   We find that NL objects are more obscured compared to BL objects, as expected from Unification models.  However, there is a  notable exception as the hardest source  in the sample is associated with a BL AGN (0140.1+0628). The \\xmm spectrum presented here confirms  the column density derived from the \\chandra hardness ratio analysis (Nandra et al. 2004). The observed colour (B--R=0.8) suggests little absorption  in the optical: $A_V\\approx1.5$ assuming B--R=0.3 at z=0.5 according to the  SDSS QSO template spectrum of Vanden Berk et al. (2002).   On the other hand the column density measured  in the X-ray spectrum corresponds  to $A_V\\sim 55$ assuming the Galactic  dust-to-gas ratio (Bohlin, Savage \\& Drake 1978).  This large discrepancy can be naturally explained by sublimation  of the dust (eg Granato, Danese \\& Franceshini 1987).  This may be the case in Broad-Absorption-Line (BAL) QSOs  which usually present large absorbing column densities  in the X-ray spectra, with typically  $N_H>10^{23}$ \\cunits ~ (Gallagher et al. 2002), while they present moderate reddening in their colours (Brotherton  et al. 2001). Actually, from the existing  optical spectrum alone we cannot exclude the possibility that our object is a BAL. The most efficient diagnostic  for the classification of a source as a BAL is  the presence of a broad CIV line (1549 \\AA)  which is outside the spectral window at the redshift  of  our object.  Finally,  one the softest X-ray sources is  associated with a NL source (AXJ1511.7+0758). It is puzzling why the  optical reddening which should be responsible for the  obliteration of the broad-line region does not produce significant absorption in X-rays. Typically, NL   AGN (Seyfert-1.8-1.9.-2.0) present column densities  higher than $10^{23}$ \\cunits (eg Awaki et al. 1991). One possibility is that at the low luminosity regime   - our source has $\\rm L_x\\sim 10^{42}$ \\lunits - no broad-line  region is formed (Nicastro 2000). Indeed, there are a few   bona-fide examples of Seyfert-2 galaxies, (see Panessa \\& Bassani 2002,  Georgantopoulos \\& Zezas 2003) with  no obscuration in X-rays.  Alternatively, the possibility that, at least, the X-ray emission comes  primarily from star-forming processes cannot be ruled out. As discussed earlier the X-ray spectrum requires the presence of a soft thermal component.  This would naturally explain the relatively low level  of X-ray luminosity (see e.g. Zezas, Georgantopoulos \\& Ward 1998)  as well as the soft X-ray spectrum. Nevertheless, the weak $H\\beta$ relative to the [OIII] line as well as the relative strength of the NII relative to the Ha suggest against the star-formation activity."
    },
    "0601/astro-ph0601559_arXiv.txt": {
        "abstract": "The motion of the solar system barycenter with respect to the cosmic microwave background (CMB) induces a very large apparent dipole component into the CMB brightness map at the 3 mK level.  In this Letter we discuss another kinematic effect of our motion through the CMB: the small shift in apparent angular positions due to the aberration of light.  The aberration angles are only of order $\\beta \\simeq 0.001$, but this leads to a potentially measurable compression (expansion) of the spatial scale in the hemisphere toward (away from) our motion through the CMB.  In turn, this will shift the peaks in the acoustic power spectrum of the CMB by a factor of order $1 \\pm \\beta$.  For current CMB missions, and even those in the foreseeable future, this effect is small, but should be taken into account.  In principle, if the acoustic peak locations were not limited by sampling noise (i.e., the cosmic variance), this effect could be used to determine the cosmic contribution to the dipole term. ",
        "introduction": "\\label{sec:intro}  The motion of the solar system barycenter with respect to the cosmic microwave background (CMB) induces a very large apparent dipole component into the CMB brightness map (3.353 mK, Bennett et al. 1996), and a much fainter, but still readily detectable, quadrupole signature at the $\\sim 1\\, \\mu$K level (e.g., Lineweaver et al. 1996).  The empirically fitted dipole term is subtracted out before computing the spherical harmonics to the residual CMB fluctuations (e.g., Bennett et al. 2003).  The concomitant kinematically induced quadrupole term is considerably smaller than the cosmic quadrupole term and can be subtracted off as well.  One consequence of this operation is that the resultant CMB power spectrum yields no intrinsic dipole term, which would otherwise be of considerable interest for cosmological models.  The intrinsic cosmic dipole is of particular interest as it represents structure on the largest observable scale, and there may well be interesting phenomena to study on the scales of the lowest multipoles. Recent analyses of the first year all-sky WMAP datasets confirm the low quadrupole amplitude (Gaztan{\\~n}aga et al. 2003;  Bennett et al. 2003; Tegmark et al. 2003) as was observed in the earlier COBE DMR maps (Smoot et. al. 1992; Hinshaw et al. 1996). In addition, several papers have pointed out the unlikely alignment of the lowest multipoles for a random Gaussian field, defining the so-called \"Axis-of-Evil\" (Land \\& Magueijo 2005; de Oliveira-Costa et al. 2004; Copi et al. 2004; Bielewicz et al. 2005).   Possible explanations have been suggested, from selection of foreground-free sky (Slosar \\& Seljak 2004) to weak lensing of the cosmic dipole by local structures (Vale 2005).  Future CMB datasets may settle the matter, but a measurement or constraint of the amplitude and direction of the intrinsic cosmic temperature dipole could be essential to reach a final resolution.  In this Letter we discuss another kinematic effect of our motion through the CMB, i.e., the small shift in apparent angular positions due to the aberration of light (Bradley 1729).  The aberration angles are only of order $ \\beta$ (where $\\beta = 0.00123$; Lineweaver et al. 1996), but this leads to a potentially measurable compression (expansion) of the spatial scale in the hemisphere toward (away from) our motion through the CMB.  In turn, this will shift the peaks in the angular temperature power spectrum of the CMB by a factor of order $1 \\pm \\beta$.  The effect amounts to about 1 multipole ($\\ell$) bin at the 5th harmonic of the CMB acoustic spectrum.  In any event, the aberration effect can and should be corrected even though it leads to shifts in the position of the CMB features by only up to $\\sim 4^\\prime$.  If it were not for inherent uncertainties in determining the centroids of the acoustic peaks due to sampling limitations (i.e., the cosmic variance), this effect could, in principle, yield an independent measurement of the kinematic contribution to the dipole and allow a correction to measure the intrinsic power in the cosmic dipole. ",
        "conclusions": "\\label{sec:sum}  We have shown that the aberration of the detected CMB radiation, due to the motion of the barycenter relative to the CMB, can shift the centroids of the acoustic peaks in the forward vs. the backward directions by a measurable amount.  This effect may already be marginally detectable with the WMAP data, and should be detectable with the data to be acquired with the Planck Mission.  To put the magnitude of this aberration effect into perspective, it is anticipated that the Planck CMB maps will be constructed with $\\sim 10^\\prime$ angular bins (Sandri et al. 2004).  Locations in CMB features will be shifted by $2.^\\prime 1$, $3.^\\prime 0$, and $4.^\\prime 2$ at $\\theta^\\prime = 30^\\circ$, $45^\\circ$, and $90^\\circ$, respectively.  If there is, in fact, a significant measurable difference between the locations of the peaks in the CMB power spectrum for the forward and backward directions, then this will serve to demonstrate the expected kinematic effect due to motion of the barycenter.  At present, it does not appear possible to measure this shift with better than about 1 part in 4000 accuracy due to the cosmic variance.  It may be worth noting, nonetheless, that if it {\\em were} somehow possible to achieve a factor of $\\sim$100 better in accuracy, this would be sufficient to allow the kinematic contribution to the dipole term to be determined independently with $\\sim$1\\% accuracy.  This, in turn would enable a determination of the cosmic contribution to the dipole term.  Also, we would like to emphasize that a detection and measurement of the aberration of the CMB is complementary to other effects which arise due to the motion of the solar system. In particular, the intensity quadrupole explored by Kamionkowski \\& Knox (2003) presents an independent observable to constrain the contribution of the kinematic temperature dipole.  Finally, we remark that in the same spirit of searching for direction dependent aberration effects, it might be generally worthwhile to measure the CMB power spectrum over a number ($\\sim$10) of independent regions of the sky and intercompare the results.  The isotropy of the temperature power spectrum has been investigated (Eriksen et al. 2003, Hansen et al. 2004), and regions of excess asymmetry have been reported.  The next generation of isotropic tests should account for the aberration of the CMB to avoid systematic errors over the sky."
    },
    "0601/astro-ph0601590_arXiv.txt": {
        "abstract": "The first FUV ($\\lambda\\lambda$ 1350--1750{\\AA}) spectral imaging observations of the Eridanus superbubble, obtained with the \\emph{SPEAR/FIMS} mission, have revealed distinct fluorescent emission from molecular hydrogen. In this study, the observed emission features were compared with those from a photo-dissociation region model with assumed illuminating stellar fields. The result showed rather high line ratios of ${\\rm I_{1580}/I_{1610}}$, which may imply the existence of high-temperature molecular clouds in the region. The H$_2$ fluorescence intensity showed a proportional correlation with H-$\\alpha$ emission, indicating that the fluorescence and the recombination emission have similar physical origins. ",
        "introduction": "The Eridanus region contains a large nearby interstellar superbubble. This region was discovered as an extended $15^\\circ$ X-ray `hot spot' \\citep{williamson1974,naranan_1976}. Subsequently, the Eridanus Loop was identified as an expanding shell of H{\\footnotesize~I} by \\citet{heiles1976} and was extensively imaged in H-$\\alpha$ emission \\citep{reynolds1998,boumis2001}, revealing a large filamentary shell structure. \\citet{reynolds_ogden1979} argued that Barnard's Loop and the Eridanus shell were both created by one or more supernovae in the Orion OB association and were maintained by the intense stellar winds and ionizing radiation from those stars. {\\it ROSAT} observations \\citep{snowden1995} support the idea that the huge interstellar cavity is filled with several-million-degree gas. Surveys of CO \\citep{magnani1985, aoyama2002} show that the cavity is bounded in part by molecular gas. The CO clouds generally coincide with H{\\footnotesize~I} filaments; however, on a smaller scale, the clouds are located where H{\\footnotesize~I} is weak within the filament. In the FUV spectral range, early \\citep{murthy1993} and improved \\citep{kregenow} spectral observations of diffuse emission from the Eridanus Loop over limited angular fields show that gas at intermediate temperatures, T$\\sim10^5$K, exists between the hot superbubble interior and the surrounding material.  We present the first FUV spectral imaging observations of the entire Eridanus Loop. The data were obtained with the Spectroscopy of Plasma Emission from Astrophysical Radiation (\\emph{SPEAR}), also known as the Far-ultraviolet Imaging Spectrograph (\\emph{FIMS}). These {\\it SPEAR} data show H$_{2}$ fluorescent emission throughout the Eridanus Loop field and provide new insights into the geometry and physical conditions of the molecular environment around this superbubble. ",
        "conclusions": "Our detection of H$_2$ fluorescence in Eridanus is consistent with previous observations in other locations in that fluorescence is found where H$_2$ is exposed to a strong UV radiation field. The correlations among the subregions' fluorescent emission intensity, the H-$\\alpha$ emission intensity, and the color excess value are shown in Figure \\ref{fig3}. The H$_2$ fluorescence intensity shows a strong correlation with both the H-$\\alpha$ emission intensity and the color excess. This suggests that the H-$\\alpha$ emission and hydrogen molecular fluorescence have similar physical origins and are governed globally by a strong and similar UV radiation field, likely arising from the I Orion OB association. We expect that the H$_2$ abundance is  proportional to that of both hydrogen atoms and dust particles, since dust particles are known as the main birth place of H$_2$ in the interstellar medium.  The most conspicuous features in the H-$\\alpha$ emission map shown in Figure \\ref{fig1}, are the filamentary structures elongated along $\\alpha=4^{\\rm h}$(Arc A; Region 3) and $\\alpha=3^{\\rm h}20^{\\rm m}$(Arc B; Region 5). It appears from Figure \\ref{fig3}(a) and (b), and from Figure \\ref{fig1}, that UV photons from the Orion OB association undergo extinction by the Arc A and Arc B filaments, since Region 6 shows relatively low H-$\\alpha$ and H$_2$ fluorescence, while having a color excess value comparable to those of Region 1, 2, and 5. The H$_2$ fluorescent emission from Region 5 (Arc B) is relatively high among regions with comparable dust extinction (Regions 2 and 6), and it is similar to that of Region 4, even though the dust extinction for Region 4 is higher and the line-of-sight distance to the I Orion OB FUV sources is likely smaller. This observation can be explained if the gas in Region 3 (Arc A) is not in the line of sight that extends from the radiation source (the I Orion OB association) to the gas in Region 5 (Arc B), i.e., the distances to Arcs A and B are different and the Eridanus shell structure is open from the I Orion to Arc B. This geometry is consistent with the considerations of \\citet{boumis2001} and \\citet{welsh2005} in that both studies argue that Arcs A and B are part of a complex of individual shells viewed along the same line of sight but at different distances (Arc A $>$500pc, Arc B $\\sim$150pc). The observational results do not rule out Reynolds and Ogden's (1979) description, in that the both Arc A and Arc B can still be exposed to the radiation from the I Orion OB association that causes the H$_2$ fluorescent emission.  The FUV continuum intensity does not show a strong correlation with H-$\\alpha$ emission or H$_2$ fluorescent emission intensities (See Table \\ref{tbl-1}). Instead, the continuum intensity seems closely tied to the geometry of the region. The strong continuum of Region 1 can be attributed to the proximity of Region 1 to the radiation source. Region 3 is remarkable in that it shows both a strong continuum intensity and a strong H$_2$ fluorescent emission intensity.  The observationally derived shell expansion velocity of the Eridanus region, of between 15 and 23 km s$^{-1}$ \\citep{heiles1976,reynolds_ogden1979} is consistent with the survival of the H$_{2}$ in the Eridanus region. The observed ${\\rm I_{1580}/I_{1610}}$ value suggests the existence of high-temperature molecular gas in the region. According to \\citet{draine1983}, when an interstellar shock with $v_s = 25$ km/s is running into molecular gas with a density of $n=1\\times10^2$ cm$^{-2}$, the region in which the gas temperature exceeds 200 K will extend only $2\\times10^{16}$ cm, implying that interstellar shocks are not good at producing large column densities of warm gas. An alternative explanation, for the heating of molecular gas, might be selective UV excitation by UV lines radiated by the hot gas in the bubble, or it may be that other heating mechanisms are able to raise the gas temperature where UV pumping of the H$_{2}$ is taking place. Further studies are needed to clarify the contribution and role of UV radiation, supernova shocks, and other possible heating mechanisms.  In summary, we have presented the first detection of H$_2$ fluorescence in the Eridanus region. The fluorescence emission features, with a PDR spectrum prediction model, imply the possible existence of molecular gas with an excitation temperature as high as 1,000 K, which is about an order of magnitude higher than excitation temperatures generally found for molecular gases in the Galactic disk \\citep{savage1977}. We note that high-excitation temperature H$_2$ has also been detected in high-velocity shocked gas associated with the Monoceros Loop \\citep{welsh2002} and in a galactic translucent cloud \\citep{snow2000}. The observed spectrum of the H$_2$ fluorescence in the Eridanus region is different in that a substantial column density of H$_2$ ($\\sim1.0\\times10^{20}$ cm$^{-2}$) is required to produce the observed fluorescence. Thus, most of the H$_2$ undergoing photon pumping must have a high rotational temperature, if collisional excitation is to account for the observed fluorescent feature ratios. A more rigorous observation and model study are expected to reveal the physical conditions of the molecular gas in the Eridanus supershell."
    },
    "0601/astro-ph0601073_arXiv.txt": {
        "abstract": "The formula for dark energy density derived by Gurzadyan and Xue is the only formula which provides (without a free parameter) a value for dark energy density in remarkable agreement with current cosmological datasets, unlike numerous phenomenological scenarios where the corresponding value is postulated. This formula suggests the possibility of variation of physical constants such as the speed of light and the gravitational constant. Considering several cosmological models based on that formula and deriving the cosmological equations for each case, we show that, in all models source terms appear in the continuity equation. So, on the one hand, GX models make up a rich set covering a lot of currently proposed models of dark energy, on the other hand, they reveal hidden symmetries, with a particular role of the separatrix $\\Omega_m=2/3$, and link with the issue of the content of physical constants. ",
        "introduction": "With the establishment of observational evidence favoring dark energy domination in the recent history of the cosmological expansion, numerous different models were suggested to account for it. For most of them the value of the corresponding density parameter is chosen to fit the observational one. In contrast, Gurzadyan and Xue formula \\cite{GX} for dark energy predicts the observed value for the density parameter of the dark energy. The formula reads \\cite{GX,GX2} \\begin{equation} \\rho _{GX}=\\frac{\\pi }{8}\\,\\frac{\\hbar c}{L_{p}^{2}}\\,\\frac{1}{a^{2}}=\\frac{% \\pi }{8}\\,\\frac{c^{4}}{G}\\,\\frac{1}{a^{2}},  \\label{rhoLambda} \\end{equation}% where $\\hbar $ is the Planck constant, the Planck length is $L_{p}=\\left( \\frac{\\hbar G}{c^{3}}\\right) ^{\\frac{1}{2}}$, $c$ is the speed of light, $G$ is the gravitational constant, $a$ is the scale factor of the Universe. Following the original idea by Zeldovich \\cite{zelvac}, it corresponds to the cosmological term \\begin{equation} \\Lambda _{GX}=\\frac{8\\pi G\\rho _{GX}}{c^{2}}=\\frac{8\\pi G}{c^{2}}\\frac{\\pi }{% 8}\\,\\frac{c^{4}}{G}\\,\\frac{1}{a^{2}}=\\pi ^{2}\\left( \\frac{c}{a}\\right) ^{2}. \\label{Lambda} \\end{equation}% Therefore, to keep it constant, the speed of light should vary with cosmological expansion as $c\\propto a$. One can consider other possibilities as well, assuming the presence of some fundamental physical quantity and admitting variation of basic physical constants \\cite{Ver05} in the spirit of Dirac approach. Since the Planck constant does not appear in (\\ref% {rhoLambda}), these can be the gravitational constant and the speed of light.  In the literature there is a long-standing discussion how to understand possible variability of physical constants, linking this point to the meaning of world constants and units in physics \\cite{Duf02}. Clearly experimentally only dimensionless constants variation can be detected, but the meaning and the content of the underlying physical theory changes depending on the issue which dimensionful quantity varies.  In fact, models with varying physical constants are among the currently discussed ones (e.g. \\cite{B}). However, models based on the formula (\\ref% {rhoLambda}) are very different from various phenomenological models of dark energy, phantoms etc., because of the underlying \\emph{empirical basis}, namely agreement between predictions for dark energy density parameter and the value, following from the array of recent observations.  The classification of various cosmological models following from GX-formula is given in \\cite{Ver05}. Furthermore, it is shown \\cite{VY06b} that GX models provide indeed good fit to supernovae and radio galaxies data. These models contain hidden invariance, the separatrix dividing cosmological solutions into two classes: singular Friedmannian-type and non-singular with vanishing initial density \\cite{VY06b},\\cite{VY06a}. In this paper we provide cosmological equations for the simplest cases discussed in \\cite% {Ver05} and generalize them, including radiation field, and compare them to the standard cosmological model.  In the section 2 cosmological framework with varying physical constants is formulated. In the section 3 we derive cosmological equations for particular cases, discussed in \\cite{Ver05}. We provide analytical solutions for energy density in all models, and also in some models for the scale factor. Conclusions follow in the last section. ",
        "conclusions": "In this paper we explored various cosmological models with Gurzadyan-Xue dark energy formula (\\ref{rhoLambda}), assuming also variation of physical constants such as the speed of light and the gravitational constant. It is shown that unlike standard Friedmann models, various interesting possibilities appear. In particular in the model 1 the continuity equation is reduced to the energy transfer from dark energy component into the usual matter. For some cases (models 1 and 4), the source term in the continuity equation can be neglected for sufficiently large scale factors (or sufficiently large present density of matter), and the discussion of the dynamics in this asymptotic cases can be given in terms of the standard cosmology.  Solutions for the density in terms of the scale factor for all models with pressureless matter and dark energy contain separatrix which divides solutions in two classes. The first class is analogous to the standard Friedmann solutions with singularity in the beginning of expansion. The second class, in contrast, does not contain singularity; solutions start with vanishing density and non-zero scale factor.  We also considered generalization of GX models, including radiation, assuming that dynamics of the latter is not influenced by the variation of constants and by dark energy. Interestingly, for the model 1 we obtained the behavior qualitatively similar to models without radiation. For the rest of models separatrix in the sense discussed above is absent and all solutions belong to the second class, namely non-Friedmannian solutions. The density in these solutions may be larger in the past, but if so, it has a maximum and all solutions start with vanishing density in radiation-dominated epoch. However for models 2 and 3 with radiation the same value of density parameter $\\Omega _{m}=2/3$ that parametrizes separatrix in simple GX models, corresponds to constant mass density of pressureless matter in matter dominated epoch. For larger density parameter the mass density decreases with expansion and for smaller density parameter it increases.  In fact, not all GX models with non-Friedmannian behavior ($\\Omega_m<2/3$) pass age constraint and have deceleration parameter which fits the value followed from observations. At the same time model IV with $\\Omega_m\\geq 2/3$ pass both constraints and represent a good candidate for dark energy model \\cite{VY06b} which in the limit of vanishing energy transfer from dark energy to the usual matter has $a^{-1}$ term in Friedmann equation. In the same way models I and III fit these constraints for density parameter slightly smaller than the separatrix value. This circumstance once more points out on a special role of the separatrix in GX models.  Notice another interesting feature of GX models with radiation. It is well known that Friedmannian models with zero curvature have arbitrary scale factor. The same holds also for simple GX models (model 2). However, this degeneracy breaks down in GX models with radiation, where the scale factor can always be computed in terms of present density parameters, and present values of (varying) physical constants.  We are thankful to the referees for helpful comments."
    },
    "0601/hep-th0601028_arXiv.txt": {
        "abstract": "{A wide variety of vacua, and their cosmological realization, may provide an explanation for the apparently anthropic choices of some parameters of particle physics and cosmology. If the probability on various parameters is weighted by volume, a flat potential for slow-roll inflation is also naturally understood, since the flatter the potential the larger the volume of the sub-universe. However, such inflationary landscapes have a serious problem, predicting an environment that makes it exponentially hard for observers to exist and giving an exponentially small probability for a moderate universe like ours. A general solution to this problem is proposed, and is illustrated in the context of inflaton decay and leptogenesis, leading to an upper bound on the reheating temperature in our sub-universe. In a particular scenario of chaotic inflation and non-thermal leptogenesis, predictions can be made for the size of CP violating phases, the rate of neutrinoless double beta decay and, in the case of theories with gauge-mediated weak scale supersymmetry, for the fundamental scale of supersymmetry breaking.}  \\end{center} \\end{titlepage} ",
        "introduction": "Our complicated world is based on hundreds of elements that have different properties.  The hundreds of elements reflect the existence of hundreds of stable nuclei. The variety of stable nuclei requires fine choices of parameters such as $\\alpha_{\\rm QCD}, \\alpha_{\\rm QED}, m_u, m_d$ and $m_e$. Even a deviation of a few \\% in $\\alpha_{\\rm QCD}$ from the current value could destabilize or stabilize such nuclei as ${}^2 {\\rm H}$, ${}^2 {\\rm He}$ and di-neutrons, completely changing the thermonuclear processes, chemical abundances of the universe and lifetime of stars \\cite{Barrow}. Such fine tuning cannot be explained by the naturalness principle. Rather, it is easily explained by a combination of the many-universe idea (a cosmological diversity in the choice of theories and their parameters) with anthropic selection \\cite{manyA,manyB}. It is only in the sub-universes with fine-tuned parameters that observers made of many species of atoms exist. This combination also provides a simple solution to both the why-small and why-now problems of the cosmological constant \\cite{CC1,CC2}.  Once we accept that the combination of many-universes with anthropic selection plays a role in determining the values of observed parameters, then we need to change our way of thinking. The combination not only confirms that the observed values of parameters are consistent with the existence of observers, but, because the notion of naturalness is replaced by probability, it changes to some extent the questions that should be asked. The fine tuning of parameters and initial conditions of inflation is one prominent example: if the probability distribution is weighted by the volume of the sub-universes, such a fine tuning is rather probable, as long as theory space allows the inflaton mass to scan \\cite{Vilenkin95}, or provided appropriate initial conditions are realized somewhere in the universe \\cite{LindeChaotic}. We call such a space of theories containing lots of inflationary regions an inflationary landscape.  The problem of inflationary landscapes \\cite{FHW,GV} is that we are no longer able to choose the parameters of the inflation model by hand to fit the observed data; rather these parameters (and consequently the normalization of the density perturbations) are predicted by a combination of statistics from the landscape, cosmological dynamics and anthropic conditions \\cite{manyB}. In particular, it is generically true that as the inflaton mass becomes smaller, so the inflated volume becomes exponentially larger. This exponential behaviour of the volume-weighted probability distribution implies not only that a small inflaton mass is very probable \\cite{Vilenkin95}, but also that the inflaton mass will become as small as possible until it is prevented by some anthropic reason. Thus, a naive prediction of inflationary landscapes would be that the inflaton mass is so small that the environment is {\\it exponentially} hard for observers to exist \\cite{FHW} because, for example, the density perturbations are too small or too large (this was called ``$\\sigma$ problem'' in \\cite{FHW}), the baryon asymmetry is too small or there is not enough matter. Whatever the reason, this picture apparently does not lead to our universe.\\footnote{After understanding astrophysics and biology better, however, the naive prediction may turn out to be the solution to Fermi's paradox ``Where is everybody?'' \\cite{everybody}.} Since the slow-roll volume factor selects vacua with smaller and smaller inflaton mass, we call this the inflation runaway problem.  The $\\sigma$ problem may be solved \\cite{FHW,GV} if the density perturbations arise from fluctuations of light fields other than the inflaton \\cite{curvaton,modulated}. Even then, the inflaton mass has to be predicted at a moderate value, or otherwise, there could be\\footnote{When the inflaton mass does not affect late time cosmology except for the number of e-foldings, as in the hybrid inflation model in \\cite{FHW}, such a problem does not exist.} the same problem as above, for instance, in the baryon asymmetry or dark matter abundance. Thus, this is quite a generic problem of inflationary landscapes. Note, however, that the probability distribution on observable parameters depends sensitively on the boundary values of unobservable parameters and amplitudes being scanned. Thus, the distribution is not exponentially sensitive to observables \\cite{Linde-curvaton} in some landscapes where the observable--unobservable correlation is sufficiently week \\cite{FHW}.  In this article, we present an idea of how to stop the runaway behaviour of the inflaton mass, along with its concrete implementation. In section \\ref{ssec:rev} we introduce the runaway problem, and in section \\ref{ssec:idea} we point out a caveat that allows a solution. It is crucial to distinguish anthropically relevant parameters from fundamental parameters, and further \\begin{enumerate} \\item [(a)] to have a sharp function, or threshold behaviour, of the former in terms of the latter, and \\item [(b)] to have an anthropic reason that exponentially disfavours some range of relevant parameters. \\end{enumerate} The runaway problem can be overcome when both (a) and (b) are satisfied. In section \\ref{ssec:inflatondecay}, we see that (a) can be realized as the reheating temperature $T_R$ depends upon the inflaton mass. If the inflaton decays dominantly to two heavy particles, $T_R$ drops rapidly as the inflaton mass becomes close to twice of the mass of the heavy particle. The inflaton mass is predicted to be close to twice the mass of the heavy particle. In this article we assume that only the inflaton mass and the cosmological constant are scanned in the landscape, not the mass of the heavy particle.\\footnote{We do so as a first attempt to address the runaway problem in the simplest and easiest situation. There may be room to lift some of these strong assumptions, but this is beyond the scope of this article.} Section \\ref{sec:exp} explains how $T_R$ can be a relevant parameter. If the baryon asymmetry is generated by thermal leptogenesis, a very low reheating temperature leads to a low baryon asymmetry and, if the baryon asymmetry is too low, galaxies cannot form before protons decay. Thus, too low a $T_R$ is exponentially disfavoured. Therefore, both (a) and (b) can be realized, and the runaway problem is solved. In section \\ref{sec:neutrino} we study the possibility that the heavy particles produced by inflaton decay are identified with one of the right-handed neutrinos, hoping for constraints and predictions in the neutrino sector. The last section is devoted to conclusions and discussion, together with a summary of our predictions. The appendix investigates one of the uncertainties in the analysis in sections \\ref{sec:exp} and \\ref{sec:neutrino}, the effect of evaporation of particles from proto-galaxies. ",
        "conclusions": "\\label{sec:concl}  If the fundamental theory allows more than one vacuum, and if many vacua are realized cosmologically, then cosmology may play an important role in choosing the low energy theory and its parameters. In inflationary landscapes, where models and parameters of inflation are scanned cosmologically, sub-universes with a small inflaton mass dominate the volume of the universe, making slow-roll inflation a probable consequence. The problem of inflationary landscapes \\cite{FHW,GV}, however, is that the probability distribution tends to be exponentially sensitive to some late-time cosmological parameters, and hence does not correctly predict their observed values. We call this the inflation runaway problem.  Landscapes and many-universes are about the world outside the horizon of our universe, and we usually think that no sign of them can be seen. One also usually considers that the anthropic conditions depend on only a limited number of relevant parameters, most of which have already been measured, so that there are no predictions. Without a well-understood top-down tool to figure out the statistic distribution of vacua on a landscape, it seems almost impossible to determine theories or parameters that are yet to be tested in experiments. It is important to note, however, that inflationary landscapes have a generic prediction: the probability distribution is exponentially sensitive to some parameters. This is where the runaway problem comes from. The problem is so severe that there may be very few possible solutions, from which indications may be obtained for particle physics to be explored in the future.  Some ideas have already been proposed to evade this problem. For instance, eternal inflation of false-vacuum type or chaotic-type may strongly select only one inflationary region in a landscape \\cite{FHW,GV}. The observational consequences for the curvature of the universe and for lower multipoles of CMB spectrum are discussed in \\cite{Susskind05} and references therein. Maybe the fundamental theory of phsycis has an inflationary landscape carefully designed to avoid the volume factor exponentially sensitive to some parameters \\cite{FHW,Linde-curvaton}. If such landscapes involve chaotic type inflation, we may expect tensor perturbations. A fundamental cut off scale parametrically larger than the observed Planck scale may also have something to do with the apparent fine tuning problems of certain low-energy supersymmetric theories \\cite{IIY}. An alternative is to have density perturbations generated by a curvaton or by modulated reheating \\cite{FHW,GV}, whose observational consequences are discussed in the literature \\cite{curvaton,nonGauss,Linde-curvaton}.  In this article, we have proposed a new solution to the runaway problem, which is connected to the {\\em particle physics} of inflaton decay and baryogenesis. In spite of the exponential distribution from the inflationary landscapes, observables are predicted successfully in the middle of the allowed anthropic windows. The key ingredient is to distinguish between parameters describing inflation and those that are directly observed. The latter are not necessarily given by simple power-law functions of the former. For instance, the reheating temperature drops sharply as the inflaton mass approaches the threshold for two body decay. In this case, the inflationary landscape predicts the inflaton mass to be close to twice the mass of the threshold particle.  A sufficiently sharp threshold behaviour for $T_R$ requires small couplings in operators relevant to reheating, and hence a low $T_R$ in our universe that depends on the cosmological scenario: $T_R \\simlt 10^{12}$ GeV in the thermal leptogenesis scenario of section \\ref{sec:exp}, and $10^6 \\; \\GEV \\simlt T_R \\simlt 10^7 \\; \\GEV$ in the inflaton-decay leptogenesis scenario of section \\ref{sec:neutrino}. Chaotic inflation is assumed. In both sections, the anthropic limit was set by requiring that bremsstrahlung and runaway hierarchical cooling of proto-galaxies start before protons decay. Although we have ignored various uncertainties, the upper bound obtained on $T_R$ illustrates a completely new way to obtain information on particle cosmology from anthropic reasoning.  In section \\ref{sec:neutrino} the particle produced by inflaton decay was identified with a right-handed neutrino, allowing the inflaton mass, and some other parameters including the normalization of density perturbations, to be set by physics of the neutrino sector. This provides a mechanism for understanding the values of inflation parameters in terms of observable particle physics. In this particular scenario, there are several predictions that again demonstrate how this line of anthropic argument can lead to observable tests. One CP violating phase of the neutrino sector must be of order 0.1 or larger. The low energy neutrino mass matrix is forced to be essentially rank 2, so that the rate for neutrinoless double beta decay will be constrained by a particular relationship involving the neutrino mixing angles and the observed neutrino masses. Furthermore, in supersymmetric theories with gauge mediation, dark matter is predicted to be gravitinos with mass in the range 20--200 MeV. This corresponds to a fundamental scale of supersymmetry breaking of $10^8$--$10^9$ GeV, which can be tested by precise measurements of the superpartner spectrum and by the nature of the decay of the next-to-lightest superpartner \\cite{DESY}.  It should be noted that the above predictions are based on an assumption that the anthropic limit is set by the comparison between the rate of bremsstrahlung cooling and proton decay. When various other anthropic conditions are imposed, some of which are listed in footnotes in section \\ref{sec:exp}, the anthropic limit in the $Y_B$--$Q$ plane may be stronger, and the upper bounds on $T_R$ may become weaker. Including other uncertainties in galaxy formation and in the anthropic conditions, the above predictions may be modified. There is no doubt that further studies in astrophysics will be important in making the predictions more precise and reliable. As a step in this direction, we discuss possible effects of particle evaporation from proto-galaxies in the appendix, concluding that the upper bound on $T_R$ will not be greatly changed, partly because,  in scenarios in section \\ref{sec:exp} and \\ref{sec:neutrino}, $Y_B$ is highly sensitive to $T_R$.  As the astrophysics and anthropic conditions become better understood, the upper bounds of the reheating temperature will change numerically. But, the idea for stopping the runaway by a combination of threshold effects and exponential anthropic factors will survive; extra anthropic conditions  would only strengthen our conclusion that the exponential distribution of inflationary landscapes can be tamed. No matter how sophisticated and complicated the astrophysical analysis becomes, the upper bound on the reheating temeperature persists as long as the threshold behaviour of the inflaton decay plays a crucial role."
    },
    "0601/astro-ph0601114.txt": {
        "abstract": "% OK OK  We use data obtained with the {\\it Far Ultraviolet Spectroscopic Explorer} ({\\it FUSE}) to determine the interstellar abundances of D\\,I, N\\,I, O\\,I, Fe\\,II, and \\hmol~along the sightlines to WD\\,1034$+$001, BD$+$39\\,3226, and TD1\\,32709. Our main focus is on determining the D/H, N/H, O/H, and D/O ratios along these sightlines, with log $N$(H) $>$ 20.0, that probe gas well outside of the Local Bubble. {\\it Hubble Space Telescope} ({\\it HST}) and {\\it International Ultraviolet Explorer} ({\\it IUE}) archival data are used to determine the H\\,I column densities along the WD\\,1034+001 and TD1\\,32709 sightlines, respectively. For BD$+$39\\,3226, a previously published $N$(H\\,I) is used. We find (D/H)$\\times10^5$ = 2.14 $\\pm~^{0.53}_{0.45}$, 1.17 $\\pm~^{0.31}_{0.25}$, and 1.86 $\\pm~^{0.53}_{0.43}$, and (D/O)$\\times10^2$ = 6.31 $\\pm~^{1.79}_{1.38}$, 5.62 $\\pm~^{1.61}_{1.31}$, and 7.59 $\\pm~^{2.17}_{1.76}$, for the WD\\,1034$+$001, BD$+$39\\,3226, and TD1\\,32709 sightlines, respectively (all 1$\\sigma$). The scatter in these three D/H ratios exemplifies the scatter that has been found by other authors for sightlines with column densities in the range 19.2 $<$ log $N$(H) $<$ 20.7. The D/H ratio toward WD\\,1034$+$001 and all the D/O ratios derived here are inconsistent with the Local Bubble value and are some of the highest in the literature. We discuss the implications of our measurements for the determination of the present-epoch abundance of deuterium, and for the different scenarios that try to explain the D/H variations. We present a study of D/H as a function of the average sightline gas density, using the ratios derived in this work as well as ratios from the literature, which suggests that D/H decreases with increasing gas volume density. Similar behaviors by other elements such Fe and Si have been interpreted as the result of depletion into dust grains. ",
        "introduction": "The present day abundance ratio of deuterium to hydrogen places important constraints on Big Bang nucleosynthesis (BBN) and the chemical evolution of galaxies. Because deuterium is only produced in appreciable amounts in primordial BBN and destroyed in stellar interiors, the measurement of D I/H I in the interstellar medium (ISM) places a lower limit on the primordial abundance of deuterium. In addition, by comparing the ISM abundance of deuterium to its abundance in high-redshift intergalactic gas we should be able to better understand the effects of astration and chemical evolution of galaxies. Measurements of the D/H ratio in intervening clouds of gas seen toward distant quasars have yielded a range of values D/H = (1.65 -- 4.0)$\\EE{-5}$ \\citep[and references therein]{2001ApJ...552..718O,2001ApJ...560...41P,2002ApJ...565..696L}. \\citet{2003ApJS..149....1K} measured D/H = (2.42 $\\pm~^{0.35}_{0.25}$)$\\times10^{-5}$ toward Q1234$+$3047 ($z$ = 2.526). These authors calculate the value they believe is the best estimate of the primordial D/H abundance, D/H$_{\\rm prim}$ = (2.78 $\\pm~^{0.44}_{0.38}$)$\\times10^{-5}$ (1$\\sigma$ errors in the mean), by taking the weighted mean of five D/H measurements toward QSOs ranging from $z$ = 2.079 -- 3.572. This value is in good agreement with that determined from the cosmic microwave background measurements performed with $WMAP$~and previous missions \\citep{2003ApJS..148..175S}. \\citet{2004ApJS..150..387S} found D/H = (2.2 $\\pm$ 0.7)$\\times10^{-5}$ for Complex C, a high velocity cloud falling into our galaxy, which has low metallicity and has presumably experienced more stellar processing than the gas seen toward QSOs.  Measurements of D/H in the local ISM have been made with {\\it Copernicus} \\citep[e.g.][]{1973ApJ...186L..95R}, \\hst, \\citep[e.g.][]{1995ApJ...451..335L},~\\imaps,~\\citep{1999ApJ...520..182J,2000ApJ...545..277S}, and more recently~\\fuse~\\citep[e. g.][and references therein]{2002ApJS..140....3M}. A nearly constant ratio of D/H = $(1.5~\\pm~0.1)\\EE{-5}$ (1$\\sigma$~on the mean) has been obtained in the Local Interstellar Cloud (LIC) by \\citet{1998SSRv...84..285L}; recent measurements inside the Local Bubble \\citep[LB, log $N$(H\\,I) $\\le$ 19.2][]{1999A&A...346..785S} appear to be consistent with a single value for D/H in the LB \\citep{2002ApJS..140....3M,2003ApJ...587..235O}. From a compilation of measurements from the literature \\citet{2004ApJ...609..838W} derive (D/H)$_{\\rm LB}$ = (1.56 $\\pm$ 0.04)$\\times10^{-5}$. Using (D/O)$_{\\rm LB}$ = (3.84 $\\pm$ 0.16)$\\times10^{-2}$ from \\citet{2003ApJ...599..297H} and (O/H)$_{\\rm LB}$ = (3.45 $\\pm$ 0.19)$\\times10^{-4}$ from \\citet{2005ApJ...625..232O} one can derive indirectly (D/H)$_{\\rm LB}$ = (1.33 $\\pm$ 0.09)$\\times10^{-5}$. The direct and indirect determinations of (D/H)$_{\\rm LB}$ are only consistent when one considers the 2$\\sigma$ uncertainties in both quantities. A detailed discussion of the possible causes for this discrepancy can be found in \\citet{2003ApJ...599..297H}.  Outside the Local Bubble, however, there still is not a consistent picture of the D/H behavior. Measurements performed along sightlines probing gas outside the LB % performed by \\citet{1979ApJ...229..923L,1983ApJ...264..172Y,1999ApJ...520..182J,2000ApJ...545..277S} suggest variations of the interstellar D/H ratio beyond the Local Bubble, at the distance of a few hundred parsecs. D/H varies for 19.2 $\\le$ log $N$(H\\,I) $\\le$ 20.7 \\citep[e.g.][]{1999ApJ...520..182J,2000ApJ...545..277S,2002ApJS..140...37F} and apparently remains constant for log $N$(H\\,I) $\\ge$ 20.7, albeit with a lower value. The possibility of a trend towards low D abundance at high H column densities was noted in the D/O survey by \\citet{2003ApJ...599..297H}. \\citet{2004ApJ...609..838W}, using additional measurements and values in the literature, confirmed the trend for D/H. So far, only five measurements for sightlines with log $N$(H\\,I) $\\ge$ 20.7 have been published. \\citet{2003ApJ...586.1094H}, \\citet{2004ApJ...609..838W} and \\citet{2005guillaume} measured D/H along five extended lines of sight ($d \\ge$ 500 pc). The weighted mean of these five measurements yields D/H = (0.87 $\\pm$ 0.08)$\\times10^{-5}$, in clear disagreement with the LB value.  The value of the present-epoch Milky Way abundance of deuterium, (D/H)$_{\\rm PE}$, is still an issue subject to debate. Two opposite explanations have emerged. \\citep{2003ApJ...599..297H} argue for a value of the current-epoch D/H ratio lower than D/H$_{\\rm LB}$, on the basis of the generally low values found for more distant sightlines. On the other hand, Linsky et al. (2005, in prep) argue that an important fraction of deuterium along many sightlines is trapped in a population of grain material and hence the current-epoch deuterium abundance (gas plus grains) is (D/H)$_{\\rm PE}$ $\\ge$ (2.19 $\\pm$ 0.27)$\\times10^{-5}$. The astration factors, (D/H)$_{\\rm prim}$/(D/H)$_{\\rm PE}$, derived from both of these scenarios ($\\sim$4 and $\\sim$1.2, respectively) challenge current models of galactic chemical evolution.  Because the O/H variations in the diffuse ISM from the LB out to 1 Kpc, are small, consistent with the measurement uncertainties \\citep{1998ApJ...493..222M,2003ApJ...591.1000A,2004ApJ...613.1037C,2005ApJ...625..232O}, O\\,I has been used as a proxy for H\\,I \\citep[for a more detailed discussion see][and references therein]{2003ApJ...599..297H}. Thus D/O and D/H are expected to trace each other in the diffuse ISM. However, how well this works at high column densities is a subject of discussion \\citep[][and $\\S$\\ref{litvalues} of this paper]{2006guillaume}.  At the present time, definitive conclusions are prevented by the limited number of D/H and D/O measurements along high column density sightlines. In this work we present D/H and D/O measurements along three sightlines (WD\\,1034$+$001, BD$+$39\\,3226, and TD1\\,32709) with log $N$(H\\,I) $\\ge$ 20.0. Using data obtained with $FUSE$, $IUE$, and the Goddard High Resolution Spectrograph (GHRS, onboard {\\it HST}) we derive column densities of atomic and molecular species which are then used to determine important ratios (D/H, N/H, O/H, etc.) that can be compared to values in the literature.  This paper is organized as follows. The three targets used for the analyses are described in \\S\\ref{targets}. The observations and data processing are presented in \\S\\ref{obs}. In \\S\\ref{ana} we determine the column densities of H\\,I, D\\,I, N\\,I, O\\,I, Fe\\,II, and \\hmol~along the WD\\,1034$+$001, BD$+$39\\,3226, and TD1\\,32709 sightlines. Lower limits to the column densities of other elements are also reported. The D/H, D/N, D/O, N/H, O/H, and O/N ratios for the three sightlines are presented in $\\S$\\ref{results} and discussed in \\S\\ref{discussion}, where evidence for above average D abundance is presented for the WD\\,1034$+$001 and TD1\\,32709 sightlines. The D/O ratios for the three sightlines are substantially higher than the mean LB value, and thus they considerably increase the scatter in the published D/O ratios. Implications of the derived ratios for hypothetical mechanisms, such as dust depletion and infall of metal poor gas, that try to explain the D/H variations in the Galactic disk are considered in \\S\\ref{implica}. Our derived D/H, D/O, and O/H ratios are compared with previously published ratios in \\S\\ref{litvalues}. In $\\S$\\ref{iron} we consider the relationship between D/O and Fe/O. In \\S\\ref{nh_disc} we perform a study of D/H as a function of the average sightline gas density ($N$(H)/$d$). Our study indicates that D/H is constant up to densities of 0.10 cm$^{-3}$ and decreases with increasing density after this point, similarly to what has been observed for other elements such as Fe and Si. We summarize our findings in \\S\\ref{summary}. All uncertainties are quoted at the 1$\\sigma$ level unless stated otherwise. ",
        "conclusions": "\\label{discussion}   \\subsection{WD\\,1034$+$001} \\label{wd1034_disc}  The WD\\,1034$+$001 sightline has one of the highest D/H ratios reported in the literature, D/H = (2.14 $\\pm~^{0.53}_{0.45}$)$\\times10^{-5}$, 1.29$\\sigma$ above D/H$_{\\rm LB}$ derived by \\citet{2004ApJ...609..838W}. However, its O/H and N/H ratios are consistent at the 1$\\sigma$ level with previous measurements \\citep{1998ApJ...493..222M,1997ApJ...490L.103M}. A D/H ratio consistent with the LB value ($\\sim$1.56$\\times10^{-5}$), would require log$N$(H\\,I) = 20.21. In such case, the O/H and N/H ratios (2.47 $\\pm~^{0.80}_{0.63}$)$\\times10^{-4}$ and (5.66 $\\pm~^{2.06}_{1.60}$)$\\times10^{-5}$, respectively) would still be consistent with the values derived by \\citet{1998ApJ...493..222M} and \\citet{1997ApJ...490L.103M}, while not affecting the high D/O ratio (the same percentage uncertainties were used for the required $N$(H\\,I) as the ones derived in this work). In this scenario $\\sim$ 34\\% of the H\\,I would have been missed in our analysis. It seems unlikely that $N$(H\\,I) along this sightline could be underestimated by such a large factor. On the other hand, a value of log$N$(D\\,I) = 15.26 (38\\% smaller than our derived value) together with our derived $N$(H\\,I) would also lead to D/H consistent with the LB value, without affecting the O/H and N/H ratios. In this scenario, using the same uncertainties for the required $N$(D\\,I) as the ones derived in this work, we would find D/O = (4.57 $\\pm~^{1.43}_{1.16}$)$\\times10^{-2}$. This is lower than our derived value of (6.31 $\\pm~^{1.79}_{1.38}$)$\\times10^{-2}$, and consistent with the Local Bubble value derived by \\citet{2003ApJ...599..297H}. We derived $N$(D\\,I) along this sightline using three different methods, AOD, COG, and PF (see Table \\ref{atomicdata}), and it is unlikely that $N$(D\\,I) has been overestimated.  Linsky et al. (2005, in prep) looked at correlations between the D/H ratio and other parameters, such as the depletions of Fe and Si, and the kinetic temperature of \\hmol~($T_{01}$), in order to try to understand what is the total deuterium abundance in the local Galactic disk. They found that sightlines with higher Fe and Si depletions correspond in general to lower D/H ratios and that sightlines with high  $T_{01}$~have typically higher D/H ratios.  The Fe/H ratio for this sightline, Fe/H = (1.07 $\\pm~^{0.33}_{0.27}$)$\\times10^{-6}$ (corresponding to log (Fe\\,II/H\\,I) = $-$5.97), is lower than would be expected for this D/H ratio according to Figure 3 of Linsky et al. (2005, in prep). Using the solar abundance of Fe from \\citet{2004astro.ph.10214A}, [Fe/H]$_{\\odot}$ = $-$4.55, the depletion of Fe along this sightline is D(Fe) = $-$1.42. This is an unusual sightline, for which non-stellar absorption by C\\,IV, N\\,V, O\\,VI, Si\\,IV, S\\,III, and S\\,IV, has been detected \\citep[][and this work]{1985ApJ...292..477S,1995A&A...298..567W}, in addition to the large elliptic halos seen in emission (see $\\S$\\ref{wd1034_target}). %{\\bf Data was obtained in TTAG, inspection of LWRS data shows no OVI emission, Birgit obtained 2sigma upper limits on the emission, can this be used to conclude something about photoionization along this sightline?}  %log$N$(H\\,I) = 19.93 for D/H = 1.56 implies O/H = 4.65 and N/H = 1.06$\\times10^{-4}$. On the other hand if D is overestimated then require log$N$(D\\,I) = 15.26 for D/H =1.56, D/O = 4.6 still high. On the other hand for D/O = 3.84 need D = 15.18 which will lead to D/H = 1.3.   The presence of highly ionized species, the high D/H and lower than expected Fe/H, and the typical O/H and N/H ratios are consistent with a scenario in which D is released from the grains, but not much Fe. Perhaps the high radiation field of WD\\,1034$+$001 plays a special role for this sightline. A similar mechanism has been proposed by Linsky et al. (2005, in prep) to explain the high D/H and low Fe/H ratios along the sightline to $\\gamma^2$ Vel.    \\subsection{BD$+$39\\,3226} \\label{bd39_disc}  For the BD$+$39\\,3226 sightline the D/H ratio is only slightly inconsistent with the LB value (1.15$\\sigma$ below). For this sightline, D/H and Fe/H ((1.17 $\\pm~^{0.34}_{0.28}$)$\\times10^{-6}$) follows the correlation derived by Linsky et al. (2005, in prep) for these quantities. Even though interstellar O\\,VI is detected along this sightline, we found no evidence of other highly ionized species (such as C\\,IV, N\\,V, Si\\,IV, S\\,III, or S\\,IV) in either the {\\it FUSE} or {\\it IUE} data for this star. In this case the separation between stellar and interstellar absorption is high enough ($\\sim$255 \\kms) to allow for a definitive detection, should an interstellar feature be present. It is possible that along this sightline O\\,VI is formed in evaporative interfaces between cool clouds and the hot and diffuse gas in the Local Bubble. This mechanism has been invoked by \\citet{2005ApJ...622..377O} to explain the presence of O\\,VI along the sightlines to several white dwarfs inside the Local Bubble. D/O along this sightline, (5.62 $\\pm~^{1.61}_{1.31}$)$\\times10^{-2}$, is not consistent at the 1 $\\sigma$ level with the D/O ratio derived for the LB by \\citet{2003ApJ...599..297H}, D/O = (3.84 $\\pm$ 0.16)$\\times10^{-2}$. This is however, in part due to the low O/H ratio derived here for this sightline, O/H = (2.09 $\\pm~^{0.72}_{0.58}$)$\\times10^{-4}$, which is 1.86$\\sigma$ below the O/H ratio derived by \\citet{1998ApJ...493..222M}. This O/H ratio is lower than what has been derived in similar studies (see $\\S$\\ref{litvalues}) and might be indicative of variations of the O/H ratio. Using the O/H ratio derived by the authors above yields D/O $\\sim$ 3.4$\\times10^{-2}$, in closer agreement with the D/O LB value reported above.    \\subsection{TD1\\,32709} \\label{td1_disc}  The TD1\\,32709 sightline presents a D/H ratio consistent at the 1$\\sigma$ level with the LB, while D/O is inconsistent within the quoted uncertainties at the 2$\\sigma$ level. O/H is lower than the value derived by \\citet{1998ApJ...493..222M} by 1.09$\\sigma$. If we assume that the D/H and D/O ratios are high due to $N$(D\\,I) being overestimated then we would need log$N$(D\\,I) = 15.22 to bring D/H closer to the LB value. Using the same uncertainties for the required D\\,I as the ones derived in this work (0.05 dex), the D/O ratio would still be high at (6.31 $\\pm~^{1.80}_{1.47}$)$\\times10^{-2}$, even with this reduced value for $N$(D\\,I). It is unlikely that $N$(H\\,I) has been over or underestimated, as a lower $N$(H\\,I) would lead to an even larger D/H ratio, while a larger $N$(H\\,I) would lead to an even smaller O/H ratio. Increasing $N$(O\\,I) so that O/H agrees with the \\citet{1998ApJ...493..222M} value still leads to a high D/O ratio, (5.43 $\\pm~^{1.55}_{1.26}$) $\\times10^{-2}$ (assuming that the increased $N$(O\\,I) has the same uncertainties as the one derived in this work). The only possibility to reconcile the D/H and D/O ratios along this sightline with the LB values would be for D and O to be simultaneously over and underestimated by 1.6 $\\sigma$ and 1.5 $\\sigma$, respectively. However this is unlikely given the number of transitions and methods used to determine $N$(D\\,I) and $N$(O\\,I).     For this sightline the Fe/H ratio ((0.83 $\\pm~^{0.30}_{0.24}$)$\\times10^{-6}$) is lower than would be expected for the D/H ratio according to Figure 3 of Linsky et al. (2005, in prep). We searched the {\\it IUE} data for this star for absorption by highly ionized species such as C\\,IV, N\\,V, and Si\\,IV. However because the separation between stellar and interstellar absorption is low ($\\sim$13 \\kms) we cannot determine if the detected absorptions are of stellar or interstellar origin (or both), similarly to what we had concluded for the O\\,VI absorption detected in the {\\it FUSE} data (see $\\S$\\ref{uv0904atomic}).      \\subsection{Implications for models that explain D/H variations} \\label{implica}  The high D/H ratios derived for the WD\\,1034$+$001 and TD1\\,32709 sightlines increase the sample of sightlines for which a high D/H ratio has been found. These measurements strengthen the argument that these high ratios can not be due to some systematic problem associated with the different analyses (see also discussion below) or that they are peculiar cases. In particular, it has been argued that the high D/H ratios found in the ISM could be due to systematic effects that affect the $N$(H\\,I) measurements \\citep{2005guillaume}. Because the D/O ratios derived for these two sightlines are also high and independent of $N$(H\\,I), the high D/H ratios are likely due to a high abundance of D and not to a systematic problem with the $N$(H\\,I) determination (see also discussion in $\\S$\\ref{litvalues}). In addition, the high D/O ratios, together with O/H ratios consistent with the \\citet{1998ApJ...493..222M} results, contradict the argument by \\citet{2006guillaume} against a high present-epoch D/H ratio, on the basis that no correspondingly high D/O ratios have been measured. It is also hard to explain these high D/H ratios on the basis of localized infall of metal poor gas, as one would expect the O/H and N/H ratios for these sightlines to be affected as well. In addition, several studies of the O/H ratio \\citep{1998ApJ...493..222M,2003ApJ...591.1000A,2004ApJ...613.1037C,2005ApJ...625..232O} have shown that this ratio is constant across a large range of $N$(O\\,I), distance to targets, and average sightline gas density. However, note that in a special case where the infalling gas had a high D/H ratio and metallicity comparable to the local ISM value of \\citet{1998ApJ...493..222M}, such a mechanism could not be ruled out. Variable astration in which the amount of D burned in the interior of stars and O produced, varies from sightline to sightline, seems also to be an unlikely mechanism for explaining the high D/H ratios. In the case of WD\\,1034$+$001 and TD1\\,32709 less deuterium would have to have been destroyed by astration and consequently less O would have been produced by supernovae. This scenario seems implausible because it implies that astration rates would have to be variable on short distance scales in order to explain the nearby sightlines with LB-like D/H ratios and also because it is not supported by our derived O/H and N/H ratios. In addition, in both the infall and variable astration scenarios, these processes would have to occur in a time smaller than the typical mixing time-scale of 350 Myr \\citep{2002ApJ...581.1047D}. Our results support the idea that some sightlines in the Milky Way ISM have high D/H ratios. Linsky et al. (2005, in prep) have used the existence of these sightlines together with the deuterium dust-depletion model of \\citet{2004oee..symp..320D} to argue for a high present-epoch abundance of deuterium. Linsky et al. (2005, in prep) estimate the total abundance of D (in gas $+$ dust) in the local disk of the Galaxy, D/H $\\ge$ (2.19 $\\pm$ 0.27)$\\times10^{-5}$, by taking the weighted average of the ratios for the five sightlines with the highest values. Our new measurements for WD\\,1034$+$001 and TD1\\,32709 are in good agreement with this estimate. Below we compare our new measurements with ratios from the literature.    \\label{summary}  We have used data obtained with {\\it FUSE} together with archival data from {\\it IUE} and {\\it HST} to derive column densities and ratios of column densities along the lines of sight to WD\\,1034$+$001, BD$+$39\\,3226, and TD1\\,32709. The D/H derived here for two of these sightlines are not consistent at the 1 $\\sigma$ level with the Local Bubble value; none of the D/O ratios are consistent, at the 1 $\\sigma$ level with the Local Bubble value, and present some of the highest values in the literature. Considered along with the ratios published for other sightlines, our results reinforce the scatter in the D/H measurements and indicate that there is also scatter in the D/O ratio, implying that the high D/H ratios derived here and along other sightlines are unlikely to be due to problems with the $N$(H\\,I) determinations. We estimate that the additional, presently unknown, systematic errors on the determination of $N$(H\\,I) would have to be of the order of 40\\% of each D/H ratio, in order to bring all the D/H ratios into agreement. Thus, our results support the idea that some sightlines in the Milky Way ISM have high D/H ratios. For sightlines with log$N$(H) $>$ 20.7 the O/H ratios are not consistent with a single value, and are inconsistent with the local ISM value derived by \\citet{1998ApJ...493..222M}. The hint of anti-correlation between D/H and O/H for these sightlines could be an indication of astration.  The trend between decreasing D/H with decreasing Fe/H observed by Linsky et al. (2005, in prep) is also observed when D/O versus Fe/O are considered. These ratios are not subject to possible systematic uncertainties in the determination of $N$(H).  In order to try to understand the behavior of the deuterium abundance we performed a  study of D/H versus the sightline average gas density. Our results indicate that as long as the gas probed by the sightlines is in warm diffuse clouds with $\\langle n_{\\rm H}\\rangle <$0.10 cm$^{-3}$ D/H seems to be constant and have the Local Bubble value. In addition, our study shows that D/H seems to decrease with increasing sightline gas density, similarly to what has been observed for other elements such as Fe and Si, which strongly supports the idea that D might be depleted into dust grains. Finally, a few sightlines do not follow the trend, but show exceptionally high D/H ratios."
    },
    "0601/astro-ph0601067_arXiv.txt": {
        "abstract": "We show that Be stars belong to a high velocity tail of a single B-type star rotational velocity distribution in the MS. This implies that: 1) the number fraction N(Be)/N(Be+B) is independent of the mass; 2) Bn stars having ZAMS rotational velocities higher than a given limit might become Be stars. ",
        "introduction": "\\label{zrv}  The observed $V\\!\\sin i$ of 127 galactic Be stars were corrected for uncertainties due to the gravitational darkening effect with the {\\sc fastrot} models (Fr\\'emat et al. 2005). Using the same models we calculated the today true rotational velocity $V$ of each star by determining its inclination angle $i$. With evolutionary tracks for rotating stars we determined the mass and age of the studied stars (Zorec et al. 2005). From each $V$ we estimated the corresponding $V_{\\rm ZAMS}$ by taking into account four first order effects affecting the evolution of equatorial velocities: 1) variation induced by the time-dependent stellar radius; 2) angular momentum (AM) loss due to mass-loss phenomena in stars with mass $M\\ga15M_{\\odot}$ and conservation of AM for $M\\la15M_{\\odot}$; 3) changes of the inertial momentum due to evolutionary effects and to the rotation using 2D barotropic models of stellar structure (Zorec et al. 1988); 4) internal redistribution of the AM in terms of the meridional circulation time scale parameterized from Meynet \\& Maeder's (2000) models. The obtained $V_{\\rm ZAMS}$ against the stellar mass are shown in Fig.~{\\ref{f1}a, where there is a neat mass-dependent limiting cut $V_{\\rm min}(M)$ that confines all studied Be stars in a high velocity sector. This indicates that stars need to have $V_{\\rm ZAMS}\\ga$ $V_{\\rm min}$ to become Be in the MS phase. For each mass-interval we divided the $V_{\\rm ZAMS}$ by $V_{\\rm min}$ and obtained the global histogram shown in Fig.~{\\ref{f1}b. The fit that better describes the distribution of $V_{\\rm ZAMS}(M)/V_{\\rm min}(M)$ obtained is a Gaussian tail. The histogram concerns only 17\\% roughly of stars out of the whole B star MS population. Since more than 80\\% of MS B-type stars must then be in the $V_{\\rm ZAMS}/V_{\\rm min}$ $\\la 1$ interval, it implies that Be stars do not form a separate distribution, but possibly a tail of a general distribution that encompasses the whole B-star MS population.  \\begin{figure}[t] \\centerline{\\psfig{file=jzorec1_fig1.eps,width=13.5truecm,height=5truecm}} \\caption[]{a) Distribution of the equatorial velocities in the ZAMS of the studied Be stars against the mass . b) Frequency distribution of $V_{\\rm ZAMS}/V_{\\rm min}$ and fit with a Gaussian function. No Be star has $V_{\\rm AMS}/V_{\\rm min}<1$} \\label{f2} \\end{figure}  \\begin{figure}[h!] \\centerline{\\psfig{file=jzorec1_fig2.eps,width=13.5truecm,height=8truecm}} \\caption[]{a) H$\\alpha$ emission intensity function as a function of $T_{\\rm eff}$ and $\\tau_o$. b) Number of stars around the Sun in an apparent V=7 magnitude limited volume. c) Normalized predicted number of Be stars near the Sun. d) Comparison of predicted with observed number of Be stars near the Sun} \\label{f1} \\end{figure}  \\vspace{-0.5cm} ",
        "conclusions": ""
    },
    "0601/astro-ph0601584_arXiv.txt": {
        "abstract": "Diffuse far-UV emission arising from the edge of the Orion-Eridanus superbubble was observed with the \\emph{SPEAR} imaging spectrometer, revealing numerous emission lines arising from both atomic species and H$_{2}$. Spatial variations in line intensities of C{\\sc iv}, Si{\\sc ii}, and O{\\sc vi}, in comparison with soft X-ray, H${\\alpha}$ and dust data, indicate that these ions are associated with processes at the interface between hot gas inside the bubble and the cooler ambient medium. Thus our observations probe physical conditions of an evolved thermal interface in the ISM. ",
        "introduction": "The Orion-Eridanus Superbubble (OES) is a large cavity in the interstellar medium (ISM) created by some combination of stellar winds and supernovae from the enclosed Orion OB1 stellar association, \\citep[e.g.][]{reynolds79}. Its comparable proximity and size ($\\sim$300~pc) give it a large apparent angular extent of $\\sim$30$^{\\circ}$ \\citep{burrows96}, allowing for detailed study of its structure. The complicated morphology and dynamics, described by \\cite{guo95}, indicate that the OES is still evolving and expanding into the ISM and thus it likely includes a variety of interfaces and asymmetries resulting from interacting phases of interstellar gas. The OES has clearly-delineated thermal boundaries between the ambient ISM and the suberbubble cavity. The cavity is bright in soft X-ray (SXR) emission, indicating that it is filled with hot T$\\sim$$10^{6}$K gas \\citep[see][]{heiles99}. For example, the X-ray edge of the OES toward Galactic South (lower right of Fig. \\ref{slitpositions}) coincides with a ridge of H$\\alpha$ emission produced by cooler T$\\sim$$10^4$~K gas, beyond which still-cooler neutral H and dust are seen in 21-cm \\citep{heiles99} and infrared \\citep{burrows93}.  \\begin{figure} \\epsscale{.4} \\plotone{f1.eps} \\caption{Composite 1/4~keV, H$\\alpha$, and dust map image \\citep{skyview} of an OES interface in greyscale [color image online], overlaid with \\emph{SPEAR} slit position. The \\emph{SPEAR} field spans the edge of the OES wall, sampling both hot inside and cooler outside material. The gridlines show ecliptic latitude and longitude.} \\label{slitpositions} \\end{figure}  We took a detailed far ultraviolet (FUV) emission spectrum with the \\emph{SPEAR} instrument to study the physical conditions across one edge of the OES, where hot and cooler gas  are likely to be interacting. {\\em SPEAR} (aka {\\em FIMS}) spectrally images diffuse FUV background radiation with \\ensuremath{\\lambda /\\Delta \\lambda}$\\sim$550, using two large field of view (FOV) imaging spectrographs optimized for measuring diffuse emission: the Short 'S' (900-1150 {\\AA}) and Long 'L' (1350-1750 {\\AA}) channels. The instrument and on-orbit performance are described in \\citet{edelstein05a}.\\\\ ",
        "conclusions": "\\label{conclusions}  We have discovered numerous diffuse atomic and molecular FUV emission lines emanating from both hot and cool gas toward the OES. Of these lines, only O{\\sc vi} has been previously detected in the OES at all, though not at the bubble edge. We measured intensity variations of O{\\sc vi}, C{\\sc iv}, and Si{\\sc ii} emission across the interface between the hot bubble interior and the ambient cooler media. The O{\\sc vi} and C{\\sc iv} emission are greatly enhanced at the interface, and clearly show an ionized stratification. While these ions could be formed in a fast shock, we suggest that a quiescent thermal interface model is also consistent with the FUV observations and with other previous observations finding only low-velocity gas. We have also observed a significant H$_{2}$ component, despite the moderate total hydrogen column in this direction.  The spatial variation of emission line intensity across the interface holds much promise for probing the physics of the interface, especially when applied to species of widely varying ionization potential. To further probe the physics of the OES region, we intend to extend our analysis to additional spectral lines and to two more deep observations \\emph{SPEAR} made of the OES toward nearby fields. We will also test interface models predicting the spectral components. These new data can contribute to the understanding of interstellar thermal heating sources and interfaces by comparing the FUV emission line observations with predictions made with interface models invoking shocks, conduction, and turbulent mixing."
    },
    "0601/astro-ph0601298_arXiv.txt": {
        "abstract": "{If dark matter is made of neutralinos, annihilation of such Majorana particles should produce high energy cosmic rays, especially in galaxy halo high density regions like galaxy centres.} {M31 (Andromeda) is our nearest neighbour spiral galaxy, and both its high mass and its low distance make it a source of interest for the indirect search for dark matter through $\\gamma$-ray detection.} {The ground based atmospheric Cherenkov telescope CELESTE observed M31 from 2001 to 2003, in the mostly unexplored energy range 50-500 GeV.} {These observations provide an upper limit on the flux above 50 GeV around $10^{-10}\\rm{cm}^{-2}\\rm{s}^{-1}$ in the frame of supersymmetric dark matter, and more generally on any gamma emission from M31. ",
        "introduction": "The presence of dark matter in the Universe has been known for decades. Since early measurements in galaxy clusters (Zwicky,~\\cite{zwicky}), the mass distribution of the Universe has been studied at different scales with a focus on dynamical effects by means of galaxy star rotation curves, large-scale galaxy cluster dynamics (Ostriker et al.,~\\cite{ostriker}). Recent developments in observational techniques in cosmology have resulted in independent estimates of the matter content of the Universe $\\Omega_{\\rm{m}}{h^{2}}$, through large-scale structure surveys (Hawkins et al.,~\\cite{2dF} ; Loveday et al.,~\\cite{sdss}) and measurements of the cosmic microwave background radiation (the most recent being the WMAP mission (Spergel et al.,~\\cite{wmap})). Given standard cosmology, all suggest that most of the matter in the Universe is dark, cold and non-baryonic. Such hypotheses have led to the construction of the Cold Dark Matter (CDM) paradigm (Blumenthal et al.,~\\cite{blumenthal}): dark matter would be made of Weakly Interacting Massive Particles (WIMPs) which are neutral, stable and originating from the Big Bang era (Lee \\& Weinberg,~\\cite{lee}).  Supersymmetric theories (SUSY) (see for instance Nilles,~\\cite{nilles}) offer an excellent WIMP candidate (Goldberg,~\\cite{goldberg}), the neutralino, which is a mixture of the superpartners of the neutral Higgs bosons and of the electroweak gauge bosons. We will not discuss the case of extra-dimension phenomenological theories, which also provide interesting candidates (Servant \\& Tait,~\\cite{servant}).   The nature of dark matter is presently probed in both direct searches, by means of underground experiments that could detect elastic interactions of neutralinos with nuclei, and indirect searches using ground based or satellite telescopes to detect cosmic rays (gamma, leptons or hadrons) created by neutralino pair annihilations in galactic or extragalactic media (for review, see Bergstr\\\"om~\\cite{bergstrom2}).  These different types of searches, together with collider experiments, are necessary to constrain in a wider view the quantum nature of dark matter, because they allow either additional or complementary surveys of the particle model parameter spaces.   Searching for WIMP annihilation signatures with ground based $\\gamma$-ray telescopes leads to the question of the choice of targets. A good candidate will have a large amount of dark matter, and combine as big a density as possible, as small a distance from us as possible, and finally, will transit at high elevation in the experimental sky. The Galactic centre is a prime candidate except for being too close to the horizon for CELESTE, which is in the northern hemisphere. Instead, we have chosen M31 and the Draco dwarf galaxy for our searches. M31, which is the nearest giant spiral galaxy, is very massive ($\\sim 10^{12} M_{\\odot} $), and its star rotation curve indicates a large amount of dark matter. Draco, a neighbour dwarf spheroidal galaxy dominated by a dark component (Kleyna et al.,~\\cite{kleyna}), is also a very good candidate but our attempts to study it were foiled by bad weather.   In this paper, we present the result of searches for $\\gamma$-ray emission from M31 with the CELESTE telescope. In section~\\ref{predictions}, we review the predictions made for the observations (more details in Falvard~\\cite{falvard}, hereafter F04), for which we considered CDM in the frame of minimal supergravity (mSUGRA) phenomenology. The impact of the astrophysical modelling is briefly revisited, as well as possible consequences of non-standard cosmologies. Focusing on the experimental techniques, we present in section~\\ref{celeste} the method we use to search for a $\\gamma$-ray signal, with explicit comparisons to Crab data. In the absence of a detection, $2\\sigma$ confidence level limits are computed for all studied SUSY models. ",
        "conclusions": "Observations of M31 with CELESTE provide a $2\\sigma$ upper limit on the $\\gamma$-flux above 50 GeV, depending on the expected spectrum. This limit, around $10^{-10}\\rm{ph.cm}^{-2}\\rm{s}^{-1}$, is quite far from the SUSY parameter space, but significantly constrains combinations of different enhancement factors discussed in \\S~\\ref{xtrafactors} (which are also likely to be excluded by EGRET limits, depending on the neutralino mass), and also any other model of annihilating dark matter besides SUSY. Whereas these observations have been motivated by indirect searches for SUSY CDM, this result yields a general astrophysical result: the first observation of a spiral galaxy in this energy range, somehow constraining $\\gamma$-ray emission from this class of objects.   Any gamma ray detection from a galaxy like M31 would be difficult to interpret in terms of dark matter annihilation. Spiral galaxies are known sites of non-thermal processes and cosmic ray acceleration, and the relevant physical mechanisms are not yet well understood. In this sense, a Dwarf Spheroidal galaxy like Draco is a very promising source for indirect detection, given it is clearly dominated by the dark matter component. Unfortunately, we have too few data on Draco to perform a relevant analysis.   However, the upcoming generation of $\\gamma$-ray instruments will undoubtedly further constrain dark matter models and halo models for various astrophysical sources. These searches are not only complementary to future particle physics experiments, but also very important to understand how the question of dark matter is connected to the particle content of the Universe."
    },
    "0601/astro-ph0601251_arXiv.txt": {
        "abstract": "We report on IRAC-4.5\\micpa, IRAC-8.0\\mic and MIPS-24\\mic deep observations of 16 Gamma-Ray Burst (GRBs) host galaxies performed with the \\spi \\, {\\it Space Telescope,} and we investigate in the thermal infrared the presence of evolved stellar populations and dust-enshrouded star-forming activity associated with these objects. Our sample is derived from GRBs that were identified with sub-arcsec localization between 1997 and 2001, and only a very small fraction ($\\sim$\\,20\\%) of the targeted sources is detected down to $f_{4.5 \\mu m}$\\,$\\sim$\\,3.5\\muJy and $f_{24 \\mu m}$\\,$\\sim$\\,85\\muJypa~(3$\\sigma$). This likely argues against a population dominated by massive and strongly-starbursting (i.e., SFR\\,$\\gtapp$\\,100\\,\\Msol\\,yr$^{-1}$) galaxies as it has been recently suggested from submillimeter/radio and optical studies of similarly-selected GRB hosts.  Furthermore we find evidence that some GRBs do not occur in the most infrared-luminous regions -- hence the most actively star-forming environments -- of their host galaxies.  Should the GRB hosts be representative of all star-forming galaxies at high redshift, models of infrared galaxy evolution indicate that $\\gtapp$\\,50\\% of GRB hosts should have $f_{24 \\mu m}$\\,$\\gtapp$\\,100\\muJypa. Unless the identification of GRBs prior to 2001 was prone to strong selection effects biasing our sample against dusty galaxies, we infer in this context that the GRBs identified with the current techniques can not be directly used as unbiased probes of the global and integrated star formation history of the Universe. ",
        "introduction": "It is now widely believed that the so-called ``long'' Gamma-Ray Bursts (i.e., GRBs with duration $\\gtapp$\\,2s and soft spectra, as opposed to the short and hard bursts, e.g., \\citealt{Kouveliotou93}) are intimately connected to the collapse and the cataclismic destruction of some short-lived and very massive stars.  Evidence in this regard include (i) the signature of Type Ic supernova and hypernova in the optical transient emission of GRB counterparts \\citep[e.g.,][]{Galama98,Stanek03,Hjorth03, Malesani04}; (ii) heavy elements from metal-enriched media typical of supernova remnants observed in the spectrum of several X-ray afterglows \\citep{Piro00,Reeves02}; (iii) the unquestionable starbursting nature of their host galaxies \\citep[e.g.,][]{Bloom98a,Fruchter99a,Sokolov01,Chary02,LeFloch03,Christensen04}; and (iv) the location of GRBs relative to the center of their hosts, which appears to be consistent with a population of progenitors residing in galaxy disks \\citep{Bloom02a}. Hence it has often been proposed that long GRBs could be used as powerful tracers of the global star-forming activity in the early Universe \\citep[e.g.,][]{Wijers98,Mirabel00,Blain00}. As illustrated by the recent identification of a burst at $z$\\,$=$\\,6.29 \\citep[e.g.,][]{Kawai05}, GRBs are indeed detectable up to very high redshifts \\citep{Lamb00}. They are also very little affected by dust extinction, which is known to be particularly significant in distant starburst galaxies \\citep[e.g.,][]{Blain99,Franceschini01,Chary01}.  However, this picture relies on the strong assumption that the production of GRBs in a given starburst region scales only with the rate of stars formed in this environment, with no dependency on other physical parameters that may vary from one galaxy to another or that may evolve throughout the lifetime of the Universe.  This assumption is rather major, as the occurrence of a GRB could strongly depend on the properties of the gas from which its progenitor originates (e.g., metallicity: \\citealt{MacFadyen99,Ramirez_Ruiz02b,Heger03,Hirschi05}). It could also vary according to the fraction of those progenitors involved in binary systems \\citep{Izzard04,Podsiadlowski04,Mirabel04b}, and it finally relies on a non-evolution of the Initial Mass Function with redshift. Testing this ``one-to-one'' connection between GRBs and star formation is therefore particularly crucial to guarantee an accurate understanding of the use of GRBs as quantitative tracers of galaxy evolution.  One possible approach to investigate this relation is to compare the properties of the hosts of Gamma-Ray Bursts with respect to the galaxies responsible for the bulk of the star-forming activity in the Universe as a function of redshift.  It is now well established that a significant fraction of the present-day stellar mass budget was formed during brief and violent infrared-luminous (L$_{\\rm IR}$\\,=\\,L$_{\\rm 8-1000\\mu m}$\\,$\\gtapp$\\,10$^{11}$\\,\\Lsol) episodes of star formation within massive ($\\mathcal{M}$\\,$\\gtapp$\\,5\\,$\\times$\\,10$^{10}$\\,M$_\\odot$) galaxies at 0.5\\,$\\ltapp$\\,$z$\\,$\\ltapp$\\,3 \\citep[e.g.,][]{Flores99,Blain02,Elbaz02,Dickinson03,Lagache03,Franceschini03,LeFloch05,Hammer05}. Investigating the extinction-corrected star formation rate and the mass of GRB hosts should therefore provide tight constraints on the relevance of GRBs for probing the star formation history of the Universe.  Previous observational studies on GRB hosts though have led to conflicting views about their nature.  Based on their properties at optical and near-infrared wavelengths it has been argued that GRB hosts are mostly blue, sub-luminous and low-mass galaxies with young stellar populations, characterized by a modest activity of star formation and potential selection effects due to low metallicity \\citep[e.g.,][]{Sokolov01,LeFloch03,Fynbo03,Courty04,Prochaska04,Christensen04}. On the other hand it has also been claimed that the morphology of these objects and their average radio/submillimeter properties rather indicate massive and actively star-forming galaxies \\citep{Conselice05,Berger03}.  In order to bring tighter constraints on the nature of those sources, we undertook a survey of GRB host galaxies with the \\spi \\, {\\it Space Telescope.} In this paper we present mid-infrared (mid-IR) images of 16 objects at 4.5\\micpa, 8\\mic and 24\\micpa, which allows us to constrain the presence of evolved stellar populations and dust-enshrouded star-forming activity in these sources. In Sect.\\,2 we describe the \\spi \\, data used in this study. Sect.\\,3 outlines general results derived from these mid-IR observations, while the GRB host properties are detailed on a galaxy-by-galaxy basis in Sect.\\,4.  Our findings are discussed in Sect.\\,5, and they are finally summarized in Sect.\\,6.  Throughout this work, we assume a $\\Lambda$CDM cosmology with H$_0$\\,=\\,70~km~s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_m$\\,=\\,0.3 and $\\Omega_{\\lambda}\\,=\\,0.7$ \\citep{Spergel03}. ",
        "conclusions": "\\subsection{The origin of the infrared emission in the GRB host galaxies}  The IRAC and MIPS instruments on-board \\spi \\, have recently opened new exciting perspectives for tackling the evolution of galaxies in the early Universe. With unprecedented sensitivity at 3.6--8.0\\micpa, the IRAC channels can sample the rest-frame near-IR emission of very distant sources, while the MIPS imager can detect the hot dust emission of active starbursts and AGNs up to $z$\\,$\\sim$\\,2--3.  Both cameras provide therefore new insights to study the evolved stellar populations and the importance of dust-obscured star formation at high redshift.  Interpreting the observed mid-IR emission of GRB-selected galaxies can be however more subtle than analyzing the other mid-IR sources detected in the field. As mentioned in Sect.\\,3.2 \\citet{Venemans01} have argued that optical and UV afterglows can be significantly absorbed by dusty material surrounding GRBs, and that the reprocessed IR light can be observed several years after a burst because of the very slow time-scale variations of the heated dust emission.  The SED of this dust component would be characterized by a steep increase from the short wavelengths up to $\\sim$\\,8--10\\mic rest-frame, followed by a gentle decline in the mid-IR and far-IR \\citep[see Fig.\\,5 of ][]{Venemans01}.  Up to redshift $z$\\,$\\sim$\\,1--2, it could dominate the total luminosity of the host galaxy, then questioning the IRAC and MIPS flux measurements as mass and star formation rate indicators.  However we did not find any obvious evidence for such ``GRB-heated'' dust emission among the \\spi-detected GRB hosts. In the resolved GRB\\,980425 host galaxy at $z$\\,=\\,0.0085 we do not see any signature of this effect at the location of the hypernova SN\\,1998bw, while the infrared emission detected in the complex environment of the host of GRB\\,980613 ($z$\\,=\\,1.10) does not coincide with the galaxy fragment where the burst occured.  At the redshift of the other sources (i.e., 0.84\\,$\\leq$\\,$z$\\,$\\leq$1.1), the IRAC-8\\mic and MIPS-24\\mic bands constrain the rest-frame 4\\micpa/12\\mic flux ratios, which appear to be too red compared to the SED of the GRB dust component predicted by the models.  The detectability of this GRB dust emission obviously depends on the efficiency of the absorption of the UV/optical afterglow.  As already noted by \\citet{Venemans01}, it should be preferentially observed toward the host galaxies of those ``dark bursts'' originating from dust-enshrouded star-forming regions\\footnote{The so-called ``dark bursts'' can also be GRBs with intrinsically faint afterglows \\citep[e.g.,][]{Fox03} or very high redshift bursts with optical counterparts suppressed by Ly$\\alpha$ absorption \\citep[e.g.,][]{Lamb00}.}.  Even though three~of our galaxies fall indeed in this category (i.e., the hosts of GRB\\,970828, GRB\\,981226 and GRB\\,990506), our sample has been mostly built from GRBs pinpointed with optical afterglows, and our selection could also induce a bias against the detection of this burst-remnant dust emission.  As a result, we conclude that the contribution of this component is negligible in our data.  The IRAC and MIPS detections should truly reflect the properties of the stellar populations and the global star-forming activity within the observed host galaxies.    \\subsection{A panchromatic view on the nature of the GRB hosts}  A few sources from our sample display clear signatures of evolved stellar populations and intense starbursting activity. An example is the host of GRB\\,990705 at $z$\\,=\\,0.84, which harbors active star formation ($SFR_{\\rm \\, ir}$\\,$\\sim$\\,32\\,\\Msol\\,yr$^{-1}$, see Sect.\\,\\ref{sec:990705}) and appears as a large and massive Sc spiral galaxy typical of the disk-dominated systems at these redshifts. Other cases of active starbursts, like the host of GRB\\,000418, have previously been found at radio and submillimeter wavelengths \\citep{Berger03}.  On average though, our MIPS observations combined with the submillimeter and radio published photometry do not really favor a population of GRB host galaxies dominated by very active and luminous dusty starbursts.  With the exception of the host of GRB\\,980703 that we discussed in Sect.\\,4.5, our 24\\mic non-detections are consistent with the existing SCUBA and VLA data and they are even more constraining if one assumes typical IR-luminous and starburst-dominated SEDs for the host of GRBs (see Figure\\,7). Our measurements argue against the presence of numerous LIRGs and ULIRGs in the GRB host population and they also point to lower star formation rates (see Sect.\\,3.2 and Table\\,1) which are not consistent with some conclusions obtained by other groups.  For instance \\citet{Berger03} have recently claimed that 20\\% of GRB host galaxies have $SFR \\sim 500$\\,\\Msol\\,yr$^{-1}$, which is obviously in disagreement with our \\spi \\, results.  Similarly our IRAC data argue very clearly against a population of sources with massive and evolved stellar populations dominating their near-IR rest-frame emission.  Assuming a standard conversion between the mass and the rest-frame near-IR absolute luminosity \\citep[i.e., --0.25\\,$\\ltapp$\\,log$_{10}$\\,$(\\mathcal{M}/L_K)$\\,$\\ltapp$\\,+0.15, ][]{Bell03}, the fluxes or upper limits that we measure at 4.5\\mic translate for most cases into masses $\\mathcal{M}$\\,$\\ltapp$\\,5$\\times$10$^{9}$\\,\\Msol.  The IRAC data are also consistent with the constraints derived from the optical and the near-IR published photometry assuming typical SEDs of blue star-forming galaxies (see Figure\\,7), and our results globally agree with the rather small masses determined by \\citet{Chary02} using deep $K$-band observations of GRB hosts at Keck.  Our interpretation, on the other hand, strikingly contrasts with the conclusions recently obtained by \\citet{Conselice05} who argue for a trend toward massive sources at $z$\\,$\\gtapp$\\,1 based on an analysis of the GRB host morphologies.  Hence the \\spi \\, view on the GRB-selected galaxies rather suggests low mass sources characterized by a relatively modest amount of dust obscuration.  Similar conclusions have already been derived based on the optical and near-IR properties of these objects \\citep[e.g.,][]{LeFloch03,Courty04,Christensen04}.  \\subsection{Implications to GRBs as SFR probes}  Because of their dust-penetrating power and their relation with the death of young and massive stars it has been argued that the long GRBs could be used as an unbiased probe of the star formation history of the Universe. On a pure statistical point of view, their host galaxies should be therefore representative of the sources producing the bulk of the stellar mass at high redshift.  Are our results consistent with this picture\\,?  The detection of the cosmic infrared background and the recent surveys performed at infrared/submillimeter/radio wavelengths have revealed that a significant fraction of the present-day stellar mass was formed in the past during short-live and dust-obscured episodes of intense star formation within infrared-luminous galaxies \\citep[e.g.,][]{Puget96,Blain99,Chary01,Cowie04}.  At $z$\\,$\\sim$\\,1, these infrared-luminous starbursts (i.e., \\Lir\\,$\\geq$\\,10$^{11}$\\,\\Lsol) are responsible for $\\sim$70\\% of the star-forming activity in the Universe \\citep{LeFloch05} and their contribution is believed to be even more important at higher redshifts \\citep[e.g.,][]{Blain02,Lagache03}. Using \\spi-updated IR galaxy evolution models \\citep[e.g.,][]{Lagache04,Chary04} we estimate that more than half of the global star-forming activity throughout the lifetime of the Universe has occured within galaxies characterized by $f_{24 \\mu m}$\\,$\\gtapp$\\,100\\muJypa. In this context, the significant fraction of non-detections in our MIPS data is therefore surprising.    Similarly, the downsizing evolution of the cosmic star formation history reveals that the bulk of the star-forming activity has been moving from massive galaxies at high redhifts to low-mass objects in the present-day Universe \\citep[e.g.,][]{Cowie96,Juneau05}.  Should the long GRBs trace the whole population of distant starbursting sources, we should therefore detect these bursts preferentially toward massive systems easily accessible to IRAC.  This apparent discrepancy between the nature of the GRB host galaxies as a whole and the sources dominating the high-redshift star-forming activity has already been noted from a comparison between their colors and luminosities at optical and near-IR wavelengths \\citep[e.g.,][]{LeFloch03} or their star formation rate recovered from cosmological simulations \\citep{Courty04}. For instance GRB hosts appear blue and optically sub-luminous compared to the massive dusty starbursts. Very often they also display H$\\delta$ in emission \\citep{Djorgovski98,Soderberg04,Prochaska04,Gorosabel05}, which reveals star formation episodes characterized by younger populations than those typically observed in the distant luminous-infrared galaxies \\citep{Flores99,Hammer01,Hammer05}. As a result we conclude that the GRB afterglows, as they are currently selected, can not be considered as unbiased probes of the integrated activity of star formation in the high redshift Universe.  \\subsection{A bias in the \\spi \\, sample ?}  As we have described in more detail earlier, our sample consists mostly of galaxies pinpointed using optical afterglows identified prior to 2001, when GRB follow-ups were not as prompt and efficient as they are currently.  This could bias the selection toward the most luminous transients, and therefore against dust-enshrouded GRBs occuring within luminous-infrared galaxies. A relevant case illustrating this hypothesis is given by the dark burst GRB\\,970828 and its host galaxy (see Sect.\\,\\ref{sec:970828}).  The accurate coordinates of this GRB were determined thanks to the detection of its radio afterglow, and the absence of optical transient emission was interpreted as an evidence for dust extinction in the host galaxy \\citep{Djorgovski01a}. The detection of this object at 24\\mic clearly supports this interpretation, and it suggests that a sample of GRB hosts purely selected with optical afterglows could be biased toward dust-free sources. In this context, the XRT instrument on-board the recently-launched {\\it Swift\\,} satellite is now routinely localizing X-ray GRB afterglows with an accuracy of $\\sim$\\,5\\arcsec\\, on the sky. In the near-future, the comparison between the properties of GRB host galaxies selected from optical and X-ray afterglows will provide better insights into this possible bias affecting our current sample.   On the other hand, we have not detected the two other host galaxies selected from radio afterglows with no optical transient (i.e., the hosts of GRB\\,981226 and GRB\\,990506), which shows that not all dark GRBs   originate from  dusty galaxies. A similar conclusion was reached by \\citet{Barnard03} based on the non-detection of four dark GRB hosts with the SCUBA camera at 850\\micpa. In fact, many of these dark bursts could simply be associated with intrinsically very faint and/or fast-decaying optical afterglows \\citep[e.g.,][]{Fynbo01,Fox03}.  In addition to this observational bias that could affect our selection, another explanation might be directly related to the physical properties characterizing the local environments where long GRBs take place. Based on theoretical simulations and the connection between GRBs, hypernovae and Type Ic supernovae, it has been argued that long GRBs should more likely occur within binary systems \\citep{Izzard04,Podsiadlowski04,Mirabel04a,Mirabel04b}, whose frequency in star-forming regions may vary with redshift.  Furthermore it has been suggested that long gamma-ray bursts would be more efficiently produced -- and they would also appear more luminous -- if they originate from low metallicity progenitors \\citep[e.g.,][]{MacFadyen99,Ramirez_Ruiz02b,Meynet05,Hirschi05}. Massive stars with metal-poor envelopes keep a high angular momentum in the latest stage of their evolution, and they are also less subject to mass loss. After the final collapse, this would favor the formation of a fast-rotating black hole with accretion of material, that could more easily lead to a bright GRB event. In this case GRBs would be preferentially observed in young and chemically-unevolved galaxies, which would explain this lack of massive and dusty starbursts in our \\spi \\, data.  In fact, direct evidence for low metallicity has already been observed in several GRB host galaxies \\citep{Soderberg04,Prochaska04,Gorosabel05}.  Furthermore, the global properties of these objects such as their blue colors, their relatively low luminosities, their Ly$\\alpha$ emission and their high {\\it specific\\,} star-formation rate \\citep{LeFloch03,Fynbo03,Christensen04} do support a picture where GRBs occur in low-mass, young and hence metal-poor starbursts.  This is also corroborated by the cosmological simulations obtained by \\citet{Courty04} who identify the GRB hosts as the most efficient star-forming objects (i.e., sources with the highest specific star-formation rate) and not as galaxies with obvious high $SFR$.  In fact a good illustration of this global property can be given by the complex environment of the host of GRB\\,980613. This burst did not occur in the region detected by MIPS and harboring therefore the most intense star-forming activity within the galaxy (see also \\citealt{Hjorth02} for a similar conclusion based on optical HST data). Note that the same interpretation can be derived from the characteristics of the GRB\\,980425 host galaxy, since the hypernova SN\\,1998bw did not occur in the most active region as revealed by the very luminous mid-IR point source detected with \\spi."
    },
    "0601/astro-ph0601121_arXiv.txt": {
        "abstract": "Many clues about the galaxy assembly process lurk in the faint outer regions of galaxies.  The low surface brightnesses of these parts pose a significant challenge for studies of diffuse light, and few robust constraints on galaxy formation models have been derived to date from this technique.  Our group has pioneered the use of extremely wide-area star counts to quantitatively address the large-scale structure and stellar content of external galaxies at very faint light levels. We highlight here some results from our imaging and spectroscopic surveys of M31 and M33. ",
        "introduction": "The study of galaxy outskirts has become increasingly important in recent years. From a theoretical perspective, it has been realised that many important clues about the galaxy assembly process should lie buried in these parts. Cosmological simulations of disk galaxy formation have now been carried out by several groups and have led to testable predictions for the large-scale structure and stellar content at large radii -- for example, the abundance and nature of stellar substructure (e.g.  Bullock \\& Johnston 2005, Font et al. 2005), the ubiquity, structure and content of stellar halos and thick disks (e.g. Abadi et al. 2005, Governato et al. 2004, Brook et al. 2005) and the age distribution of stars in the outer regions of thin disks (e.g.  Abadi et al. 2003).  Bullock \\& Johnston (2005) find that Milky Way-like galaxies will have accreted 100-200 luminous satellites during the last 12~Gyr and that the signatures of this process should be readily visible at surface brightnesses of V$\\sim 30$ magnitudes per square arcsec and lower. Although traditional surface photometry at such levels (roughly 9 magnitudes below sky) remains prohibitive, star count analyses of nearby galaxies have the potential to reach these effective depths (e.g. Pritchet \\& van den Bergh 1994). The requirement of a large survey area (to provide a comprehensive view of the galaxy) and moderate-depth imagery can now be achieved in a relatively straightforward manner using wide-field imaging cameras attached to medium-sized telescopes. ",
        "conclusions": ""
    },
    "0601/astro-ph0601317_arXiv.txt": {
        "abstract": "We consider the thermal structure and radii of strongly irradiated gas giant planets over a range in mass and irradiating flux. The cooling rate of the planet is sensitive to the surface boundary condition, which depends on the detailed manner in which starlight is absorbed and energy redistributed by fluid motion. We parametrize these effects by imposing an isothermal boundary condition $T \\equiv T_{\\rm deep}$ below the photosphere, and then constrain $T_{\\rm deep}$ from the observed masses and radii. We compute the  dependence of luminosity and core temperature on mass, $T_{\\rm deep}$ and core entropy, finding that simple scalings apply over most of the relevant parameter space. These scalings yield analytic cooling models which exhibit power-law behavior in the observable age range $0.1-10\\ {\\rm Gyr}$, and are confirmed by time-dependent cooling calculations. We compare our model to the radii of observed transiting planets, and derive constraints on $T_{\\rm deep}$. Only HD 209458 has a sufficiently accurate radius measurement that $T_{\\rm deep}$ is tightly constrained; the lower error bar on the radii for other planets is consistent with no irradiation. More accurate radius and age measurements will allow for a determination of the correlation of $T_{\\rm deep}$ with the equilibrium temperature, informing us about both the greenhouse effect and day-night asymmetries. ",
        "introduction": "\\begin{deluxetable*}{lccccccr} \\tabletypesize{\\small} \\tablewidth{0pt} \\tablecaption{Transiting Extrasolar Planets \\label{tab1} } \\tablehead{ \\colhead{object} & \\colhead{$a$(au)} & \\colhead{$M_{\\rm p}$($M_{\\rm J}$)} & \\colhead{$R_{\\rm p}$($R_{\\rm J}$)} & \\colhead{ $T_{\\rm eq}[K]$ \\tablenotemark{a} } & \\colhead{ $T_{\\rm deep}[K]$\\tablenotemark{b}} & \\colhead{Age (Gyr)} & \\colhead{Reference}  } \\startdata OGLE-TR-132  & 0.031 & $1.19 \\pm 0.13$ & $1.13 \\pm 0.08$ & 2100& $\\leq 2200$   & 0--1.4  & 1 \\\\ OGLE-TR-56   & 0.023 & $1.24 \\pm 0.13$ & $1.25 \\pm 0.08$  & 2100 & $ 1000-3100$ & $3 \\pm 1$ & 2,3,12 \\\\ HD 209458    & 0.046  & $0.69 \\pm 0.05$  & $\\rm 1.31^{+0.05}_{-0.05}$ & 1500 & 2200-2800 & 4--7  & 4,5\\\\ OGLE-TR-10   & 0.042  & $0.63 \\pm 0.14$  & $1.14 \\pm 0.09$   & 1500 & $\\leq 2600$  &  -- & 6,12  \\\\ OGLE-TR-113 & 0.023 & $1.35 \\pm 0.22$ & $\\rm 1.08^{+0.07}_{-0.05}$ & 1300 & $\\leq 2100$ & -- &  7 \\\\ TrES-1      & 0.039 & $0.73 \\pm 0.04$ & $\\rm 1.08^{+0.05}_{-0.05}$    & 1200 & $\\leq 1000$  & $2.5 \\pm 1.5$ & 5,8 \\\\ OGLE-TR-111  & 0.047 & $0.52 \\pm 0.13$ & $\\rm 0.97 \\pm 0.06$    &1000 & $\\leq 1200$ & -- & 9,12 \\\\ HD 149026 \\tablenotemark{c}    & 0.042 & $0.36\\pm 0.04$ & $0.725 \\pm 0.05$ & 1700 &  &  $2.0 \\pm 0.8$  & 10 \\\\ HD 189733 &  0.031 & $1.15\\pm0.04$ & $1.26\\pm 0.03$ & 1200 & $\\leq 3200$ & &  11 \\\\ \\enddata \\tablenotetext{a}{Here $T_{\\rm eq} \\equiv T_{\\rm *}(R_*/2a)^{1/2}$. See the discussion following eq.\\ (\\ref{eq:Teq}).} \\tablenotetext{b}{Allowed range of $T_{\\rm deep}$ given range of mass, radius and age. If no age range given in the literature, we (arbitrarily) give the maximum value of $T_{\\rm deep}$ for an age less than 10Gyr. However, given an accurate age range, the figures in \\S \\ \\ref{sec:applications} can be used to obtain stronger constraints than given here.} \\tablenotetext{c}{HD 149026's small radius clearly indicates a large core size or heavy element abundance. The present paper does not include heavy element cores, so we do not discuss HD 149026 further.} \\tablerefs{(1) \\cite{2004A&A...424L..31M}, (2) \\cite{2004ApJ...609.1071T}, (3) \\cite{2003ApJ...596.1327S}, (4) \\cite{2002ApJ...569..451C}, (5) \\cite{2005ApJ...621.1072L}, (6) \\cite{2005ApJ...624..372K}, (7) \\cite{2004A&A...421L..13B}, (8) \\cite{2004ApJ...616L.167S}, (9) \\cite{2004A&A...426L..15P}, (10) \\cite{2005ApJ...633..465S},(11)\\cite{2005astro.ph.10119B}, (12) \\cite{2006astro.ph..1024S} } \\end{deluxetable*}  Following the discovery of the planet orbiting 51 Peg \\citep{1995Natur.378..355M,1995AAS...187.7004M},  more than 160 planets have been found around nearby stars using precision Doppler spectroscopy. \\footnote{For up to date catalogs, see http://exoplanets.org/ and http://obswww.unige.ch/~udry/planet/planet.html.} Theories of planet formation now have the demanding task of explaining the existence of gas giants with semi-major axes one hundred times smaller than Jupiter, others with order unity orbital eccentricities, a detailed spectrum of (minimum) planet masses, and metallicity correlations with the parent star.  The discovery of transiting planets in the last five years (see Table \\ref{tab1}) challenges not only theories for the origin of short-period gas giants, but also their structure and thermal evolution, spectrum, and interior fluid dynamics. Measurements of planetary mass, radius, and (stellar) age test cooling models which predict radius as a function of mass and age. The atmospheres of two planets have been directly observed. For HD 209458b, absorption lines (due to stellar photons  passing through planet's atmosphere) have been found \\citep{2002ApJ...568..377C, 2003Natur.422..143V, 2004ApJ...604L..69V}, and the first detections of photons emitted by planets outside our solar system have been made for the thermal emission from HD 209458b \\citep{2005Natur.434..740D} and TrES-1 \\citep{2005ApJ...626..523C}. These observations directly constrain the atmospheric structure, temperature profile and chemical composition near the photosphere.  Evolution of the short orbital period transiting exoplanets is significantly different than for Jupiter and Saturn due to proximity of the parent star \\citep{1996ApJ...459L..35G}. Irradiation increases the photospheric temperature by  nearly an order of magnitude relative to an isolated planet. Irradiation also decreases the cooling rate, and hence the rate of shrinkage, by  altering the surface boundary condition \\citep{2000ApJ...534L..97B}. This is immediately apparent in Table \\ref{tab1} as many transiting extra-solar giant planets (EGP's) have radii significantly larger than Jupiter. As short period planets are expected to be tidally synchronized \\citep{1996ApJ...459L..35G,1997ApJ...481..926M}, the strong day-night temperature contrast will drive winds to transport heat from the day to the night side \\citep{2002A&A...385..166S}. Hence the atmospheric temperature profile depends on a combination of detailed radiative transfer calculations for absorption of starlight, and hydrodynamics to model day-night winds and dissipation of wind kinetic energy. Lastly, tides raised on the planet by the parent star may significantly affect its thermal evolution \\citep{2001ApJ...548..466B,2002A&A...385..166S}. The  free energy available by synchronizing the  planet's spin or circularizing the orbit are comparable or larger than the thermal energy. Hence {\\it if} the heat can be deposited sufficiently deep in the planet in less than a cooling timescale, the cooling can be slowed, or even reversed. However, it is uncertain if tides can deposit heat deep in the planet \\citep{1997ApJ...484..866L, 2004ApJ...610..477O, 2004astro.ph..7628W} .  Evolutionary models show that the cooling rate is quite sensitive to the uncertain surface boundary condition \\citep{2002A&A...385..156G}. This boundary condition has been implemented using various approximations. Full radiative transfer calculations \\citep{2001ApJ...556..885B, 2003ApJ...594.1011H} of {\\it static} atmospheres include the stellar irradiation self-consistently, and determine the temperature structure for a given cooling flux from the deep interior. These calculations compute (rather than assume) the albedo, and determine the temperature rise due to absorption of starlight (the greenhouse effect). Such detailed radiative transfer solutions have been incorporated as boundary conditions for some evolutionary calculations \\citep{2003A&A...402..701B, 2003ApJ...594..545B}. However, as day-night and equator-pole winds are not included, assumptions must be made about how the stellar flux is deposited over the surface of the planet (only day-side versus evenly over the entire surface, etc.) which directly affect the temperature profile. Other evolutionary calculations (e.g. Bodenheimer et al. 2003) solve the radiation diffusion equation and set the temperature at (infrared) optical depth $2/3$ to be the equilibrium temperature, ignoring additional temperature increase due to absorption of starlight. Lastly, a number of groups \\citep{2002A&A...385..166S, 2003ApJ...587L.117C, 2005ApJ...618..512B, 2005ApJ...629L..45C, 2005A&A...436..719I} are beginning to model the day-night winds on tidally locked, short orbital period planets, and the role of clouds \\cite{2003ApJ...589..615F} . As we stress here, the crucial parameter for the cooling rate is the temperature at the radiative-convective boundary, which is orders of magnitude deeper in pressure than the photosphere.  The plan of the paper is as follows. The uncertain surface boundary condition is discussed in \\S \\ \\ref{sec:surfacebc}, motivating the surface isotherm used in our models. Details of cooling models and microphysical input are described in \\S \\ \\ref{sec:numerics}. In \\S \\ \\ref{sec:luminosity} we compute the dependence of the luminosity on planet mass, core entropy and irradiation.  An analytic solution for the temperature profile in the radiative zone, and the position of the radiative-convective boundary are derived in \\S \\ \\ref{sec:transition}. These results are collected together in \\S \\ \\ref{sec:cooling} to derive an analytic cooling model which exhibits simple power-law dependence on time. The radii of irradiated gas giant planets are discussed in \\S \\ \\ref{sec:radiusev}, and the analytic formula for the radius given in eq.\\ (\\ref{eq:Rtfit}).  We apply our models to the observed transiting planets and give constraints on the temperature of the deep surface isotherm in \\S \\ \\ref{sec:applications}. Our main conclusions are summarized in \\S \\ \\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We have presented calculations of cooling and radius evolution for strongly irradiated planets. Novel aspects of this model are the following:  \\begin{itemize}  \\item We argue that the generic outcome of strong surface heating, whether it be due to absorption of stellar flux or dissipation of winds and tidal flow, is that a deep isothermal region exists above the radiative-convective boundary. The thermal time in this layer is sufficiently long that the temperature profile is approximately spherically symmetric, irrespective of the size of the asymmetry near the photosphere. We assign this region the temperature $T_{\\rm deep}$ and treat it as a boundary condition for the cooling models.  \\item We show that the cooling flux is determined at the radiative-convective boundary, which is much deeper than the photosphere. Scalings of the flux with core entropy, $T_{\\rm deep}$, and mass are computed.  \\item These scalings allow us to derive an analytic model for the cooling, which shows power-law decrease over a large range of parameter space. The part of the radius which changes in time (the deviation from the isothermal planet) is also a power-law. An analytic formula for radius evolution is given in eq.\\ (\\ref{eq:Rtfit}), with coefficients in Table \\ref{tab:fits}.  \\item While we have used the SCVH EOS and Allard et.al.(2001) opacities for numerical estimates, the analytic solution makes it particularly clear  which quantities need to be evaluated for a given opacity table and EOS. The luminosity is sensitive only to the local conditions at the radiative-convective boundary, while the core temperature involves building static (i.e. not time-dependent) models.   \\end{itemize}  We have compared our theory to observed masses and radii for the transiting planets in Table \\ref{tab1} (except for HD 149026, which clearly has a large abundance of heavy elements). Our findings are as follows: \\begin{itemize}   \\item Figure \\ref{fig:R_vs_M_data} shows mass versus radius for eight transiting planets, compared to our model. The radii can be broadly explained with solar composition, ages in the range $1-10\\ {\\rm Gyr}$, and temperatures deep in the atmosphere  $T_{\\rm deep} \\leq 3000\\ {\\rm K}$. The largest radii requiring the most irradiation to explain are HD 209458, HD 189733 and OGLE-TR-56.  \\item Figures \\ref{fig:HD209458TRES1}, \\ref{fig:OGLETR56OGLETR10}, \\ref{fig:OGLETR111OGLETR113} and \\ref{fig:OGLETR132HD189733} show constraints on $T_{\\rm deep}$ using measured masses, radii, and ages (when available), and their uncertainties. We find that only HD 209458b is well constrained, with $T_{\\rm deep}=2200-2800\\ {\\rm K}$. OGLE-TR-56 has a weak lower limit, and the other six planets have only upper limits, due to the large measurement uncertainty in the radius, or lack of an age determination. These constraints are summarized in Table \\ref{tab1}.  \\item The equilibrium temperature $T_{\\rm eq}$ is measurable, but plays no part in our model. The deep isothermal temperature $T_{\\rm deep}$ is not measurable, but is crucial for the cooling rate. Radiative transfer calculations find these two temperatures should be strongly correlated. Figure \\ref{fig:TeqvsTdeep} shows $T_{\\rm eq}$ versus $T_{\\rm deep}$. As only upper limits on $T_{\\rm deep}$ are available for all but HD 209458b and OGLE-TR-56, it is difficult to draw conclusions at the present time.  It is in principle possible  for a correlation to exist given the current error bars, however, it is not required.  \\end{itemize}  We hope that our models have illuminated the need for more accurate ages and radii. Once those are in hand, our calculations will provide a measurement of $T_{\\rm deep}$ of adequate accuracy to compare to $T_{\\rm eq}$, thus constraining greenhouse physics and day-night transport.     \\vspace{4cm}"
    },
    "0601/astro-ph0601471_arXiv.txt": {
        "abstract": "In this paper, we consider a class of five-dimensional Ricci-flat vacuum solutions, which contain two arbitrary functions $\\mu(t)$ and $\\nu(t)$. It is shown that $\\mu(t)$ can be rewritten as a new arbitrary function $f(z)$ in terms of redshift $z$ and the $f(z)$ can be determined by choosing particular deceleration parameters $q(z)$ which gives early deceleration and late time acceleration. In this way, the $5D$ cosmological model can be reconstructed and the evolution of the universe can be determined. ",
        "introduction": "Recent observations of high redshift Type Ia supernovae reveal that our universe is undergoing an accelerated expansion rather than decelerated expansion \\cite{RS,TKB,Riess}. In addition, the discovery of Cosmic Microwave Background (CMB) anisotropy on degree scales together with the galaxy redshift surveys indicate $\\Omega _{total}\\simeq 1$ \\cite{BHS} and $\\Omega _{m}\\simeq \\left. 1\\right/ 3$. All these results strongly suggest that the universe is permeated smoothly by 'dark energy', which violates the strong energy condition with negative pressure. The dark energy and accelerating universe has been discussed extensively from different points of view \\cite{Quintessence,Phantom,K-essence}. In principle, a natural candidate to dark energy could be a small cosmological constant. However, there exist serious theoretical problems: fine tuning problem and coincidence problem. To overcome the fine tuning problem, some self-interact scaler fields $\\phi$ with an equation of state (EOS) $w_{\\phi}=p_{\\phi}\\left/ \\rho_{\\phi}\\right.$ were introduced dubbed quintessence, where $w_{\\phi}$ is time varying and negative. In nature, the potentials of the scalar field would be determined from the underlying physical theory, such as Supergravity, Superstring/M-theory etc.. However, disregarding these underlying physical theories just from the phenomenal level, we can design many kinds of potentials to solve the concrete problems \\cite{sahni}. Once the potentials are given, EOS $w_{\\phi}$ of dark energy can be found. On the contrary, the potential can be reconstructed by a given EOS $w_{\\phi}$ \\cite{GOZ}. Then, the forms of the scalar potential are obtained by the observations data with given EOS $w_{\\phi}$. However, it is mentioned that the same models have early deceleration epoch. And, the acceleration of the universe just begins at not distance past. So, starting from the deceleration parameter, we can also reconstruct the cosmological models \\cite{Banerjee}. This is the main motivation of this paper.  The idea that our world may have more than four dimensions is due to Kaluza \\cite{Kaluza}, who unified Einstein's theory of General Relativity with Maxwell's theory of Electromagnetism in a $5D$ manifold. In 1926, Klein reconsidered the idea and treated the extra dimension as a compact small circle topologically \\cite{Klein}. However, up to now, there is no experimental evidence and indication for the existence of extra dimensions. Nevertheless, there is strong theoretical motivation for considering our world spacetime with more than three spatial dimension. $M$-theory is a higher dimensional theory with seven extra spatial dimensions, which try to incorporate quantum gravity in a consistent way. Recently, in higher dimension frame, brane world model has been studied extensively. In brane world model, our four dimensional world is a hypersurface (brane) embedded in a higher dimensional manifold (bulk) and all forces are confined on the brane but gravity can propagate in the bulk. Also, it has interesting consequence in cosmology, for recent review seen \\cite{brane-review}. Also, in higher dimensional frame, the Space-Time-Matter (STM) theory as a modern Kaluza-Klein theory is designed to incorporate the geometry and matter by Wesson and his collaborators \\cite{Wesson}. In STM theory, our world is hypersurface embedded in five-dimensional Ricci flat ($R_{AB}=0$) manifold and all the matter in our world are induced from the higher dimension. The theory requires the metric components containing the extra dimension, saying $y$. And, the compactification or not of the extra dimension is not necessary in STM theory, for Ricci flat ($R_{AB}=0$) five-dimensional manifold. The consequent cosmology in STM theory is studied in \\cite{LiuW}, \\cite{STM-cosmology}, \\cite{WLX}. ",
        "conclusions": "A general class of $5D$ cosmological models is characterized by a big bounce as opposed to the big bang in $4D$ standard cosmological model. This exact solution contains two arbitrary functions $\\mu(t)$ and $\\nu(t)$. By careful observation, one finds that the dimensionless density, deceleration parameters and EOS of dark energy are independent on $\\nu(t)$ explicitly which is included in $A$. In terms of redshif $z$, $\\mu(t)$ can be rewritten into a new arbitrary function $f(z)$. Once the forms of the arbitrary function $f(z)$ are specified, the universe evolution will be determined. In this paper, we reconstruct the arbitrary function $f(z)$ by assuming a particular form of deceleration parameter $q(z)$. In this way, the $5D$ cosmological models are reconstructed and the evolution of the universe can be determined."
    },
    "0601/astro-ph0601192_arXiv.txt": {
        "abstract": "We study the X-ray emission of baryon fluid in the universe using the WIGEON cosmological hydrodynamic simulations. It has been revealed that cosmic baryon fluid in the nonlinear regime behaves like Burgers turbulence, i.e. the fluid field consists of shocks. Like turbulence in incompressible fluid, the Burgers turbulence plays an important role in converting the kinetic energy of the fluid to thermal energy and heats the gas. We show that the simulation sample of the $\\Lambda$CDM model without adding extra heating sources can fit well the observed distributions of X-ray luminosity versus temperature ($L_{\\rm x}$ vs. $T$) of galaxy groups and is also consistent with the distributions of X-ray luminosity versus velocity dispersion ($L_{\\rm x}$ vs. $\\sigma$). Because the baryonic gas is multiphase, the $L_{\\rm x}-T$ and $L_{\\rm x}-\\sigma$ distributions are significantly scattered. If we describe the relationships by power laws $L_{\\rm x}\\propto T^{\\alpha_{LT}}$ and $L_{\\rm x}\\propto \\sigma^{\\alpha_{LV}}$, we find $\\alpha_{LT}>2.5$ and $\\alpha_{LV}>2.1$. The X-ray background in the soft $0.5-2$ keV band emitted by the baryonic gas in the temperature range $10^5<T<10^7$ K has also been calculated. We show that of the total background, (1) no more than 2\\% comes from the region with temperature less than $10^{6.5}$ K, and (2) no more than 7\\% is from the region of dark matter with mass density $\\rho_{\\rm dm}<50 \\bar{\\rho}_{\\rm dm}$. The region of $\\rho_{\\rm dm}>50\\bar{\\rho}_{\\rm dm}$ is generally clustered and discretely distributed. Therefore, almost all of the soft X-ray background comes from clustered sources, and the contribution from truly diffuse gas is probably negligible. This point agrees with current X-ray observations. ",
        "introduction": "In the linear regime of gravitational clustering of cosmic matter, the evolution of baryonic gas follows dark matter point by point (Bi et al. 1992; Fang et al. 1993; Nusser \\& Haehnelt 1999; Nusser 2000). That is, the density and velocity distributions of baryonic matter can be obtained by a similar mapping from the dark matter field. The similar mapping has also been used in modeling the gas in clusters of galaxies (Kaiser 1986). However, it has been found that the similarity between the mass and velocity fields of baryonic gas and dark matter is broken in clusters that are formed through highly non-linear evolution. The mechanism of the breaking of similarity is due to the different dynamical behavior of dark matter and baryonic gas. The former is collisionless, while the latter is approximately an ideal fluid. The velocity field of collisionless dark matter particles is multivalued at the intersection of their trajectories, while baryon fluid always has a single-value velocity field. Thus, shocks in baryon fluid will appear at the intersection and break the similarity between the mass and velocity fields of baryon fluid and dark matter (Shandarin \\& Zel'dovich 1989).  At later times, it has been recognized that shocks occur not only in high-, but also in middle- and even low-density regions. This point has been shown by the dynamical equation of baryonic gas. Although cosmic baryonic gas is a Navier-Stokes fluid, its evolution is dominated by the growth mode, which is approximately governed by a random-force-driven Burgers equation, and the random force is produced by the gravity of the random field of dark matter (Gurbatov et al 1989; Berera \\& Fang 1994; Vergassola et al. 1994; Jones 1999; Matarrese and Mohayaee 2002; Pando et al. 2002, Pando et al 2004). When the Reynolds number is large, shock-caused turbulence, called Burgers turbulence, will develop in the fluid (L\\\"assig 2000). The dynamics and thermodynamics of baryonic gas will be substantially affected by the Burgers turbulence. Dynamically, it will lead to the discrepancy between baryonic matter and dark matter, like the discrepancy of a passive substance from the underlying field during nonlinear evolution (e.g. Shraiman \\& Siggia 2000). Thermodynamically, the Burgers turbulence leads to the multiple phases of thermal properties and to converting the bulk kinetic energy of the fluid to thermal energy. For the system of cosmic baryonic and dark matter, a large Reynolds number actually corresponds to the onset of the non-linear regime of the gravitational clustering. Therefore, the similarity between baryonic and dark matter will inevitably break in the non-linear regime, and the dynamical and thermodynamical effects of the Burgers turbulence must be considered.  Some effects of the Burgers turbulence have been detected with observational data and/or hydrodynamic simulation samples of cosmic baryonic gas. First, the statistical decoupling between the density and velocity fields of baryonic gas and dark matter is found to be significant on scales larger than the Jeans length (Pando et al 2004; Kim et al 2005). The importance of the discrepancy in the early nonlinear evolution has also been noted by Yoshida et al. (2003). Second, the transmitted flux of QSOs' Ly$\\alpha$ absorption is found to be remarkably intermittent (Jamkhedkar et al 2000, 2003, 2005; Pando et al 2002; Feng et al 2003), which is consistent with the prediction of the intermittence of fully developed Burgers turbulence (Polyakov 1995; Balkovsky et al. 1997; Frisch et al 2001). Third, the temperature field of the baryonic gas is multiphase, and the heating caused by the shocks of the Burgers turbulence is substantial. Consequently, the distribution of the baryon fraction on large scales is nonuniform (He et al. 2005), and high-entropy gas is produced (He et al. 2004).  In this paper, we extend these studies to the X-ray emission of the hot baryonic gas in the universe. X-ray emission of galaxy groups is considered as important evidence of the similarity breaking between baryon fluid and dark matter. The relation between X-ray luminosity $L_{\\rm x}$ and temperature $T$ predicted by the similarity takes the form of a power law, $L_{\\rm x}\\propto T^{\\alpha}$, with an index of $\\alpha=2$ (Kaiser 1991), while the observed result is $\\alpha > 2$ (Edge \\& Steward 1991; David et al. 1993; Wu et al 1999; Helsdon \\& Ponman 2000; Xue \\& Wu 2000; Croston et al. 2005). Moreover, the X-ray background in the soft band given by similarity is found to be much higher than the observed upper limit. In order to solve these problems, various models of non-gravitational heating for baryonic gas have been proposed (e.g. Valageas \\& Silk 1999; Tozzi \\& Norman 2001; Voit et al 2002; Zhang \\& Pen 2003; Xue \\& Wu 2003). The common goal of these models is to violate the similarity by heating baryonic gas before it falls into gravity potential wells. The amount of heating is of the order of 1 keV nucleon$^{-1}$ (Pen 1999; Wu et al 1999).  However, most non-gravitational heating models {\\it assume} that the similarity will still hold without the non-gravitational heating sources. This point, as pointed out above, is actually inconsistent with the dynamics of the cosmic baryonic gas in the non-linear regime. Baryonic gas, either in high- or in low-density regions, will be heated when Burgers turbulence develops, regardless of whether non-gravitational heatingi is added. Typical shocks of the Burgers turbulence can convert the kinetic energy of a baryon fluid with a speed of a few hundred km s$^{-1}$ into thermal energy (He et al. 2004; Kim et al. 2005), which is of the order of 1 keV nucleon$^{-1}$. Therefore, the thermodynamical effect of the Burgers turbulence should be substantial on the $L_{\\rm x}-T$ relation and soft X-ray background. We need, at least, to estimate the effect of the hydrodynamic heating on the X-ray emission before considering extra non-gravitational heating sources.  The paper is organized as follows. \\S 2 describes the features of the baryonic gas as a Burgers fluid with samples of hydrodynamic cosmological simulation. \\S 3 presents the analysis of the relationships between X-ray luminosity and mass density, temperature, and velocity dispersion. The soft X-ray background radiation is addressed in \\S 4. Finally, the conclusions and discussion are given in \\S 5. ",
        "conclusions": "The evolution of baryonic gas can be roughly divided into three stages. The first is the linear stage. It can be simply described by the similarity between the baryonic and dark matter. When the non-linear evolution takes place, Burgers turbulence develops. Baryonic gas will decouple with dark matter, and becomes multiphase. The kinetic energy of baryon fluid will be dissipated by the shocks of turbulence. Some cosmic baryonic gas finally falls into massive halos and evolves in the thermal equilibrium state in gravity wells. Some heated baryonic gas remains outside of gravity wells. The Burgers turbulence heating is not uniform. A uniform heating would contradict with the Ly$\\alpha$ forest, which is produced by hydrogen clouds with a temperature of about $10^4$-$10^5$K. Therefore, baryonic gas undergoing the Burgers turbulence evolution must be multiphase. In the stage of the Burgers turbulence, the dominant component in volume fraction is the gas in the low-temperature phase (He et al. 2004).  We show that the X-ray emission of baryonic gas can be well modeled by the above-mentioned scenario. That is, the observed distributions of $L_{\\rm x}$ versus $T$ and $L_{\\rm x}$ versus $\\sigma$ of galaxy groups can be reproduced by the simulation sample of the $\\Lambda$CDM model without adding extra heating sources. We also find that almost all of the soft-band X-ray background radiation is from clustered regions, and the contribution of diffuse gas with $\\rho_{\\rm dm}<50$ is negligible. This is also consistent with current X-ray observations. These results show that the Burgers turbulence heating alone seems to be enough to account for the basic features of X-ray emission of baryonic gas in the universe.  We should also point out that star formation and its feedback on the baryonic gas are not considered in our simulation. Roughly, there are two types of feedback: (1) photoionization heating by the UV emission of stars and AGNs and (2) injection of hot gas and energy by supernova explosions, or other sources of cosmic rays. Actually, the photoionization heating can be properly considered, if the UV background is adjusted by fitting the simulation with the observed mean flux decrement of QSOs' Ly$\\alpha$ absorption spectrum (Feng et al. 2003). The significant effect of injecting hot gas and energy by supernovae is mostly on dwarf galaxies. Recent Ly$\\alpha$ observations of protoclusters (Adelberger et al. 2003) suggest that AGN heating does not drastically affect the gas in clusters. Therefore, this heating mechanism may not be strong enough to change the basic features of the multiphase scenario given in this paper. Moreover, the shocks may lead to the electron temperature $T_{\\rm e}$ being lower than the ion temperature $T_{\\rm i}$. Since the mean of the ratio $T_{\\rm i}/T_{\\rm e}$ is generally less than 3 (Fox \\& Loeb 1997; Takizawa 1998; Yoshida et al. 2005), this effect may lead to an uncertainty of the luminosity $L_{\\rm x}$ that is less than a factor of 2. Therefore, all conclusions should basically be held."
    },
    "0601/astro-ph0601537_arXiv.txt": {
        "abstract": "\\begin{changemargin}{0.5cm}{0.5cm} \\textit{Test beam data collected in October 2004 at CERN PS to validate the AMS 02 Ecal Intermediate Board (EIB) are analyzed. After describing the experimental setup and the event samples, results concerning noise measurement, trigger efficiency and threshold accuracy are presented. They demonstrate that the EIB fulfils the physics requirements. Therefore the analog part of the trigger is validated, and hardware choices are also made towards the final device.} \\end{changemargin} ",
        "introduction": "\\subsection{Introduction} This analysis concerns data taken between September $16^{th}$ and October $7^{th}$ 2004 at CERN. The test beam involved both Tracker and Ecal subgroups, and partly aimed to combine energy measurement with the Ecal and tracking with some silicon ladders in a magnetic field.\\\\ Once installed on the ISS, AMS02 main trigger signal will be provided by the Time Of Flight detectors, which sign the passage of charged particles. In order to perform $\\gamma$ ray astronomy with AMS02, methods based on the conversion of photons in the upper detector, requiring the $e^{\\pm}$ pair rigidity measurement in the Tracker have been made up. In addition to this, the Ecal working group developed a specific trigger only involving the calorimeter, which we call standalone $\\gamma$ trigger (\\cite{note0}, \\cite{note1}, \\cite{note2}). Its signal has to be fast in order to be able to shortcut the main trigger. This method gives a complementary $\\gamma$ detection method, providing a very good energy resolution and sensitivity in the GeV-TeV \\cite{loic}. This test beam was the opportunity to validate some parts of the trigger chain, the Ecal Intermediate Board (EIB) prototype and to test the performance of the standalone $\\gamma$ trigger: stability, effective thresholds, noise level and efficiency.  \\subsection{Ecal instrumentation} AMS02 Ecal consists of alternated planes of lead and scintillating fibers. When a particle crosses the detector, photomultiplier tubes (PMT) collect light from fibers, each PMT having 4 distinct pixels. \\begin{figure}[H] \\centering \\includegraphics[width=10cm]{lightcoll2.ps} \\includegraphics[width=4cm]{photoPM2.ps} \\caption{Shower light collection in the Ecal} \\label{lightcol} \\end{figure} For each PMT, 9 wires provide digitized signals from front end electronics (EFE) : 2 from each pixel (low gain / high gain) and 1 dynode. In addition, the analog dynode signal which is not shaped and digitized (then faster) is amplified, compared to a threshold and used for the trigger. \\begin{figure}[H] \\centering \\includegraphics[width=12cm]{Ecal_instr.eps} \\caption{Ecal PMT readout} \\label{readout} \\end{figure} By recording data from both the EFE and the comparator response, we aim to measure the performance of this trigger system. We expect the trigger to be efficient for a 100 MeV local energy deposition, it has been previously measured \\cite{loic} that it corresponds to a 2.5 mV dynode signal. With the 3 V supply used, an amplification by a factor of 10 has been chosen. 2 types of amplifiers  have been tested and the final choice is to be one of this test beam results. Next section give a complete description of the methods and measurements. ",
        "conclusions": ""
    },
    "0601/astro-ph0601701_arXiv.txt": {
        "abstract": "We investigate the finite source size effect in the context of the wave optics in the gravitational lensing. The magnification of an extended source is presented in an analytic manner for the singular isothermal sphere lens model as well as the point mass lens model with the use of the thin lens approximation. The condition that the finite source size effect becomes substantial is demonstrated. As an application, we discuss possible observational consequences of the finite source size effect on astrophysical systems. ",
        "introduction": "Gravitational lensing is a characteristic phenomenon of the general relativity and has become a very important tool in the fields of cosmology and astrophysics \\cite{SEF,FT}. For example, the existence of the massive compact halo objects (MACHO) in the Galaxy was revealed by the detection of the lensed amplification of stellar objects \\cite{Alcock}, and the recent measurements of the cosmic shear field provide a constraint on the matter distribution independently of the clustering bias \\cite{Bacon,Wittman,Kaiser,Maoli}. Promisingly the gravitational lensing will play a more important role with progress in the capability of observational facilities in future. For example, it will be a useful probe of the nature of the dark energy \\cite{YF,YKMF,LSST,Bartelmann}.  As a fundamental aspect of the gravitational lensing, the effect of wave optics has been investigated by many authors \\cite{Peters:1974gj,DeguchiA,DeguchiB,SEF}. Very recently, this subject is revisited by several authors, motivated by a possible phenomenon which might be observed in the future gravitational wave experiments \\cite{DN,TTN,BHN,PINQ,Ruffa,TN,Takahashi,TNB,TSM,JPM,KYA,KYB,KYC}. In the context of the wave optics of gravitational lensing, the argument on the distance-redshift relation is also revisited \\cite{Yoo}.  The present paper focuses on the finite source size effect in the wave optics of the gravitational lensing. In the first half part of this paper, we present an analytic solution for the wave equation with the singular isothermal sphere lens model. In general, it is difficult to obtain an exact solution for general lens model, excepting a few simple lens models \\cite{Suyama,KYB}. Therefore it is useful to obtain such analytic solution for the wave equation. The wave optics of the singular isothermal sphere lens has been investigated by Takahashi and Nakamura using a numerical technique \\cite{Thesis,TNSIS}. We present the analytic expression for the amplification factor, which is the first aim of the present paper. Then, as an application of the analytic formula, we investigate the finite source size effect in the wave optics, which is the other aim of the present paper. Using the analytic formula, we consider  the energy spectrum from an extended source with a Gaussian distribution of surface brightness \\cite{Krzysztof}. We investigate the condition that the finite source size effect becomes important in the wave optics, including the case near the caustic in the limit of the geometrical optics.  This paper is organized as follows: In section 2, we review the basic formulas for the wave optics in gravitational lensing. The limit of the geometrical optics is also reviewed for self-containment. Then, we present analytic expressions of the amplification factor for the point mass lens model as well as the singular isothermal sphere lens model in section 3. In section 4, we investigate the finite source size effect with the use of the analytic formulas. In section 5, the validity of the approximation of the point source is discussed for the gravitational wave from a compact binary. The femtolensing of the gamma ray burst source is also revisited \\cite{Krzysztof,Gould}, and the finite source size effect is considered. The last section is devoted to summary and conclusions. Throughout this paper, we use the unit in which the light velocity equals $1$. ",
        "conclusions": "In this paper we investigated the finite source size effect on the wave optics in the gravitational lensing. First we presented the analytic expression of the magnification for the SIS lens model as well as the point mass lens model. Based on the result, we evaluated the magnification of the finite-size source, assuming a Gaussian profile for the surface intensity. The analytic expression of the magnification is given in terms of the expansion with respect to the source size. This expression is useful to understand how the finite source size effect works on the spectral feature of the magnification in the wave optics. The condition that the finite source size effect becomes significant is discussed. As application of the result, we considered the finite source size effect on the wave optics in lensing of the gravitational wave from a compact binary and the femtolensing. For the lensing of the gravitational wave, it is demonstrated that the finite source size effect can be negligible as long as we consider the gravitational wave of the frequency around $1$Hz. For the femtolensing of the gamma ray burst, we confirmed the result by Stanek et al. \\cite{Krzysztof}, for the point mass lens model. We also considered the femtolensing by the hypothetically very small halo. The femtolensing might imprint the lensing signature on energy spectra if occurred, however, the finite source size effect is crucial. If the source size is larger than $10^5$km, the finite source size effect becomes significant and will not allow the detection of the interference signature in the energy spectra.  It is worthy noting the finite source size effect near the caustic. In the wave optics of the lens configuration of the Einstein ring, the maximum magnification of the point source is not divergent, but is in proportion to the frequency of the wave. But, for an extended sources, the finite source size effect becomes substantial for $a_Sw \\simgt1$, in which the maximum magnification is suppressed by the value $\\sqrt{\\pi/2}/a_S$ as long as $a_S\\ll 1/\\sqrt{w}\\ll1$. This finite source size effect would be influential in the femtolensing near the caustic too, if the source size is larger than $10^5$km.   \\ack All the numerical computation presented in this paper were performed with the help of the package MATHEMATICA version 5.0. The authors thank an anonymous referee for useful comments which helped improve the manuscript. We are also grateful to Y. Kojima, R. Yamazaki, M. Sakagami, K. Nakao, C. Yoo and R. Takahashi for useful comments and conversations related to the topic in the present paper."
    },
    "0601/astro-ph0601270_arXiv.txt": {
        "abstract": "We observed the embedded, young  8--10~\\Mo\\ star AFGL~490  at subarcsecond resolution with the Plateau de Bure Interferometer in the \\co\\ (2--1) transition and found convincing  evidence that AFGL~490 is surrounded by a rotating disk. Using two-dimensional modeling of the physical and chemical disk structure coupled to line radiative transfer, we constrain its basic parameters. We obtain a relatively high disk mass of 1 \\Mo\\ and a radius of $\\sim$ 1\\,500 AU. A plausible explanation for the apparent asymmetry of the disk morphology is given. ",
        "introduction": "The study of very young massive stars is of particular importance for star formation since it is not yet firmly established by what process, disk accretion or stellar merging, such stars form \\citep*[e.g.,][]{YS02,Dobbsea05}. Previous authors have demonstrated that massive stars can be formed via  accretion, leading to the formation of large ($\\la 10\\,000$~AU) and massive ($\\la 1$~\\Mo) circumstellar disks. However, most of these stars are located more than 2~kpc away, which makes it challenging to  confirm {\\em unambiguously} the presence of accretion disks around single objects with radio interferometry due to the lack of spatial resolution.    We focus on AFGL 490 -- one of the best-studied young stars that are in a transition stage to Herbig~Be stars. This very young star is located nearby \\citep[$1\\pm0.3$~kpc,][]{Snellea84}, has a mass of 8--10~\\Mo, a bolometric luminosity of $\\sim 2\\times10^3$~\\Lo\\ \\citep[B2--3 star,][]{Panagia73}, is still embedded in its parent molecular cloud \\citep[A$_{\\rm V} \\sim 40$~mag,][]{ACK89}, and drives a high-velocity outflow \\citep{Mitchellea95}. Interferometric measurements by e.g.\\ \\citet[hereafter Paper~I]{Schreyerea02} as well as IR observations of \\citet*{Campbellea86} and \\citet*{Bunnea95} have revealed that the star is located inside an ionized $\\sim 100$~AU cavity cleared up by a fast stellar wind and surrounded by a $\\la 500$~AU disk-like structure, which is enshrouded in a 22\\,000~AU $\\times$ 6\\,000~AU envelope. All these structures are further embedded in an even more extended envelope.    In this Letter, we present subarcsecond-resolution  PdBI observations of AFGL~490 that prove the existence of a rotating disk around AFGL~490 and determine its orientation, size, and mass by using comprehensive physical, chemical, and line radiative transfer modeling. ",
        "conclusions": "Why do we see such a clumpy disk structure? This object seems to be relatively young and could still be in a perturbed state that remains from an earlier, non-steady accretion phase. The viscous dissipation timescale for the AFGL~490 disk is estimated to be $\\sim$ 1~Myr \\citep{Pringle81}, which far exceeds its age. This hypothesis is further supported by the fact that the $\\sim 10^4$~years old, large-scale CO outflow consists of single moving gas clumps, and thus accretion is indeed non-steady \\citep{Mitchellea95}. Another explanation could be that the \\co\\ (2--1) emission traces the densest parts of the spiral arms that are easily excited in massive disks by gravitational instabilities. Indeed, the morphology of the \\co\\ emission resembles the density structure of a self-gravitating disk with double arms (upper corner in Fig.~\\ref{co}a) calculated by \\citet{Fromangea04}. Moreover, the numerical simulations by \\citet*{Fromangea04a} imply that such disks should also develop a dual structure composed of {\\em an inner thin Keplerian disk} fed by a thicker self-gravitating torus with nearly uniform rotation. Finally, a clumpy disk structure could be a result of a recent encounter with a nearby star or due to the gravitational interaction with {\\rm a wide low-mass companion}. The complex multiple outflow systems seen in the close vicinity of AFGL~490 strongly support this idea (see Fig.~4b in Paper~I).  In this Letter, we present clear evidence that $\\sim10$~\\Mo\\ stars can be surrounded by Keplerian disks in their earliest evolutionally phase \\citep[see also][]{shepherd}. In contrast to the previously reported huge disks around such stars with radii of $\\la 10\\,000$~AU (e.g., Chini et al.\\ 2004, but see Sako et al.\\ 2005) the AFGL~490 disk is much smaller, $R\\sim 1\\,500$~AU. Using advanced theoretical modeling, we constrain basic disk parameters: (1) the inclination and positional angles are $\\sim 30\\degr$, (2) the surface density profile is $\\sim-1.5$, (3) the mass is $\\sim$ 1.5~\\Mo\\ ($\\approx 50\\%$ uncertainty), and (4)  the disk rotation is close to  Kepler's law."
    },
    "0601/astro-ph0601046_arXiv.txt": {
        "abstract": "We present preliminary results on spectroscopic data of the candidate pre-cataclysmic variable LTT 560. A fit to the flux-calibrated spectrum reveals the temperature of the white-dwarf primary to be $T_{\\rm eff} = 7000 - 7500$ K, and confirms the result of previous studies on the detection of an M5V secondary star. The analysis of radial velocity data from spectral features attributed to the primary and the secondary star show evidence for low-level accretion. ",
        "introduction": "LTT 560 is a high proper motion system from the \\inlinecite{luyt57} catalogue. A first photometry has been presented by \\inlinecite{egge68}, who marked the object as a possible white dwarf (WD). It was later included as a nova-like variable in \\inlinecite{vogt89}, and appeared as such in the \\inlinecite{down+97} catalogue of cataclysmic variables (CVs) with the designation Scl2. A first spectrum was published by \\inlinecite{hoarwach98}, who found narrow H$\\alpha$ emission and a red continuum that was fitted well by an M5 dwarf. Their data did not cover the blue part of the spectrum, so that they could not identify any contribution from the WD. However, a variable doubling of the H$\\alpha$ emission in their spectra let them deduce the binary nature of the object. Light curves presented by \\inlinecite{tapp+04} showed ellipsoidal variation with a period $P_{\\rm ph} = 3.54$ h, thus revealing LTT 560 as a potential pre-CV. ",
        "conclusions": "The general composition of LTT 560 appears clear. It is a binary star consisting of an M5-6V secondary star and a cold ($T_{\\rm eff} \\approx 7000$ K) WD primary. The current orbital period amounts to 3.54 h, and, if the pre-CV configuration applies, the system should thus evolve into a semi-detached CV within Hubble time.  The derivation of other system parameters still need a more thorough examination of the present data. The velocities of the WD absorption line are not yet reliable due to an insufficiently precise measurement of the WD absorption line. They will therefore be revised using a more sophisticated method. Furthermore, the H$\\alpha$ velocities are probably distorted due to the presence of an additional component.  Emission line doubling has been detected, e.g. also in the candidate pre-CV HS 2237+8154 \\cite{gaen+04}, but there it is due to a component that appears to be stationary at the system velocity. In LTT 560, this is clearly not the case. Since it is roughly in counterphase with the other component and the TiO band, it has to origin in some region beyond the centre of mass as seen from the secondary star, i.e. in the vicinity of the WD. A rough estimate of its semi-amplitude yields $K \\approx 100 ~{\\rm km~s}^{-1}$, thus, with respect to TiO, leading to much more plausible system parameters, e.g.\\ a mass ratio $q \\approx 0.5$. Still, this has to be confirmed by a more detailed analysis.  We conclude that our data shows evidence for emission from, or close to, the WD. It might be caused by sporadic accretion, e.g. due to wind from the secondary star. Accretion due to Roche lobe overflow appears less likely, since the late spectral type of the secondary implies that a semi-detached configuration is not expected for periods longwards of 2 h (e.g. \\opencite{beue+98})."
    },
    "0601/gr-qc0601001_arXiv.txt": {
        "abstract": "We  investigate the  non-linear coupling between  radial   and   non-radial oscillations of static  spherically  symmetric neutron stars as a possible  mechanism for the generation of gravitational waves that may lead to observable signatures. In this paper we concentrate on the axial sector of the  non-radial perturbations.  By using a multi-parameter perturbative framework we introduce a complete description of the non-linear coupling between radial and axial non-radial  oscillations; we  study the gauge   invariant character of the associated perturbative   variables  and  develop  a   computational  scheme  to evolve the non-linear coupling perturbations in the time domain.  We present results of simulations corresponding to different  physical  situations  and  discuss the dynamical behaviour of this non-linear coupling.  Of particular interest is the occurrence of signal amplifications in the form of resonance phenomena when a frequency  associated  with the radial pulsations is close  to a frequency  associated  with  one of the  axial  $w$-modes of  the star. Finally, we mention possible extensions  of  this work  and improvements towards more astrophysically motivated scenarios. ",
        "introduction": "Gravitational wave astronomy is currently developing as a new way of studying the cosmos, with a number of different experiments starting operating or being developed: earth-based interferometers (LIGO~\\cite{ligo}, VIRGO~\\cite{virgo}, GEO600~\\cite{geo} and TAMA~\\cite{tama}), resonant bars (EXPLORER, AURIGA, NAUTILUS, ALLEGRO and NIOBE; see~\\cite{bars}) and spheres such as MiniGRAIL~\\cite{minigrail}, as well as the laser interferometer space antenna LISA~\\cite{lisa}.  The scientific success of these detectors in producing new physical and astrophysical knowledge depends on the amount of {\\em a priori} available theoretical understanding about the different sources of gravitational waves.  It is then of outmost importance to have good theoretical descriptions of the systems that are likely to be observed by the ongoing gravitational-wave experiments.   Due to the weak character of the gravitational interactions there is only a few number of systems that can generate gravitational waves detectable for the experiments just mentioned.  In general, the dynamics of these systems involves strong gravitational fields that need to be described, using different levels of approximations, by general relativity.  This is the case in the modeling of compact objects such as neutron stars and supernov\\ae\\ core collapse, where non-linearity plays a role in the dynamics.  For these systems, relativistic perturbation theory is an excellent approach.  The linear theory has been used for a long time to study their oscillations and instabilities~\\cite{Andersson:2002ch,Kokkotas:2002ng}.  However, we know relatively little about non-linear dynamical effects (mainly due to numerical studies~\\cite{Sperhake:2001xi,2002PhRvD..65b4001S,2002ApJ...571..435M,2002PhRvD..65h4039L,2003ApJ...591.1129A,Font:2001ew,Stergioulas:2003ep,Dimmelmeier:2005zk}). Despite the fact that strong non-linear effects require a fully non-perturbative approach, it is reasonable to expect that there are interesting physical phenomena that only involve a mild non-linearity for which a second-order treatment should be perfectly suited.  In this sense, a non-linear perturbative approach is therefore timely and may lead to a better understanding of the mechanisms of generation of gravitational waves.   An interesting scenario to study is a neutron star which is oscillating radially and non-radially.  At first order radial oscillations of a spherical star don't emit {\\it per se} any gravitational waves, but they can drive and possibly amplify non-radial oscillations and then produce gravitational radiation to a significant level. In addition one may expect the appearance of non-linear harmonics, which may also come out at lower frequencies than the linear modes~\\cite{Zanotti:2004kp,Dimmelmeier:2005zk}, where the Earth-based laser interferometers have a higher sensitivity. An additional motivation is that there are a number of studies aiming at investigating if non-radial oscillations of stars can be excited by external sources (see e.g.\\ \\cite{Gualtieri:2001cm,Poisson:1993vp}). However, our idea is to investigate whether non-radial oscillations can be driven or even amplified through coupling by an internal radial oscillation, regardless of the presence of an external source.  These non-linear processes can occur, for instance, in a proto-neutron star that is still pulsating.  A mainly radial pulsation could for example drive the non-radial oscillations, either naturally present, or excited through fall-back accretion.   In recent papers~\\cite{Bruni:2002sm,Sopuerta:2003rg} a non-linear multi-parameter perturbative formalism has been developed, which can be applied to the study of the non-linear dynamics of the oscillations of neutron stars, as shown in~\\cite{Passamonti:2004je} where it has been further developed to study a particular second-order effect: the coupling of radial and non-radial first order perturbations of a compact spherical star.  One of the main ideas behind this work is that it is very convenient to treat the two sets of perturbations, radial and non-radial, as separately parametrized and then to study, at the next perturbative order, the way in which they couple, which corresponds to a particular sector of the second-order perturbations, the one formed by the product of radial and non-radial perturbations. In~\\cite{Passamonti:2004je}, the equations describing the coupling were obtained for the case of polar non-radial perturbations and gauge-invariant variables were also found (having fixed the gauge for the radial ones for simplicity).   In this paper, we extend this study to the case of axial perturbations. At first order, axial perturbations are decoupled from fluid perturbations, but at the next perturbative order, the terms describing the coupling are driven by the radial pulsations. Here we discuss the results of simulations in the time domain that show how the coupling mechanism works and under which conditions we may have a situation in which this coupling produces amplifications due to resonant phenomena.  The plan of this paper is the following: In Section~\\ref{framework} we describe the perturbative framework we use in this work.  In Section~\\ref{ingredients} we introduce all the necessary ingredients required by the perturbative scheme in order to study the coupling of radial and axial non-radial oscillations of neutron stars.  In Section~\\ref{numerics} we describe the structure of the computational framework we have implemented to solve the perturbative equations in the time domain. In particular, we discuss in detail the construction of initial data and the numerical treatment of the boundary conditions. In Section~\\ref{results} we present the results of our simulations for two different physical scenarios: (i) the scattering of a gravitational wave packet by a differentially rotating star, and (ii) the non-linear coupling in a differentially rotating and radially oscillating compact star. In particular we discuss the appearance of amplifications due to resonances of the system.  We conclude in Section~\\ref{conclusions}, where we summarize the main results of this paper and discuss potential extensions of the approach we have used.  In the Appendices we give some proofs of the gauge invariance character of some of the perturbative quantities (Appendix~\\ref{gaugeinvariance}), expressions for the source terms in the perturbative equations describing the coupling (Appendix~\\ref{sourceterms}), and the equations for radial perturbations in terms of the {\\em tortoise} fluid coordinate (Appendix~\\ref{equationstortoise}).  Throughout this work we use the following conventions: We use Greek letters to denote spacetime indices; capital Latin letters for indices in the time-radial part of the metric; lower-case Latin indices for the spherical sector of the metric.  We use physical units in which $8\\pi G = c = 1\\,$. ",
        "conclusions": "\\label{conclusions} In this paper we have investigated non-linear effects in the dynamics of compact stars due to the coupling of radial and axial non-radial oscillations, and the potential impact that this mechanism may have in the emission of gravitational radiation.   To that end, we have used and extended the multi-parameter non-linear perturbative scheme introduced in~\\cite{Bruni:2002sm,Sopuerta:2003rg}, and further developed in~\\cite{Passamonti:2004je} for the case in which the coupling is between radial and polar non-radial oscillations.  The non-linear perturbative equations and the gauge-invariant character of the perturbative quantities describing the coupling have been derived by using the 2-parameter perturbation theory in connection with the GSGM formalism~\\cite{Gerlach:1979rw,Gundlach:1999bt,Martin-Garcia:2000ze}. As expected, in the stellar interior, the perturbative variables describing the coupling satisfy inhomogeneous linear equations.  The form of these equations is such that the associated homogeneous equations are constructed from the same linear operators as the equations for  the first order axial non-radial perturbations.   The source terms are quadratic in the first-order radial and axial non-radial perturbative quantities.  In the exterior we do not have source terms and hence, the dynamics is described by Regge-Wheeler-type equations.  Interior and exterior communicate through the junction conditions at the surface of the star. We have discussed in detail the initial data setup used in the simulations presented in this paper as well as the different boundary conditions. The initial data for the radial oscillations corresponds to specific eigenmodes.  In the case of the axial non-radial perturbations we have used two types of initial data: (i) static profiles representing differential rotation of the star from the relativistic $j$-constant rotation law, and (ii) Gaussian packets containing a number of axial eigenmodes. Regarding the boundary conditions, of particular importance are the conditions at the stellar surface.   In order to avoid negative values of the mass energy density near the surface due to the Eulerian description of the radial and non-radial polar perturbations, we have adopted an approximation already used in the literature~\\cite{Sperhake:2001si,Sperhake:2001xi,2003PhRvD..68b4002H}, that consists in removing outer layers of the neutron star, which is equivalent, in most of the cases we deal with, to neglect less than one percent of the total gravitational mass of the star. This approximation leads to a description of the second-order coupling effects which is accurate to better than five per cent.  We have presented a computational framework for investigating, in the time domain, the evolution of the axial non-linear perturbations describing the coupling.  Our numerical codes are based on finite difference methods and standard explicit evolution algorithms.  Their structure is hierarchical: First, we solve the TOV equations; then, we evolve a time step the first-order radial and non-radial perturbations, and with the result we update the source terms for the non-linear perturbations; eventually, we evolve them in the same time step.  We have then studied two different physical scenarios: (i) The scattering of a gravitational wave by a radially pulsating star.  (ii) A differentially rotating and radially oscillating star.  The simulations for the first scenario have shown that the correction onto the gravitational wave signal due to the non-linear coupling are of the order of $2\\,\\%$ or less when the radial oscillations are in the fundamental mode, and of the order of $0.1\\,\\%$ or less when we consider higher radial overtones.  Moreover, the properties of the waveforms and spectra are essentially unchanged with respect to the first-order ones.   The simulations for the second scenario have provided more interesting results.  The waveforms we obtain have the following pattern: at the early stages of the evolution they present an initial $w$-mode burst related to the initial data choice. At later stages, the signal becomes periodic and is driven by the radial pulsations through the sources.  This is confirmed by an analysis of the spectra, which show that the gravitational wave signal contains the frequencies of the radial eigenmodes prescribed in the initial data.   More interestingly,  we have found a resonance effect that takes place as the frequencies of the radial eigenmodes considered get close to the frequency of the first $w$-mode. For the stellar model used in this paper, the amplitude of the gravitational wave signal associated with the fourth radial overtone (the one with the closest frequency to the first $w$-mode) is about three orders of magnitude higher than that associated with the radial fundamental mode.  We have also given a rough estimate of the damping times of the radial pulsations due to the gravitational wave emission. The values found depend strongly on the presence of resonances.  In the case of differential rotation with a $10$ ms rotation period at the axis and a differential rotation parameter $A=15$ km, the fundamental radial mode gets damped after about ten billion oscillation periods, while the fourth overtone after only ten.  This is not surprising, and shows that the coupling near resonances is a very effective mechanism for extracting energy from the radial oscillations.  In this sense, it is important to mention that the detection of such a gravitational wave signature would provide important information about the stellar physical properties, specially because it has information on radial pulsations, which can be determined easily for a large class of equations of state.   The work presented in this paper represents a first step in the study of the non-linear coupling of oscillations of compact relativistic stars and its impact for gravitational wave physics.  The next step in this study is the description of the coupling between radial and polar non-radial oscillations. The spectrum of the polar non-radial perturbations is richer than the axial case due to the presence of the fluid modes, which may have frequencies lower than the spacetime modes and a longer gravitational damping.   As a result, we may expect a more effective coupling with the radial pulsation modes and more channels for possible resonant amplifications.  Other extensions of this work include the consideration of more realistic models of the stellar structure, by taking into account the effects of rotation, composition gradients, magnetic fields, dissipative effects, etc.  In particular, it would be interesting to investigate a proto-neutron star model to compare the damping rates due to gravitational radiation produced by the coupling of oscillations with the strong damping induced by the presence of a high-entropy envelope surrounding the newly created neutron star. By including rotation we may find new interesting non-linear effects due to the different behaviour of radial and non-radial modes in a rotating configuration. While the radial modes are only weakly affected so that their spectrum is essentially the same as in a non-rotating configuration (when scaled by the central density), the non-axisymmetric modes manifest a splitting similar to the Zeeman effect in atomic spectra. The rotation removes the mode degeneracy in the azimuthal harmonic number of the non-rotating case.   The details of this frequency separation depend on the stellar compactness and rotation rate.   It may then happen that for a given stellar rotation rate and compactness the non-radial frequencies cross the sequence of radial frequencies~\\cite{Dimmelmeier:2005zk} so that the frequencies of these two kinds of modes are comparable, and possible resonances or instabilities could influence the spectrum and the gravitational wave signal. Finally, as we have already mentioned above, it would be also interesting to explore the impact of gravitational back-reaction on the dynamical coupling mechanism (see, e.g.~\\cite{Balbinski:85be,Yoshida:95ye}).    \\[ \\] {\\bf Acknowledgements:} We thank Nils Andersson, Valeria Ferrari, Ian Hawke, Kostas Kokkotas, Pablo Laguna, Giovanni Miniutti, Ulrich Sperhake and Nick Stergioulas for fruitful discussions. AP work on this project is partly funded by a ``VESF'' grant. MB work was partly funded by a \"Cervelli\" fellowship from  MIUR (Italy). CFS acknowledges the support of the Center for Gravitational Wave Physics funded by the National Science Foundation under Cooperative Agreement PHY-0114375, and support from NSF grant PHY-0244788 to Penn State University.    \\appendix"
    },
    "0601/astro-ph0601336_arXiv.txt": {
        "abstract": "In the region of Sgr B2 there are several condensations heated externally by nearby hot stars. Therefore H$_2$O far--IR lines are expected to probe only an external low--density and high temperature section of these condensations, whereas millimeter-wave lines can penetrate deeper into them where the density is higher and T$_k$ lower. We have conducted a study combining H$_2$O lines in both spectral regions. First, $\\textit{Infrared Space Observatory}$ observations of several H$_2$O thermal lines seen in absorption toward Sgr~B2(M) at a spectral resolution of $\\sim$35~km~s$^{-1}$ have been analyzed. Second, an \\textit{IRAM}--30m telescope map of the para--H$_2$O $3_{13}-2_{20}$ line at 183.31~GHz, seen in emission, has also been obtained and analyzed. The H$_2$O lines seen in absorption are optically thick and are formed in the outermost gas of the condensations in front of the far--IR continuum sources. They probe a maximum visual extinction of  $\\sim$5 to 10 mag. Radiative transfer models indicate that these lines are quite insensitive to temperature and gas density, and that IR photons from the dust play a dominant role in the excitation of the involved H$_2$O rotational levels. In order to get the physical conditions of the absorbing gas we have also analyzed the CO emission toward Sgr~B2(M). We conclude, based on the observed CO $J$=7--6 line at 806.65~GHz with the \\textit{Caltech~Submillimeter~Observatory}, and the lack of emission from the far--IR CO lines, that the gas density has to be lower than $\\sim$10$^4$~cm$^{-3}$. Using the values obtained for the kinetic temperature and gas density from OH, CO, and other molecular species, we derive a water column density of  (9$\\pm$3)$\\times$10$^{16}$~cm$^{-2}$ in the absorbing gas. Hence, the water vapor abundance in this region, $\\chi$(H$_2$O), is $\\simeq$(1-2)$\\times$10$^{-5}$. The relatively low H$_2$O/OH abundance ratio in the region, $\\simeq$2-4, is a signature of UV photon dominated surface layers traced by far--IR observations. As a consequence the temperature of the absorbing gas is high, T$_K\\simeq$300-500 K, which allows very efficient neutral--neutral reactions producing H$_2$O and OH. On the other hand, the 183.31 GHz data provide a much better spatial and spectral resolution than the far-IR ISO data. This maser line allows to trace water deeper into the cloud, i.e., the inner, denser ($n(H_2)$$\\ge$10$^{5-6}$~cm$^{-3}$) and colder (T$_k$$\\sim$40 K) gas. The emission is very strong toward the cores. The derived water vapor abundance for this component is a few$\\times$10$^{-7}$. There is also moderate extended emission around Sgr~B2 main condensations, a fact that supports the water vapor abundance derived from far--IR H$_2$O lines for the outer gas. ",
        "introduction": "\\label{sct:intro} The determination of water  abundance in space is a long standing problem in Astronomy. Theoretical models predict that  water can be the most abundant species in \\textit{warm} molecular clouds after H$_2$ and CO \\citep{Neu95}. Therefore, the determination of its spatial distribution and abundance contributes to a better knowledge of the chemical  and physical processes that take place in the interstellar medium (ISM).  Unfortunately, water is an abundant molecule in our  atmosphere making particularly difficult the observation  of its rotational lines and vibrational bands from Earth. Even so, some observations of water lines have been performed from ground--based and airborne telescopes: the $6_{16}-5_{23}$  at 22~GHz \\citep{Che69}, the $5_{15}-4_{22}$  at 325~GHz \\citep{Men90}, the $10_{29}-9_{36}$ at 321~GHz \\citep{Men90b} and the $3_{13}-2_{20}$ at 183.31~GHz (Watson et al. 1980; Cernicharo et al. 1990, 1994, 1996, 1999a; Gonz\\'alez-Alfonso et al., 1998). Due to the maser nature of these lines, the analysis and interpretation of the spectra is not obvious. Among these lines, only two have been used to map the extended emission of water vapor in Orion: the $3_{13}-2_{20}$ -hereafter the 183 GHz line- \\citep{Cer94} and the $5_{15}-4_{22}$ at 325 GHz \\citep{Cer99a}. However, although we know from ISO and SWAS observations that water is extended ($\\sim$25$'$$\\times$25$'$) in Sgr~B2 (Cernicharo et al. 1997, Neufeld et al. 2003, Goicoechea et al. 2004), little is known about its excitation conditions and its detailed spatial distribution. An alternative to indirectly estimate the water abundance in the Galactic Center (GC) is to use related species such as HDO \\citep{Jac90,Com03} or H$_3$O$^+$ \\citep{Phi92,Goi01}. In none of these cases, the determination of $\\chi$(H$_2$O) is straightforward.   The ISO mission \\citep{Kes96}, and specially, the \\textit{Long Wavelength Spectrometer}, LWS, \\citep{Cle96}, and the \\textit{Short Wavelength Spectrometer}, SWS, \\citep{Thi96}, have provided a unique opportunity to observe several H$_2$O lines in a great variety of astronomical environments. Nevertheless, the majority of these observations were performed at the low spectral resolution of the grating mode  ($\\sim$1000~km~s$^{-1}$), which produces a critically strong dilution in the search for molecular features in most  ISM sources. Nevertheless, the Sgr~B2 cloud has been analyzed and studied in detail with the LWS/\\textit{Fabry--Perot}, which provided a velocity resolution of $\\sim$35~km~s$^{-1}$ (Goicoechea et al. 2004, hereafter G04).    Opposite to what is found toward other star forming regions such as Orion \\citep{Cer99b}, the observations of the $2_{12}-1_{01}$ line  at $\\sim$179.5~$\\mu$m in Sgr B2 show that the line appears in absorption rather than in emission \\citep{Cer97}. Afterward, the launch of the \\textit{Submillimeter Wave Astronomy Satellite}, SWAS, \\citep{Mel00}, and ODIN \\citep{Nor03}  allowed the observation of the $1_{10}-1_{01}$ fundamental transition of both H$_{2}^{16}$O at 557 and H$_{2}^{18}$O at 548~GHz, first detected by the \\textit{Kuiper Airborne Observatory}, KAO, \\citep{Zmu95}. Although the velocity resolution is $\\sim$1~km~s$^{-1}$, the large beam of SWAS ($\\sim$4$'$) makes these observations more sensitive to the cold and less dense gas. SWAS observations have provided a reliable estimate of the water vapor abundance in the low excitation clouds located in the line of sight toward Sgr~B2 \\citep{Neu00,Neu03}.  However, the fact that only the ground-state absorption line is detected makes difficult a detailed study of water vapor excitation mechanisms in Sgr~B2 itself. This is the most massive cloud in the Galaxy, with $\\sim$10$^7$ \\Msun, \\citep{Lis90}, and a paradigmatic object in the GC  region as its geometrical properties, physical conditions and chemical characteristics make it a miniature galactic nucleus with $\\sim$15$'$ extent  (G04 and references therein).  The main star-forming regions in Sgr~B2 are located within three dust condensations, labeled as (N), (M) and (S). These condensations are embedded in a $\\sim$10~pc moderate--density ($n(H_2)=10^5-10^6$~cm$^{-3}$) cloud \\citep{Lis91,Hut93}. In addition, these structures are surrounded by lower density components of warm (T$_k$$\\geq$100 K) gas, hereafter the Sgr~B2 warm envelope. These conditions have been derived mainly from absorption observations of NH$_3$ metastable lines \\citep{Wil82,Hut95}, OH lines (Goicoechea \\& Cernicharo, 2001; 2002), and H$_2$CO lines \\citep{Mar90} in the radio domain. Nevertheless, the warm and low density gas is poorly traced by radio observations of other molecular species (generally excited by collisions in the denser regions and thus observed in emission). However, the warm envelope represents the strongest contribution to the absorption features produced by many light hydrides in the far-IR spectrum of Sgr~B2 \\citep{G04}. The origins of the observed rich chemistry  and the  heating mechanisms in the Sgr~B2 warm envelope are a subject of intense debate. The matter is complicated  due to the different observational signatures to be integrated in the same picture: high temperatures derived from NH$_3$ absorption lines  \\citep{Hut95,Cec02}, fine structure emission from the photo--ionized and photo--dissociated gas \\citep{G04}, SiO and X--ray distribution \\citep{Mar00}, etc.  In all possible scenarios, water plays a significant role. Several mechanisms allow its formation and survival in the warm envelope. Dissociative recombination of H$_3$O$^+$ leads to the production of H$_2$O and OH. This processes depend on the specific $f_{H_2O}$ branching ratio for the H$_2$O formation channel. Unfortunately, the determination of $f_{H_2O}$ with different experimental procedures has also yielded different values, from $f_{H_2O}$=0.05 \\citep{Wil96} to $f_{H_2O}$=0.25 \\citep{Jen00}, while most chemical models have used $f_{H_2O}$$\\sim$0.35. In addition, water could also be produced in the gas phase by the endothermic reaction: \\begin{equation} OH + H_2 \\rightarrow H_2O + H \\end{equation} However, the gas temperature must exceed $\\sim$300~K to overcome the activation barrier \\citep{Neu95}. At these temperatures, the reaction: \\begin{equation} O + H_2  \\rightarrow OH + H \\end{equation} also contributes to the formation of OH. Therefore, H$_2$O and OH column densities can be used to determine the role of the neutral-neutral reactions in their formation/destruction routes. Still, the exact H$_2$O/OH ratio will be determined by $f_{H_2O}$, the temperature and by photodissociation processes if UV radiation is present. As an example, Neufeld et al. \\citep{Neu03} studied a diffuse cloud ($G_0$$\\sim$1, $n(H)$=100~cm$^{-3}$) toward W51 and showed that the presence of  a warm gas component (T$_k$$\\gtrsim$400~K) could explain the observed variations of the H$_2$O/OH ratio respect to other diffuse clouds.   Finally, high oxygen depletion onto water ice mantles in dust grains could have taken place during the evolution of the cool gas in Sgr~B2. Photodesorption and/or evaporation for dust temperatures above $\\sim$90~K, could release some water back into the gas phase enhancing the H$_2$O  abundance expected from pure gas--phase formation \\citep{Ber00}. However, gas and dust are thermally decoupled in the outer layers of Sgr~B2, where the dust temperatures are significantly lower, T$_d$$\\simeq$20--30~K \\citep{Gor93,G04}, than gas temperatures, T$_k$$\\simeq$300~K \\citep{Goi02, Cec02}. Thus, T$_d$ seems too low to produce significant evaporation of water ice mantles. Therefore, a detailed  study  of the far--IR H$_2$O  lines and of the 183.31~GHz extended emission is needed to constrain the water abundance and the physical characteristics of the absorbing/emitting region.     In this work we present and analyze the far-IR observations of several thermal lines of water vapor  toward Sgr~B2(M) and the first map of the 183.31~GHz maser emission of para--H$_2$O around Sgr~B2 main condensations. The layout of the paper is as follows: In Section \\ref{sct:observ} we summarize the far--IR, submm and mm observations and data reduction. The spectra and maps are presented in Sec.\\ref{sct:results}. Section \\ref{sct:analysis} is devoted to the analysis of CO observations (\\ref{ssct:co}) and water vapor observations (\\ref{ssct:h2o}) with different radiative transfer methods. The main implications of our work are discussed in Sec. \\ref{sct:discussion}, where photochemistry models for H$_2$O and OH are also presented. A summary is given in Sec. \\ref{sct:summary}. ",
        "conclusions": "\\label{sct:analysis}  \\subsection{Carbon Monoxide} \\label{ssct:co} The warm gas present in the outer layers of Sgr~B2(M) might be expected to radiate in high--$J$ CO lines. The CO $J$=14--13 transition at 185.999~$\\mu$m is the one with the lowest energy level (E$_u$$\\simeq$581~K) within the range of ISO/LWS detectors. However, we have not detected any emission/absorption from CO (3$\\sigma$ limits are $<$2$\\times$10$^{-18}$~W~cm$^{-2}$) in the ISO/LWS spectra toward the Sgr~B2 region \\citep{G04}, with both grating and FP. CO spectroscopical data and line flux upper limits are tabulated in Table~\\ref{tab-co_lines}. Nevertheless, recent  studies of the large scale CO $J$=7--6 emission toward the GC with the \\textit{Antarctic Submillimeter Telescope and Remote Observatory} (AST/RO) have shown that the emission is concentrated toward Sgr~B and Sgr~A complexes \\citep{Kim02}. Fig.~\\ref{co_7_6_cso} shows higher spectral/angular resolution observations of this line taken with the CSO telescope toward the Sgr~B2(M) position. Hence, it seems that at a given $J_{up}$ level, the rotational CO line emission disappears from the Sgr~B2 spectrum.  \\begin{figure}[t]% \\centering \\begin{center} \\includegraphics[angle=0, width=8cm]{f3.eps} \\caption{CSO observation of the CO $J$=7--6 line at 806~GHz toward Sgr~B2(M).} \\label{co_7_6_cso} \\end{center} \\end{figure}%    \\begin{figure*}[t]% \\centering \\begin{center} \\includegraphics[width=8cm,angle=-90]{f4.eps} \\caption{Nonlocal models of CO rotational lines toward Sgr~B2. Velocity--integrated line fluxes (in 10$^{-17}$ W cm$^{-2}$) are shown as a function of $J_{up}$, the CO upper level rotational number. From $J$=14--13,  CO lines are  accessible to the ISO/LWS (vertical dashed line). No CO line has been detected. 3$\\sigma$ upper limits for CO line fluxes are shown (filled squares). Two different abundances have been considered, $\\chi$(CO)=10$^{-4}$ (upper panels) and $\\chi$(CO)=10$^{-5}$ (lower panels). Models with different kinetic temperatures (100 to 500~K) are shown in each panel for different H$_2$ volume densities: 5.0$\\times$10$^3$, 1.0$\\times$10$^4$, 2.0$\\times$10$^4$, 4.0$\\times$10$^4$, 8.0$\\times$10$^4$, 1.6$\\times$10$^5$, 3.2$\\times$10$^5$, 3.2$\\times$10$^5$ and 6.4$\\times$10$^5$~cm$^{-3}$.} \\label{co_models} \\end{center} \\end{figure*}%    We have performed nonlocal radiative transfer calculations to try to reproduce the lack of high--$J$ CO lines in the far--IR spectrum of Sgr~B2(M), and to help constraining the physical parameters needed to model the water vapor absorption. In Fig.~\\ref{co_models} we show the predictions of a nonlocal model for several high--$J$ transitions of CO. The nonlocal model used here is an adaptation of the radiative transfer code developed by Gonz\\'alez--Alfonso \\& Cernicharo \\citep{Gon93} with the inclusion of dust in the transfer \\citep{Cer00}. We have implemented it for high--$J$ levels of CO (see Sec. \\ref{sssct:nonlocal} for further details on the specific model developed for Sgr~B2(M)). The collisional rates have been taken from Flower (2001). The dust continuum emission has been modeled following Goicoechea and Cernicharo (2002) and G04 (see section 4.2.1). We have run an array of models for CO abundances\\altaffilmark{5} of 10$^{-4}$ and 10$^{-5}$ with T$_k$ varied from 100 to 500~K and $n(H_2)$ from 5.0$\\times$10$^3$ to 6.4$\\times$10$^5$~cm$^{-3}$. From these results it is clear that in order to suppress the far--IR CO emission, and to match the CO $J$=7--6 line emission, low H$_2$ density is required. In particular, if the CO $J$=7--6 line arises from a layer of gas at T$_k$$\\sim$100~K the model with $n(H_2)$=2$\\times$10$^4$~cm$^{-3}$ correctly reproduces the absence of far--IR lines. However, if it arises from the same outer layer of warm OH (T$_k$$\\sim$300~K) detected in the far--IR \\citep{Goi02}, the limit for the density will be $<$10$^4$~cm$^{-3}$. In any case, the lack of high--$J$ CO lines is an  evidence of the low density gas located in front of the far--IR continuum source. In the following section we will constrain the physical parameters of this layer by studying the available water vapor lines.   \\footnotetext[5]{An upper limit to the  $^{13}$CO fractional abundance in Sgr~B2 of 10$^{-6}$ was found by Lis \\& Goldsmith \\citep{Lis89}. According to $^{12}$C$^{18}$O/$^{13}$C$^{18}$O$\\simeq$25 in Sgr~B2 \\citep{Lan90}, the $^{12}$CO abundance would be $\\lesssim$2.5$\\times$10$^{-5}$.}  \\subsection{Water Vapor} \\label{ssct:h2o} The main observational result of this work is that the far--IR water lines toward Sgr~B2(M) appear in absorption, while the 183.31~GHz line is seen in emission around and in the main condensations. The fact that the continuum emission in the far--IR is optically thick \\citep{G04} indicates that the H$_2$O absorption lines arise from regions where the excitation temperatures (T$_{ex}$) are smaller than the dust temperatures inferred from the continuum emission. However, the problem of the H$_2$O  line excitation toward Sgr~B2 is not straightforward. Apart from self--absorption and the possible excitation by collisions with molecules (e.g. H$_2$), atoms (e.g. He) and $e^-$ (if the ionization fraction is significant), the level population can be primarily determined by the thermal emission of dust, the role of which is essential in the excitation of molecules such as H$_2$O or OH which have many rotational lines in the far-- and mid--IR. This represents a major difference with respect to the excitation treatment of molecules observed in the radio domain where, generally, one can neglect the excitation by dust photons. If the excitation is dominated by IR photons, the two transitions arising from the ground levels of ortho-water, 2$_{12}$--1$_{01}$ and 1$_{10}$--1$_{01}$, will determine how the higher energy levels will be populated. Even if collisions are important, the presence of an optically thick far--IR continuum will strongly affect the T$_{ex}$ of water % and, thus, it must be carefully taken into account in the models. In the case of Sgr~B2, this means that the external dust layers of the cloud will absorb the possible water line emission from the inner regions. Knowledge of the geometry of the region to be modeled and the relative filling factors of the dust and gas in the beam of the LWS instrument, are also important for the models. For this reason, high angular resolution ground-based observations of the 183.31 GHz water line are particularly important to model the water vapor radiative transfer in Sgr~B2.     \\subsubsection{Large Velocity Gradient modeling} \\label{sssct:lvg} In this section we analyze the H$_2$O observations with a \\textit{multi-molecule Large Velocity Gradient} (LVG) model. For Sgr~B2(M) we have adopted a spherical geometry with two components (see Fig.~\\ref{fig_h2o_geometry}): a uniform continuum core with a diameter of $\\sim$23$''$ ($\\sim$1~pc for a distance of 8.5~kpc) and a shell of variable thickness and distance to the core. The presence of an external shell of molecular gas (not resolved by the ISO/LWS beam) surrounding  a central condensation is indicated by the analysis of  the far--IR OH \\citep{Goi02} and NH$_3$ lines \\citep{Cec02} and it is also supported by H$_2$O 183.31~GHz line observations (Fig.~\\ref{mapa_183}). In particular Goicoechea \\& Cernicharo (2002) found an angular size of $\\sim$42$''$ for the OH envelope. The bulk of the H$_2$O absorption lines must arise from the OH layers, or inside them. Following the OH model geometry and Figs.~\\ref{mapa_183} and \\ref{co_7_6_cso}, a total size of $\\sim$30$''$ ($\\sim$1.2~pc) for the core$+$shell cloud has been adopted. The core is considered as a gray--body with an opacity at 80~$\\mu$m of 2.5, with a dust opacity law given by $\\tau_{\\lambda}=\\tau_{80}[80/\\lambda(\\mu m)]$, and a dust temperature of 30~K. These values are consistent with the color temperatures and dust emissivities derived from the analysis of the ISO/LWS continuum observations at the same wavelengths of the detected far--IR H$_2$O lines \\citep{G04}. Although dust temperatures can be slightly larger (T$_d$$\\sim$60--80~K) inside the cloud (from the analysis of the millimeter continuum emission) or slightly lower (T$_d$$\\sim$20~K) in the outer and colder diffuse layers (from the analysis of extended IRAS continuum emission; Gordon et al. 1993), we judged T$_d$$\\simeq$30~K as the most representative value for the dust grains coexisting with the molecular species detected in the far--IR. The bulk of the continuum emission detected by ISO can be hardly fitted with dust temperatures below 30 K. The emission of higher temperature dust arising from the innermost regions of Sgr B2 is hidden in the FIR by the huge amount of foregroung gas and colder dust.    For these models we have considered an ortho--H$_2$O column density of 1.8$\\times$10$^{16}$~cm$^{-2}$ and a para--H$_2$O column density of 0.6$\\times$10$^{16}$~cm$^{-2}$. These are lower limits suggested by the nonlocal radiative transfer models (see below). The LVG model computes the statistical equilibrium population of the rotational levels for ortho--H$_2$O and para--H$_2$O independently. The collisional rates were scaled from those of H$_2$O--He collisions \\citep{Gre93}.    \\begin{figure*}[]% \\centering \\includegraphics[ width=13cm]{f5.eps} \\caption{Large Velocity Gradient excitation models of different kinetic temperature and density for several water vapor rotational lines. The column density of ortho--H$_2$O is 1.8$\\times$10$^{16}$~cm$^{-2}$ and that of para--H$_2$O is 0.6$\\times$10$^{16}$~cm$^{-2}$. In each panel, the H$_2$O transition and wavelength$/$frequency is labeled. The thick contour corresponds to the equivalent temperature of the continuum source (T$_c$). Hence, to the left of this contour, lines will be in absorption (T$_{ex}$$<$T$_c$).} \\label{lvg_mod} \\end{figure*}%   \\begin{figure*}[]% \\centering \\includegraphics[width=17cm]{f6.eps} \\caption{Large Velocity Gradient excitation models for the 3$_{13}$--2$_{20}$ maser transition of para--H$_{2}^{16}$O at 183.31~GHz. Models with and without a central continuum source (described as a gray--body with an opacity of 2.5 at 80~$\\mu$m, with a dust opacity law given by $\\tau_{\\lambda}=\\tau_{80}[80/\\lambda(\\mu m)]$, and a dust temperature of 30~K) have been considered. For low temperature and low densities, infrared pumping seems to be efficient to increase the emerging flux of the 183.3 GHz line. In addition the maser amplifies the continuum. Hence, the resulting flux depends on the assumed opacity at 80 $\\mu$m and on the opacity law. For higher kinetic temperatures and densities collisional pumping of the maser dominates and the difference between the two cases, with and without central continuum source, is less important. For a given kinetic temperature (20 to 500~K), the expected T$_B$ as a function of the H$_2$ density is shown in each panel for different column densities of para--H$_{2}^{16}$O. The dashed line indicates the observed brightness temperature in the extended emission. Toward the main condensations, brightness temperatures of $\\sim$100 K are reached.} \\label{lvg_mod_183} \\end{figure*}%   Several LVG computations for some selected ortho-- and para--H$_2$O transitions in the THz domain are shown in Fig.~\\ref{lvg_mod}. The excitation temperature T$_{ex}$ of each transition is shown in each panel as a function of T$_k$ and $n(H_2)$. The thick contour corresponds to the equivalent temperature of the continuum core. Therefore, to the left of this contour, water lines appear in absorption. From the LVG models it is clear that high density and temperature are required to observe the far--IR water lines in emission. For the range of densities implied by the CO observations ($<$10$^4$~cm$^{-3}$), the  H$_2$O lines observed by the ISO/LWS are correctly predicted in absorption. The temperature of the absorbing layer is, however, more difficult to estimate because  the possible solutions for a given (low) density model are not very sensitive to  temperature variations, as indicated by the smooth change of T$_{ex}$ for constant density as the kinetic temperature goes from 500 to 100 K. LVG models predict that, in addition to dust photons, collisions play a role in the excitation of the lowest H$_2$O rotational levels. The higher energy levels are pumped from the lowest levels by absorption of far--IR photons. In the case of the ortho--H$_2$O 1$_{10}$--1$_{01}$ line observed by \\textit{SWAS} (Neufeld et al. 2000), only moderate densities ($>$5$\\times$10$^4$~cm$^{-3}$) are needed to observe the line in emission. This is the case of the extended  emission in the 1$_{10}$--1$_{01}$ line observed in the \\textit{180~pc Molecular Ring} around the GC between $v_{LSR}$$\\sim$$+$80 and 120~km~s$^{-1}$ (Neufeld et al. 2003). For the density conditions derived for the Sgr~B2 velocity range and for the line of sight clouds, the 557~GHz line is correctly predicted in absorption. A similar behavior is expected for the para--H$_2$O 1$_{11}$--1$_{00}$ line at $\\sim$269.3~$\\mu$m ($\\sim$1113~GHz)  that will be observed by  future heterodyne instruments such as HIFI/Herschel.  Among all the moderate excitation  para--H$_2$O lines (E$_l$$<$450~K), only the 3$_{13}$--2$_{20}$ appears in the mm domain. Contrary to other H$_2$O maser lines accessible from ground--based telescopes, relatively low T$_k$ and  density are required for the 3$_{13}$--2$_{20}$ line inversion. These conditions allowed the first detection of extended water emission in Orion (Cernicharo et al. 1994), and a strong dependence of the emission with T$_k$ was revealed. Hence, the 183.31~GHz line could be an excellent tracer of the warm gas in molecular clouds.     We have used the same LVG model to analyze the 183.31~GHz line inversion mechanism that produces  extended emission in Sgr~B2 (Fig.~\\ref{mapa_183}). Fig.~\\ref{lvg_mod_183} shows different excitation models  for the 3$_{13}$--2$_{20}$ transition. Since a fraction of the 183.31~GHz line emission may arise from the warm envelope in front of the far--IR continuum source, models with and without a central continuum source (the same described at the beginning of the section) have been considered. For a given T$_k$ (from 20 to 500~K), each panel shows computations of different para--H$_2$O column densities as a function of  T$_B$ (in K) and  H$_2$ volume density. For low temperatures (T$_k$$<$40~K) and moderate densities ($<$10$^5$~cm$^{-3}$), the line will be observable for large $N$(p--H$_2$O) values only if a continuum source is present. Due to the minor role played by collisions in the pumping of the rotational levels at these temperatures, the line intensity is almost independent of the density. Nevertheless, as the temperature increases, collisions start to be significant, and the expected intensity of the 183.31~GHz line becomes less sensitive to the models with or without the continuum source.  \\begin{figure}[t]%  \\includegraphics[angle=0, width=7cm]{f7.eps} \\centering \\caption{Sketch of the water model. The dimensions of the OH envelope are also shown (Goicoechea \\& Cernicharo 2002).} \\label{fig_h2o_geometry} \\end{figure}%   The 183.31~GHz emission is produced by the population inversion of the 3$_{13}$ and 2$_{20}$ levels. This is due to the different rates (far from the thermalization densities) at which both levels can be populated. The 3$_{13}$ level is radiatively connected with the 2$_{20}$ and 2$_{02}$  lower energy levels. The A$_{ij}$ coefficient for the 3$_{13}$--2$_{02}$ at $\\sim$138.5~$\\mu$m is $\\sim$10$^4$ times larger than the A$_{ij}$ of the 3$_{13}$--2$_{20}$ transition. The $\\sim$138.5~$\\mu$m line is seen in strong absorption in the ISO/LWS spectra (see Fig.~\\ref{obs_iso_agua}). The 2$_{20}$ level is radiatively connected with the 2$_{11}$ and 1$_{11}$ lower energy levels. The 2$_{20}$--1$_{11}$ at $\\sim$100.9~$\\mu$m produces the strongest absorption of para--H$_2$O in the far--IR (see Fig.~\\ref{obs_iso_agua}). On the other hand, the 3$_{13}$ level is radiatively connected with the 3$_{22}$, 4$_{22}$ and 4$_{04}$ higher energy levels, while the 2$_{20}$ level is connected with the 3$_{31}$ (see the $\\sim$67.1~$\\mu$m line in Fig.~\\ref{obs_iso_agua}) and with the 3$_{13}$ higher energy levels. The different radiative pathways and rates (for $n$$<<$$n_{cr}$) at which the 3$_{13}$ and 2$_{20}$ levels can be populated  produce the inversion. LVG models (see Fig.~\\ref{lvg_mod_183}) show that this mechanism can be efficient to produce 183.31~GHz line emission in regions of relatively low density.   The observed extended emission at 183.31~GHz has a brightness temperature of 10 K in average (see Fig.~\\ref{mapa_183}). Assuming that it arises from the low density regions in which the the far--IR continuum sources are embedded, with T$_k$$=$300-500 K and $n(H_2)$$\\sim$10$^4$~cm$^{-3}$, we derive a value for $N$(p--H$_2$O) of $\\gtrsim$5$\\times$10$^{17}$~cm$^{-2}$. This  value is higher than the one derived from the far-IR water lines (see sect. \\ref{sssct:nonlocal}) possibly because the 183.31 GHz line penetrates deeper into this dusty environment. Toward the main condensation, the bulk of the emission seems to arise from the cold (T$_k\\sim$40 K) and dense gas. Under these conditions, a brightness temperature of the line of $\\sim$100 K would translate into a  p--H$_2$O column density of $\\sim$10$^{19}$~cm$^{-2}$. However, an important contribution could come from the embedded  high temperature and high density core condensations where the column density should also be much larger.  \\subsubsection{Non-Local radiative transfer models} \\label{sssct:nonlocal} To take into account the  radiative coupling between regions with different physical and/or excitation conditions, the radiative transfer has to be treated with nonlocal techniques, more sophisticated than the LVG approximation. We have adapted the radiative transfer used in Sec.~5.1 for CO, to the ground vibrational states of ortho-- and para--H$_2$O respectively. The model includes all the water rotational levels with transitions between 40~$\\mu$m and 183.31~GHz. The level population is computed in statistical equilibrium considering  collisional excitation and  radiative excitation by line and continuum photons.  This is computed consistently assuming that the water molecules and the dust grains are coexistent. The geometry, core$+$shell dimensions, are the same of that considered in the LVG models. The shell was divided in 41 layers. The central dust condensation has been modeled with identical parameters than in section 4.2.1. We have considered 14 rotational levels (E$_{u}$ $<$ 608 K) of ortho-water for model with T$_k$ below 100 K , and up to 30 rotational levels (E$_{up}$ $<$ 1290 K) for the higher temperature models. We have adopted a turbulence velocity of 8 km~s$^{-1}$. The continuum radiation field has been treated as in Gonz\\'alez-Alfonso \\& Cernicharo \\citep{Gon93} by considering a spectral range of 70 km~s$^{-1}$ centered on each water transition. Collisional rates have been taken from Green et al. (1993). Obviously all the water transitions in the continuum core are thermalized to the dust temperature due to the large opacity in the far--IR. The computed line profiles are a result of convolving the brightness temperature with the angular resolution of the LWS detectors ($\\sim$80$''$). The resulting spectral resolution of the synthetic water lines is 1~km~s$^{-1}$, as no convolution with the spectral resolution has been performed. To test the sensitivity of the model to the physical parameters, models were computed with $N$(H$_2$O) and T$_k$ of 1.8$\\times$10$^{16}$, 9$\\times$10$^{16}$ and 4.5$\\times$10$^{17}$~cm~$^{-2}$, and 40, 100, 200, 300 and 500~K respectively, while densities were increased from 5$\\times$10$^{3}$ to  6.4$\\times$10$^{5}$~cm~$^{-3}$ multiplying by 2 in each step. The different nonlocal radiative transfer models for the first two H$_2$O column densities are shown in Figs.~\\ref{mod_1_h2o} and \\ref{mod_2_h2o}. Models for other column densities have been also ran. However, the observed absorption depth is not reproduced for column densities below 10$^{16}$ ~cm~$^{-2}$. In the $N$(H$_2$O)=1.8$\\times$10$^{17}$~cm~$^{-2}$ case too many lines would be in emission, contrary to the observations.   \\begin{figure*}[t]%  \\includegraphics[angle=0, width=12.7cm]{f8.eps} \\centering \\caption{Results from selected nonlocal models of H$_2$O rotational lines in Sgr~B2(M). Results for different gas temperatures (from 40 to 500~K) and H$_2$ densities (from 5.0$\\times$10$^3$ to 1.6$\\times$10$^5$~cm$^{-3}$) are shown. Note that for each model, the water abundance changes accordingly with n(H$_2$) in order to keep constant the water vapor column density (in this case, 1.8$\\times$10$^{16}$~cm$^{-2}$). For the first nine lines, the frequency is given in GHz; for the rest, the wavelength in $\\mu$m is indicated. The velocity scale is shown at the bottom of the 113.5 $\\mu$m line.} \\label{mod_1_h2o} \\end{figure*}%   \\begin{figure*}[h]% \\centering \\includegraphics[angle=0, width=13.0cm]{f9.eps} \\caption{As in Fig.~\\ref{mod_1_h2o} but for a water vapor column density of 9.0$\\times$10$^{16}$~cm$^{-2}$.} \\label{mod_2_h2o} \\end{figure*}%       The main problem to interpret the H$_2$O absorption toward Sgr~B2 arises from the large opacities of the far--IR water lines, $\\sim$10$^3$, and  even $\\sim$10$^4$ for the ortho--H$_2$O 2$_{12}$--1$_{01}$ line at $\\sim$179.5~$\\mu$m. Under these conditions, many weak water lines have to be detected to constrain the physical conditions and the column density. Therefore, we strength that it is  difficult to obtain physical parameters from the observation of few far--IR water rotational lines. In addition, as the radiative excitation dominates the population of the far--IR levels, lines with a weak dependence on the dust excitation should be investigated. Taking into account the 14 water lines detected in the far--IR, the nonlocal results imply that  models are not very sensitive to the temperature and that the only indication about the column density has to be searched in weak far--IR H$_2$O lines or in the 183.31~GHz line (see Sec.~6.1.2). Even so, the far--IR absorption  arises in the low density external layers of gas, while it is very likely that the 183.31~GHz line could have an important contribution from inner and denser regions. Another complication arises from the fact that for models of low $N$(H$_2$O) values and T$_k$$\\lesssim$200~K, it is also difficult to distinguish between different H$_2$ densities (see Fig.~\\ref{mod_1_h2o}). Limits to T$_k$ and $n$(H$_2$) have to be searched in weak  lines below 70~$\\mu$m.  The models for high water column densities (Fig.~\\ref{mod_2_h2o}) predict absorption lines in the LWS range  at $\\sim$56.3, $\\sim$57.6, $\\sim$58.7, $\\sim$78.7, $\\sim$99.5, $\\sim$125.4 and $\\sim$136.5~$\\mu$m. At the spectral resolution and sensitivity of the LWS/FP, none of these lines has been detected. This implies that $N$(H$_2$O)$\\le$4.5$\\cdot$10$^{17}$~cm$^{-2}$ toward the warm envelope of Sgr~B2. The $\\sim$136.5~$\\mu$m line (also predicted by the models with large column densities) is contaminated by the absorption produced by the C$_3$ $Q$(8) rovibrational line, which is also predicted by the models of tri--atomic carbon \\cite{Cer00}. Some of these lines are predicted (even in emission) by the  models with large $N$(H$_2$O). Models with $N$(H$_2$O)$=$1.8$\\times$10$^{16}$~cm$^{-2}$ are consistent with  ISO detections and upper limits to other ISO lines. Only the $\\sim$67.1~$\\mu$m line is weaker than the observations. Therefore, models shown in Fig.~\\ref{mod_1_h2o} give a lower limit to the water vapor column density in the outer and warm (300-500 K) envelope. Taking into account the difficulties implied by the H$_2$O modeling in the far--IR, we found that $(9\\pm3)\\times10^{16}$~cm$^{-2}$ is the best H$_2$O column density to fit the ISO observations (see Figure \\ref{mod_2_h2o}). The CO analysis (Sec.~5), and the studies in the far--IR of OH \\citep{Goi02} and the ammonia lines \\citep{Cec02} and our CO data, also support that the water absorption lines arise in the warm and low density ($\\lesssim$10$^4$~cm$^{-3}$) layer in front of Sgr~B2(M). The H$_2$O column density derived from ISO observations is below the lower limit of $N$(H$_2$O) estimated for the 183.31~GHz line for this component (see Sec. \\ref{sssct:lvg}) as it is likely that an important fraction of the 183.31~GHz emission arises from the inner and denser regions closer to Sgr~B2 main cores or even from them, as pointed out above. This component, with a mean H$_2$ density of $>$10$^{5-6}$~cm$^{-3}$, has been traced by the NH$_3$ non--metastable emission lines, T$_k$$\\simeq$100~K, \\citep{Hut93} and by the emission of HC$_3$N rotational lines, T$_k$$\\simeq$20--40~K, \\citep{Lis91}. The determination of the temperature  from mm emission lines is also complicated in these regions completely obscured to  ISO observations. Therefore, the observed differences in the 3$_{13}$--2$_{02}$ line--intensity and --shape can be a combination of $N$(H$_2$O) and/or T$_k$ variations across the region.  The results presented in this section have been compared to those obtained from another radiative transfer based on a different approach (Asensio Ramos \\& Trujillo Bueno 2003; Asensio Ramos \\& Trujillo Bueno 2006). This code is a generalization to spherical geometry of the very fast iterative Multilevel Gauss-Seidel (MUGA) and Multilevel Successive Overrelaxation (MUSOR) methods developed by Trujillo Bueno \\& Fabiani Bendicho (1995) for the case of a two-level atom and generalized by Fabiani Bendicho, Trujillo Bueno \\& Auer (1997) to multilevel atoms in Cartesian geometries. The code also allows the application of the standard Multilevel Accelerated $\\Lambda$-iteration (MALI) (Olson, Auer \\& Buchler 1986). The angular information for the calculation of the mean intensity is obtained by solving the radiative transfer equation along its characteristic curves (straight trajectories with constant impact parameter) with the aid of the short-characteristics formal solver with parabolic precision (Kunasz \\& Auer 1988) . The statistical equilibrium equations are linearized with the aid of the preconditioning scheme developed by Rybicki \\& Hummer (1991, 1992) with the introduction of an approximate $\\Lambda$ operator that can be efficiently obtained in the framework of the short-characteristics technique. The convergence rate of the MUGA and MUSOR schemes is equivalent to that obtained with the introduction of a nonlinear $\\Lambda$ operator, with the advantage of not being necessary neither to calculate nor to invert such nonlinear operator. Interestingly, the time per iteration is similar to that obtained for the standard $\\Lambda$-iteration or the MALI method. The computing time for the MUGA scheme is reduced by a factor 4 with respect to MALI, while the MUSOR scheme leads to an order of magnitude of improvement in the total computing time with respect to MALI.  The calculations have been performed with the same geometry, the same molecular data (collisional and radiative transitions), and the same physical conditions of the previously described nonlocal code. The emerging line profiles from both codes are very similar, with differences below 2-3 \\%. Both codes predict lines in absorption/emission for the same physical conditions, with identical spectral shapes in the cases where re-emission is found in the line wings. This test of consistency allows us to be very confident in the results presented in Figs.~\\ref{mod_1_h2o} and \\ref{mod_2_h2o} since they were obtained with numerical methods based on completely different approaches.    The H$_2$O (OH) column densities derived from ISO observations in the warm envelope are within an order of magnitude of the shock model predictions for the same low density region \\citep{Flo95}. Nevertheless, $C$--shocks have been invoked to explain the heating of large amounts of warm gas in the GC and also in the warm envelope \\citep{Wil82,Mar97}. In addition, the shock models of \\citep{Flo95} correctly reproduce the observed column densities of $N$--bearing species such as NH$_2$ and NH$_3$ that could be ultimately related with the dust grain chemistry \\citep{G04}.  A qualitative explanation came from the observation of large OH column densities (H$_2$O/OH$\\sim$2-4) in the warm envelope. According to Goicoechea \\& Cernicharo (2002), if a far--UV radiation field illuminates the outer regions of the cloud, water molecules could be photodissociated producing an enhancement of the OH abundance expected only from its formation through H$_3$O$^+$ dissociative recombination. The presence of such an extended far--UV radiation field with $G_0$$\\sim$10$^{3}$-10$^{4}$ is  inferred from the [\\OIII], [\\NIII], [\\NII], [\\CII], and [\\OI] extended line emission \\citep{G04}. Therefore, an important contribution from H$_2$O photodissociation may explain the observed H$_2$O/OH abundance ratio in the external layers of the envelope.\\\\  To investigate in more detail the O--chemistry in the envelope, we have run several photochemistry calculations that include depth--dependent photodissociation, H$_3$O$^+$ dissociative recombination and neutral--neutral reactions. We have used the latest version of the public and available PDR model\\altaffilmark{6} of Le Bourlot et al. \\citep{LBu93}. \\footnotetext[6]{http://aristote.obspm.fr/MIS/} The model does not include oxygen grain surface chemistry, but it is consistent with the low densities and high temperatures found in Sgr~B2 envelope. Following \\citep{G04}, we assume a $G_0$=5$\\times$10$^{3}$ radiation field and a $n_H$=$n$(H)+2$n$(H$_2$)=5$\\times$10$^{3}$ cm$^{-3}$ density. Assuming that the outer gas layers of Sgr~B2 envelope are directly illuminated by such a radiation field, the model solves the UV transfer and the chemistry up to A$_V$=20~mag. As a reference model we solve the thermal balance explicitly for an initial gas temperature of 500~K. The resulting H$_2$O and OH column densities and the H$_2$O/OH ratio are shown in Fig.~\\ref{mods_pdr} as a function of the  visual extinction through the cloud. This $pure$ PDR can only maintain the temperature above $\\sim$100~K in the first 2 magnitudes of the cloud, where the UV radiation efficiently photodissociates H$_2$O to form OH. This results in a low H$_2$O/OH ratio. Inside the cloud, the H$_2$O and OH production is rapidly dominated by  H$_3$O$^+$ dissociative recombination and the H$_2$O/OH ratio tends to a constant value that basically depends on the assumed branching ratio $f_{H_2O}$ (we have taken $f_{H_2O}$=0.25).  \\begin{figure}[t]% \\centering \\includegraphics[angle=-90, width=7.5cm]{f10.eps} \\caption{Selected PDR models in which $G_0$=5$\\times$10$^3$ and $n_H$=5$\\times$10$^3$~cm$^{-3}$. \\textit{Upper panel:} H$_2$O and OH column densities as a function of the visual extinction for an isothermal model with T$_{gas}$=500~K and for a model were the full thermal balance is solved at each depth. \\textit{Lower panel:} H$_2$O/OH column density ratio as a function of the visual extinction for different isothermal models with T$_{gas}$=500, 300, 200 and 50~K and for full thermal balance.} \\label{mods_pdr} \\end{figure}%    At least two different scenarios can produce large column densities of gas at high T$_k$. The first one is the presence of several clumpy PDRs within ISO's beam. In this scenario, far--IR H$_2$O and OH observations will only trace the PDR--clump surfaces. The second one is the presence of shocks. It is very likely that widespread low velocity shocks can locally heat the gas within the envelope and preserve a fraction of the molecular gas with temperatures up to $\\sim$500~K \\citep{Flo95}, but the same is possible in the clumpy PDR scenario. With these temperatures neutral-neutral reactions play an important role in the chemistry. To simulate this situation, we have run several iso--thermal models with temperatures ranging from 500 to 50~K. Some selected results are also shown in Fig.~\\ref{mods_pdr}. From  these models it is clear that both $N$(H$_2$O) and $N$(OH) are clearly enhanced by neutral-neutral reactions (Fig.~\\ref{mods_pdr}, upper panel). If the temperature is high, these models reproduce the observational values much better. Hence, a significant fraction of warm gas seems to be  needed to reproduce the large H$_2$O and OH column densities observed toward Sgr~B2. This conclusion agrees with the large temperatures derived from OH observations in the far--IR \\citep{Goi02}.  The gas temperature determines the contribution of neutral-neutral reactions (1) and (2). As the temperature of the gas increases, reaction (1) contributes to a enhancement of the H$_2$O/OH ratio (Fig.~\\ref{mods_pdr}, lower panel). Note that a moderate increase of the radiation field (to $G_0$=10$^4$) shifts the H$_2$O/OH curve  to larger A$_V$ inside the cloud until photodissociation becomes less important. In the shielded regions of the cloud, the predicted column densities are similar to those with $G_0$=5$\\times$10$^3$.  The H$_2$ column density (or A$_V$) responsible of the warm gas observed from far--IR absorption lines (the location of the species in optical depth) has been traditionally  difficult to establish. A comparison between the observational H$_2$O/OH ratio and photochemistry models shows that the bulk of the H$_2$O/OH absorption must arise from the surface of Sgr~B2 (a maximum A$_V$ of $\\sim$5 to 10 mag), in agreement  with the large  opacities derived from the radiative transfer models. Even assuming an homogeneous surface cloud, most of the water vapor will arise from A$_V$$<$10~mag if T$_k$$\\sim$500~K while $\\gtrsim$50$\\%$ of water can be at A$_V$$>$10~mag if T$_k$$\\sim$300~K. Therefore, an accurate description of the thermal structure of the cloud (with far--IR diagnostics tracing the same gas) will be needed to establish more detailed conclusions. Still, in a inhomogeneous medium, several PDR-like clump surfaces locally heated to T$_k$$\\sim$300--500~K by low velocity shocks could be entirely responsible for the far--IR H$_2$O/OH absorption.   Taking into account the uncertainties implied in the determination of column densities in the envelope, we take $\\chi$(H$_2$O)$\\simeq$(1-2)$\\times$10$^{-5}$ as a lower limit. There is a determination of HD column density toward Sgr B2 of $\\sim$10$^{18}$ cm$^{-2}$, which translates into $N(H_2)\\sim$10$^{24}$~cm$^{-2}$ (Polehampton et al. 2002). However, taking into account the huge absorption by dust at 112 $\\mu$m ($\\tau\\sim$4), it is unlikely that in this HD column density the cold gas, where high column densities are expected, is accounted for.  We derive, see \\citep{Goi02}, that the H$_2$ column density in the warm gas is $\\sim$10$^{22}$ cm$^{-2}$. Besides, the 183.31~GHz observations imply that the water abundance is at least an order of magnitude lower in the core regions.. However, we underline the importance of further detection of weak H$_2$O lines and/or lines with a lesser dependence from the dust emission to refine the models and better constrain the physical parameters of the region (e.g. the temperature), that also determine much of the chemistry. Specially important will be the input of the Herschel/HIFI observations for water lines below $\\sim$2~THz.\\\\   Another difficult  problem is to place the origin of the FUV radiation field revealed by the fine structure line observations in the region \\citep{G04}. In principle, FUV photons could arise from the massive stars near the Sgr~B2(M) core and/or from another stellar population within the envelope itself (not resolved yet by observations). The permeating effect of the radiation field will  be determined by the clumpyness and inhomogeneity of the medium surrounding the stars, and by the energy of the stellar photons. In addition, X--rays observations could complement this scenario of widespread low velocity shocks and UV radiation. Energetic EUV-- and/or  X--ray photons  can penetrate deeper in the neutral cloud than FUV--photons, and thus, could also play a role in the chemistry, as they can also  induce many photo--ionization and photo--dissociation  processes. The correlation found in the Sgr~B complex between the 6.4~keV  Fe$^0$ line and the SiO emission could also indicate  that the X--ray sources  also drive the shocks in the region \\citep{Mar00}.  Gas temperatures in XDRs can easily reach  T$_k$$\\sim$300--500~K because of the more efficient gas heating by X--ray--induced photoelectrons from the gas (and not from  dust grains as in PDRs). Under these conditions, neutral--neutral reactions dominate the chemistry and especially OH reaches large abundances. However, XDR models predict low H$_2$O/OH$<$0.1 abundance ratios \\citep{Mal96}, and thus, they can not be the dominant scenario explaining the far--IR water and OH lines toward Sgr~B2 \\citep{Goi02}.   Sgr~B2 shows diffuse emission in the K$\\alpha$ line of Fe$^0$ at 6.4~keV \\citep{Mur01}. A  dozen of compact X--ray sources have been detected within Sgr~B2 cloud, and they may explain the whole X--ray emission found in the region. Many of the detected X--ray sources are not associated with radio (\\HII $\\,$ regions created by massive stars) nor with IR sources. According to these observations, the intrinsic X--ray luminosity toward Sgr~B2(M) is $L_X$$\\lesssim$10$^{35}$~erg~s$^{-1}$, which translates into a X--ray flux incident on molecular gas of $F_X$$\\lesssim$0.1--0.001 erg~cm$^{-2}$~s$^{-1}$ (assuming $\\sim$0.1--1~pc from the X--ray source). Hence, the X--ray field within Sgr~B2(M) is in the low tail, or even weaker, than those studied by Maloney et al. (1996). As noted by G04, the low [\\OI]63~$\\mu$m/[\\CII]158~$\\mu$m and ([\\OI]63+[\\CII]158)/FIR intensity ratios observed in the region favor a dominant PDR origin for these lines. Both shock \\citep{Dra83} and XDR \\citep{Mal96} models predict larger ratios.       We have carried out far--IR observations of several thermal absorption lines of water vapor toward Sgr~B2(M) and have mapped the 183.31~GHz water line around the main dust condensations of the complex. The main conclusions of this work are the following:  \\begin{enumerate}  \\item The detected water absorption lines are very opaque and arise from the warm envelope around Sgr~B2(M). The observation of the CO $J$=7--6 line and the lack of far--IR CO lines at ISO's sensitivities (3$\\sigma$ limits below $\\sim$2$\\times$10$^{-18}$~W~cm$^{-2}$) imply that the density of such layer is $n$(H$_2$)$\\sim$10$^4$~cm$^{-3}$.  The para--H$_2$O 3$_{13}$--2$_{02}$ line at 183.31~GHz shows $\\sim$40$''$$\\times$40$''$ extended emission around Sgr~B2(M) and $\\sim$40$''$$\\times$20$''$ around Sgr~B2(N). This is the first observation of that line in a GC source and represents further evidence that water vapor is extended in warm molecular clouds.  \\item LVG and nonlocal radiative transfer calculations have been carried out to extract the water vapor abundance and to constrain the physical parameters of the absorbing/emitting regions. Because of the radiative excitation by dust photons, the far--IR water lines are not very sensitive to T$_k$. Taking into account the analysis of the related species OH \\citep{Goi02}, the water absorption must arise from warm gas at similar temperatures, i.e., from 300 to 500~K. For this warm envelope, we found $N$(H$_2$O)=$(9\\pm3)\\times10^{16}$~cm$^{-2}$. An important fraction of the 183.31~GHz emission arises from the inner, denser and colder gas located closer to the main cores. We estimate a water abundance of a few$\\times$10$^{-7}$ in the denser regions.  \\item Photochemistry models show that a component of warm gas, $\\sim$300-500~K, is needed to activate the neutral-neutral reactions and reproduce the large H$_2$O and OH column densities observed in the envelope. We show that OH and H$_2$O far--IR observations toward Sgr~B2 are surface tracers of the cloud (a maximum A$_V$ of 5 to 10 mag). We found $\\chi$(H$_2$O)$\\simeq$(1-2)$\\times$10$^{-5}$ in these regions. Although irradiated by FUV, and possibly more energetic photons, affecting the H$_2$O/OH ratio in the outermost layers, a clumpy structure for the PDR is needed. Alternatively, low velocity  shocks could maintain the gas heating through the envelope.  \\end{enumerate}   Due to the complexity of Sgr~B2 (and also of the GC ISM as a whole) a multiple scenario is needed to explain the  modest angular resolution far--IR observations. The input of chemical models, and higher sensitivity and larger spatial resolution far--IR observations will lead to a better understanding of the GC environment."
    },
    "0601/astro-ph0601100_arXiv.txt": {
        "abstract": "We report the discovery of a new O--type double--lined spectroscopic binary with a short orbital period of 1.4 days. We find the primary component of this binary, \\fo\\,, to have an approximate spectral type O5.5Vz, i.e. a Zero--Age--Main--Sequence star. The secondary appears to be of spectral type O9.5V. We have performed a numerical model fit to the public ASAS photometry, which shows that \\fo\\ is also an eclipsing binary. We find an orbital inclination of $\\sim 80^{\\circ}$. From a simultaneous light-curve and radial velocity solution we find the masses and radii of the two components to be $30\\pm1$ and $16\\pm1$ solar masses and $7.5 \\pm 0.5$ and $5.3 \\pm 0.5$ solar radii. These radii, and hence also the luminosities, are smaller than those of normal O-type stars, but similar to recently born ZAMS O-type stars. The absolute magnitudes derived from our analysis locate \\fo\\, at the same distance as \\ec\\,. From Chandra and XMM X-ray images we also find that there are two close X-ray sources, one coincident with \\fo\\, and another one without optical counterpart. This latter seems to be a highly variable source, presumably due to a pre--main--sequence stellar neighbour of \\fo\\,. ",
        "introduction": "In their survey for OB stars in the field of the Carina Nebula (NGC\\,3372), Forte \\& Orsatti (1981) discovered an early O type star in the darkest region of the nebula. This star ($\\alpha_{2000} = 10^{\\rm h}45^{\\rm m}36^{\\rm s}$; $\\delta_{2000} = -59^{\\circ}48'22''$; V\\,=\\,12.05), number 15 in their list of new OB stars, was assigned a spectral classification O4V on photographic image tube spectrograms. The star has been named \\fo\\, in subsequent literature.  In an infrared study  of the stellar population in the direction of the Carina Nebula, Smith (1987) identifies \\fo\\, as a member of a group of heavily reddened OB stars in the southeast border of the open cluster Trumpler 16, considering that these stars most probably also are members of this cluster. An anomalous reddening law characterized by a value of R=Av/E(B-V)=4.8 is obtained by Smith (1987).  In Fig.~\\ref{mapa} we illustrate the location of \\fo\\ in the Carina Nebula, inside the V--shaped dust lane dividing the brightest part of the nebula, between the open clusters Trumpler 16 and Collinder 228. The H$\\alpha$ image shown in Fig.~\\ref{mapa} was obtained in 1999, May, with the Curtis-Schmidt telescope at the Cerro Tololo Inter--American Observatory (CTIO), Chile.   FO15 was observed as a faint X-ray source in Einstein satellite X-ray images of the Carina Nebula (Chlebowski et al. 1989). Because one of the most often proposed mechanisms for producing X-rays from O-type stars are colliding stellar winds in binary systems, we decided to include FO15 in our ongoing optical spectroscopic observations in search for O-type binaries.  In this paper, we will present our spectroscopic observations of FO15 showing this star to be a double--lined binary system with a short orbital period. After our spectroscopic analysis was essentially complete, FO15 also appeared as an eclipsing binary in the All Sky Automated Survey (ASAS) (cf. Pojma\\'nski, 2003). We have analyzed the ASAS light curve together with our radial velocity study of the binary orbit.  \\begin{figure*} \\centering \\vspace{15cm} \\special{psfile=MF1318f1.ps hoffset=50 voffset=10 hscale=140 vscale=140} \\caption{ H$\\alpha$ Image of the Carina Nebula field surrounding \\fo\\ (indicated by an arrow). North is up and East to the left. The dashed line corresponds to the NW limit of the Spitzer Space Telescope IR observation by Smith et al. 2005 (see text).} \\label{mapa} \\end{figure*} ",
        "conclusions": "\\begin{itemize} \\item We have discovered that the O-type star \\fo\\, immersed in the active star formation site called 'southern pillars' in the Carina nebula, is a short period eclipsing binary.  \\item Both binary components are visible in the spectrum, the secondary component showing considerably weaker lines. \\item We classify the primary spectrum as O5.5Vz, i.e. as an early type Zero-Age-Main-Sequence star. The secondary seems to be of spectral type O9.5V.  \\item Analysis of the ASAS light curve of \\fo\\, fitting a binary model by the Wilson-Devinney method, yields an orbital inclination of $\\sim$80$^\\circ$.  \\item The stellar masses of the components are  $\\sim$ 30 and 16~$\\modot$.  \\item Simultaneous light and radial velocity curve analysis yields components with smaller radii and fainter absolute magnitudes when compared with normal galactic O-type stars.  These values are in agreement with recently born ZAMS O type stars.  \\item An individual determination of total-to-selective extinction ratio (R) for \\fo\\, yields a value of 4.15, which coupled with the values of absolute magnitudes determined from the light curve locate \\fo\\, at a distance of 2.2kpc, coincident with that of $\\eta$ Car.  \\item A Chandra X-ray image shows two close sources at the position of \\fo\\,. Presumably the northern source is a pre-main-sequence star with an ocassional high X-ray state, as this source apparently is not visible in the X-ray image of the same location observed by the XMM satellite. \\end{itemize}"
    },
    "0601/astro-ph0601666_arXiv.txt": {
        "abstract": "There is strong evidence that the observed kHz Quasi Periodic Oscillations (QPOs)  in the X-ray flux of neutron star and black hole sources in LMXRBs are linked to Einstein's General Relativity. Abramowicz\\&Klu\\'zniak (2001) suggested a non-linear resonance model to explain the QPOs origin: here we summarize their idea and the development of a mathematical toy-model which begins to throw light on the nature of Einstein's gravity non-linear oscillations. ",
        "introduction": "QPOs are highly coherent Lorentzian peaks observed in the X-ray power spectra of compact objects in LMXRBs. They cover a wide range of frequencies (from mHz to kHz) and  they show up alone, in pairs or in triplets. In the present review we focus on kHz QPOs which occur in pairs (the so called twin peak QPOs). Many models have been proposed in order to understand kHz QPOs: some of them involve orbital motions (e.g. \\cite{Ste98}, \\cite{Lam03}, but there are also models that are based on accretion disk oscillations (e.g. \\cite{Wag01}, \\cite{Kat01},\\cite{Rez03},\\cite{Lin04}). Here we present semi-analytical results, connected in particular with the Klu\\'zniak\\&Abramowicz model. ",
        "conclusions": "Non-linear parametric resonances occur everywhere in Nature: together with GR they could  explain the mechanism at the basis of kHz QPOs. In this way the mass and the angular momentum (e.g.\\cite{Asc04}) of the central compact object could be precisely measured, but most of all Einstein's strong gravity could be proved: a good motivation to keep on investigating on this puzzling phenomenon."
    },
    "0601/astro-ph0601385_arXiv.txt": {
        "abstract": "{} {We study the environment of active galaxies and compare it with  that of star-forming and normal galaxies.} {From the Fourth Data Release (DR4) of the Sloan Digital Sky Survey (SDSS) we extracted the galaxies in the redshift range $0.05 \\le z \\le 0.095$ and with $M(r) \\le -20.0$ (that is, $M^*$ + 1.45). Emission-line ratios and/or widths were used to separate active galactic nuclei (AGNs) from  star-forming galaxies (SFGs); AGNs were classified as Seyfert-1 (Sy1) and Seyfert-2 (Sy2) galaxies according to emission-line widths. The environmental properties, as defined by a density parameter and the number of companions, are compared for the different types of galaxies, taking the morphological type of the host galaxies into account.} {We find no difference in the large-scale environment of Sy1 and Sy2 galaxies; however, a larger fraction of Sy2 ($\\sim 2\\%$) than Sy1 ($\\sim 1\\%$) is found in systems that are smaller than  $r_{\\rm max} \\le 100$ kpc, mainly in low-density environments  (pairs or triplets). For comparison, this fraction is $\\sim 2\\%$ for star-forming galaxies and $\\sim 1\\%$ for normal galaxies.} {We find no evidence of a relation between large-scale environment properties and activity. If activity and environment are  related, this more likely occurs on small scales (e.g. galaxy interaction,  merging). ",
        "introduction": "The availability of surveys that provide very large databases (e.g., Las Campanas Redshift Survey: Shectman et al.~\\cite{shectman}; 2dF  Galaxy Redshift Survey: Colless et al.~\\cite{colless}; Sloan Digital Sky Survey: York et al.~\\cite{york}) allows for robust statistic analyses of galaxy properties, such as their clustering, luminosity, star-formation rate, and environment. As a consequence, the data from these surveys are leading to significant advances in the study of galaxy formation and  evolution (Kauffmann et al. \\cite{kauffmann}; Benson et al.~\\cite{benson}).  One major topic that can be addressed is the relationship between galaxy environment and activity (SFGs and AGNs). For example, the density in the environment of SFGs is more typical of field galaxies than cluster galaxies, which suggests that star-formation is related more to local processes such as tidal triggering. Moreover, the role of interactions in triggering nuclear starbursts is now widely accepted (e.g. Storchi-Bergmann et al. \\cite{storchi}), and an increment of the star-formation rate is observed for galaxies in close-pair systems (Lambas et al.~\\cite{lambas}; Sorrentino et al. \\cite{sorrentino}; Nikolic et al.~\\cite{nikolic}).  The situation is less clear for AGNs. Stauffer (\\cite{stauffer}) was one of the first to point out that Seyfert galaxies usually occur in groups, and Dahari (\\cite{dahari1};~\\cite{dahari2}) suggested that these galaxies have an excess of companions compared to normal galaxies. This result has been confirmed by several studies (e.g., Laurikainen et al.~\\cite{laurikainen2}; Rafanelli et al. \\cite{rafanelli}) but also contradicted by others (e.g., Fuentes-Williams \\& Stocke,~\\cite{fuentes}; de Robertis et al.~\\cite{derobertis}) in which no detectable excess of companions around Seyfert galaxies is found. Schmitt (\\cite{schmitt}) found that there is no difference in the fraction of galaxies with companions among different activity types if we consider only galaxies with similar morphological types. This result is consistent with those found by Fuentes-Williams \\& Stocke (\\cite{fuentes}) and de Robertis et al. (\\cite{derobertis}) and also with more recent results on clustering of low-luminosity AGNs at higher redshifts (Brown, et al.~\\cite{brown}; Schreier et al.~\\cite{schreier}). Other studies of Seyfert galaxies indicate that Sy2 have a larger number of companions when compared with normal galaxies, while Sy1 do not (Laurikainen \\& Salo~\\cite{laurikainen1}; Dultzin-Hacyan  et al.~\\cite{dultzin2}; Koulouridis et al. \\cite{koulouridis}). As for environment properties,  according to  de Robertis et al. (\\cite{derobertis}), within 50 kpc Sy2 inhabit  richer environments than do Sy1. On larger scales  ($< $ 1 Mpc) Koulouridis et al. (\\cite{koulouridis}) found that  Sy2 reside in less dense large-scale environments than Sy1, but this  is probably  related to the different morphological types of the host galaxies.  According to the so-called unified model (Antonucci~\\cite{antonucci}), different properties observed in AGNs are not due to intrinsic differences: in particular, an AGNs may appear as a  Type 1 or Type 2 depending on the orientation to our line of sight of a  circumnuclear torus of dust and gas. Indeed, the unified model does not imply  that other processes may not occur in the nuclear region, which may even prevail for nearby, low-luminosity AGNs (Seyfert galaxies) or for dust-obscured AGNs. Schmitt (\\cite{schmitt}) suggests that interactions {\\em are} important for triggering activity but that a starburst (SB) may prevail in the earlier phase, hiding any AGNs that might be present. Storchi-Bergmann et al. (\\cite{storchi}) propose an evolutionary link from SFGs to Sy2 galaxies, driven by interaction. They find a correlation between the presence of companions, the inner morphology, and the incidence of recent star-formation, suggesting an evolutionary scenario in which the interaction is responsible for sending gas inward, which both feeds the AGNs and triggers star-formation. The SB then fades with time and the composite Sy2 + SB nucleus evolves into a \"pure\" Seyfert nuclei that may be of Type 1 or 2 in agreement with the unified model. The existence of two different Sy2 population was finally suggested by Tran  (\\cite{tran}) from the absence of detectable polarized broad lines in a fraction of Sy2 and a comparison of their properties with those of Sy1 and Sy2 with  polarized broad lines.  A crucial question to be addressed is therefore whether AGNs and SFGs are found in similar environments, and in particular if there are differences in the environments of Type-1 and Type-2 AGNs.  In this paper we shall use the Fourth Data Release (DR4) of the Sloan Digital Sky Survey (SDSS) to investigate the environment of a complete sample of active galaxies. Spectroscopic data will be used to classify them as SFGs, type-1, or type-2 AGNs, and to compare their environmental properties. In addition, photometric parameters will be used for a morphological classification (early and late-type) of the AGNs host galaxies.  The paper is organized as follows. In  Sects. 2 and 3, we describe the data set and the extraction of the samples. The algorithms used to find the number of neighbours and compute the density are outlined in Sect. 4. The results are presented and discussed in Sect. 5, while the conclusions are in Sect. 6. ",
        "conclusions": "In this paper we have studied the environment of active galaxies (AGNs and SFGs) in the SDSS-DR4 in the redshift range $0.05 \\leq z  \\leq 0.095$ and with $M(r) \\le -20.0$. The presence of emission-lines was used to  separate active galaxies from PGs. The AGNs and SFGs were then separated according to their emission line ratios. AGNs were further separated into Sy1 and Sy2 galaxies  using the width of the emission-line [OIII]$\\lambda$5007 with respect to the Balmer lines (H$_\\alpha$, H$_\\beta$).  The environments of AGNs, SFGs, and PGs have been compared by defining a median density parameter $<\\Sigma>$. The comparison of AGNs with normal galaxies was made through matching the morphological types and the distribution of the galaxy diameters. The main results are:  \\begin{enumerate} \\item The fraction of galaxies classified as an AGNs is 2\\%. This is probably a lower limit due to the severe selection criteria. The fraction of SFGs is 7\\% and the fraction of PGs is 18\\%. UELGs, that is, emission-line galaxies that could  not be classified as SFGs or AGNs, are $\\sim$ 55\\%. \\item  There is no evidence for a difference in the fraction of neighbour galaxies in Sy1 compared to Sy2. The ratio of Sy1 to Sy2 does not change if we take into account all systems, or isolated/pair systems only, in accordance with the unified model. The median surface density is the same for Sy2 and Sy1 ($<\\Sigma> \\sim 2$). The comparison with control samples of PGs, UELGs, and SFGs does not indicate any significative difference in the environment. with the exception of close systems ($r_{\\rm max} \\le$ 100 kpc): we find a higher fraction of Sy2 in close pairs ($\\sim 2$\\%), similar to SFGs, than Sy1 ($\\sim 1$\\%). \\item The analysis of the frequency of systems with close neighbours in Sy1 and Sy2, before and after the morphological separation, shows that their distribution is different only for pairs. If we do not include pairs, the distribution is the same in Sy1 and Sy2 and is consistent  with  that in SFGs. This would imply a higher probability of finding Sy2 than  Sy1 in close pairs. \\end{enumerate}  We conclude that in our sample there is no evidence for a difference in the large-scale  environment between Sy1 and Sy2 galaxies. The only difference is found in close pairs, even if the numbers are low (21 Sy2, 10 Sy1). If these systems are interacting galaxies, the lower fraction of Sy1 may be  due to an increased probability of molecular gas being driven towards the nucleus obscuring the broad line region, as proposed by Dultzin-Hacyan et al. (\\cite{dultzin1}). This result does not seem compatible  with  the simplest formulation of the unified model for Seyfert galaxies, where both type 1 and 2 should be intrinsically alike, the only difference being the result of the orientation of an obscuring torus with respect to the line of sight.  A more detailed analysis of these systems will be the subject of a future paper."
    },
    "0601/astro-ph0601499_arXiv.txt": {
        "abstract": "{ We present a one-equation subgrid scale model that evolves the turbulence energy corresponding to unresolved velocity fluctuations in large eddy simulations. The model is derived in the context of the Germano consistent decomposition of the hydrodynamical equations. The eddy-viscosity closure for the rate of energy transfer from resolved toward subgrid scales is localised by means of a dynamical procedure for the computation of the closure parameter. Therefore, the subgrid scale model applies to arbitrary flow geometry and evolution. For the treatment of microscopic viscous dissipation a semi-statistical approach is used, and the gradient-diffusion hypothesis is adopted for turbulent transport. \\emph{A priori} tests of the localised eddy-viscosity closure and the gradient-diffusion closure are made by analysing data from direct numerical simulations. As an \\emph{a posteriori} testing case, the large eddy simulation of thermonuclear combustion in forced isotropic turbulence is discussed. We intend the formulation of the subgrid scale model in this paper as a basis for more advanced applications in numerical simulations of complex astrophysical phenomena involving turbulence. ",
        "introduction": "In the last decade, the significance of turbulence in various astrophysical phenomena from stellar to cosmological scales has been recognised.  In retrospect, this is hardly surprising, since virtually all the matter in the Universe is fluid, whereas the solid state is encountered as a rare exception. Moreover, both the integral length $L$ and the characteristic velocity $V$ in astrophysical systems are quite large compared to terrestial standards, while the viscosity $\\nu$ is comparable to what is found for liquids or gas on the Earth. For this reason, the dimensionless \\emph{Reynolds number} \\begin{equation} \\mathrm{Re}=LV/\\nu \\end{equation} becomes very large. A typical figure is $\\mathrm{Re}\\sim 10^{14}$ for turbulent stellar convection.  From the computational point of view, the number of degrees of freedom in a fluid dynamical system is given by the relation \\citep[see][]{LanLifVI} \\begin{equation} \\label{eq:landau} N\\sim\\left(L/\\eta_{\\mathrm{K}}\\right)^{3}\\sim\\mathrm{Re}^{9/4}. \\end{equation} The length scale $\\eta_{\\mathrm{K}}$ is called the \\emph{Kolmogorov scale} and specifies the smallest dynamically relevant length scale. Due to the restriction of the CFL time step for compressible flow, the total number of operations required to compute the evolution over one sound crossing time is of the order \\begin{equation} N_{\\mathrm{cr}}\\sim N^{4/3}\\sim\\mathrm{Re}^{3}. \\end{equation} Hence, $N_{\\mathrm{cr}}\\sim 10^{42}$ operations would be required to solve the problem of stellar convection in a direct numerical simulation.  The relevance of the Landau criterion~\\ref{eq:landau} has been questioned recently on the grounds of the \\emph{intermittency} of turbulence \\citep{KritNor05}.  In fact, the number of degrees of freedom $N$ refers to the ensemble of turbulent flow realisations.  At any particular instant of time, however, turbulent dynamics is concentrated in regions of fractal dimension $D<3$.  Topologically, these regions can be either vortex filaments in subsonic flow or shocklets in the case of supersonic turbulence.  The fractal dimension of vortex filaments is $D=1$, whereas $D=2$ for shocklets.  According to the $\\beta$ model \\citep[see][ Sect.~8.5]{Frisch}, the effective number of degrees of freedom is then given by \\begin{equation} N_{\\mathrm{eff}}\\sim \\mathrm{Re}^{3D/(D+1)}. \\end{equation} If an adaptive numerical scheme with maximal efficiency in tracking the intermittent turbulent regions were applied, the total computational cost could be lowered substantially in comparison to a direct numerical simulation on a static grid.  For example, it would appear feasible to treat subsonic turbulence up to a Reynolds number of $10^{10}$ with an anelastic adaptive code on high-end platforms in the near future.  On account of the limited efficiency of adaptive schemes, however, the actual constraint might be lower. Apart from that, gravity, magnetohydrodynamic effects and, possibly, reaction networks increase the work load dramatically in astrophysical applications.  Moreover, we have $N_{\\mathrm{cr}}\\sim N_{\\mathrm{eff}}N^{1/3}\\sim\\mathrm{Re}^{11/4}$ for the particularly interesting case of supersonic turbulence. As a consequence, even with sophisticated adaptive schemes, it remains intractable to resolve completely the turbulent fluid dynamics encountered in astrophysics.  In large eddy simulations (LES), on the other hand, only a limited number of degrees of freedom, which correspond to the largest scales of the system, is treated explicitly.  For the turbulent dynamics on smaller scales, a so-called \\emph{subgrid scale model} is utilised. Among astrophysicists, the most often used subgrid scale (SGS) model is \\emph{numerical dissipation}. This means that all fluctuations on length scales smaller than the resolution $\\Delta$ of the numerical grid are smoothed out and it is assumed that the dynamics on length scales larger than $\\Delta$ are more or less independent of the smaller scales. This point of view is motivated by the second similarity hypothesis of \\citet{Kolmog41}, which holds that the actual mechanism of dissipation is insignificant, provided that there is sufficient scale separation.  In other words, on length scales $l\\ga\\Delta$, it is unimportant whether energy dissipation is caused by the microscopic viscosity $\\nu$ at the length scale $\\eta_{\\mathrm{K}}$ or by numerical effects at the cutoff length $\\Delta$. The notion of numerical dissipation has been exhaustively investigated for the piece-wise parabolic method (PPM) proposed by \\citet{ColWood84}, which is one of the most popular finite-volume schemes applied in astrophysics. At least for statistically stationary isotropic turbulence, the numerical dissipation of the PPM appears to work well as an implicit SGS model \\citep{SyPort00,SchmHille05a}.  For the treatment of transient or inhomogeneous turbulent flow, however, an explicit SGS model becomes mandatory. One of the most prominent examples in contemporary theoretical astrophysics are numerical simulations of turbulent combustion in type Ia supernovae \\citep{HilleNie00}.  In this case, the velocity scale associated with SGS turbulence determines the effective propagation speed of flame fronts \\citep{NieHille95}. Part II of this paper will be dedicated to the problem of type Ia supernova simulations. The exchange of energy between resolved and subgrid scales is expected to become dynamically significant in the case of highly intermittent turbulence, for instance, in collapsing turbulent gas clouds in the interstellar medium \\citep{Larson03,LowKless04}.  In this paper (paper I), we present a general framework for the formulation of SGS models based upon the \\emph{filtering approach} of \\citet{Germano92}.  The mathematical operation of filtering smoothes the flow on length scales smaller than the prescribed numerical resolution $\\Delta$. Consequently, a scale separation is introduced, where the smoothed density, velocity, etc.\\ are identified with the resolved quantities computed in LES.  We have extended this formalism to compressible flows, using the Reynolds stress model proposed by \\citet{Canuto97} as a guideline. Next we discuss the one-equation SGS turbulence energy model \\citep{Schumann75,Sagaut}. For the energy transfer across the numerical cutoff, we introduce a \\emph{localised} eddy-viscosity closure which makes use of the dynamical procedures introduced by \\citet{GerPio91}.  Hence, the SGS model becomes independent of \\emph{a priori} structural assumptions, in particular, whether the resolved flow is homogeneous or stationary. However, it remains a necessary condition that turbulent regions become locally isotropic on length scales comparable to the numerical cutoff length $\\Delta$, because the linear eddy-viscosity closure presumes alignment between the turbulence stress and the rate-of-strain tensors. In the future, multi-parameter closures for the turbulent energy transfer which are not subject to this restriction might be adapted. For the still more complicated non-local transport of kinetic energy on subgrid scales, we use a simple gradient-diffusion closure. In contrast to what has been suggested in the literature \\citep[][ Sect.~10.3]{Pope}, we find that the optimal diffusivity parameter is larger by about one order of magnitude than the SGS viscosity parameter, i.e. the turbulent kinetic Prandtl number is large compared to unity.  Both the localised eddy-viscosity and the statistical gradient-diffusion closure, respectively, were tested by means of data from simulations of forced compressible turbulence. As a case study, we present the LES of turbulent combustion in a periodic box.  Although gravitational effects on subgrid scales, in principle, can be incorporated into the model as well, in paper I we restrict the detailed formulation and application to the case of negligible gravity. However, a simple closure which accounts for unresolved buoyancy effects in simulations of thermonuclear supernovae will be discussed in paper II. ",
        "conclusions": "The localised SGS turbulence energy model offers robustness and flexibility at relatively low computational cost. For this reason, it is particularly suitable for the application in LES of astrophysical fluid dynamics.  The energy transfer from resolved toward subgrid scales is modelled with the standard eddy-viscosity closure, where the closure parameter is computed from local properties of the flow. Hence, there are no \\emph{a priori} assumption about the resolved flow incorporated in the model. Non-local transport is treated with the down-gradient closure, using a constant statistical parameter.  With a turbulent kinetic Prandtl number significantly larger than unity, it is possible to reproduce the magnitude of diffusive flux quite well. The rate of viscous dissipation appears to be particularly challenging.  We found that a semi-statistical approach yields satisfactory results.  The SGS model was implemented in a code for the LES of turbulent thermonuclear combustion in a periodic box using the piece-wise parabolic method (PPM) for the resolved hydrodynamics and the level set method for the flame front propagation. Since PPM produces significant numerical dissipation, we found it favourable to decouple the SGS model form the momentum equation and suppressing inverse energy transfer from unresolved toward resolved scales.  In this kind of application, the SGS turbulence energy serves as a buffer between the resolved kinetic energy and the internal energy and supplies a velocity scale for calculation of the turbulent burning speed.  Furthermore, gravitational and thermal effects can be included in the SGS model, although closures specific to a certain physical system have to be formulated.  An example is presented in paper II, where the application of the SGS model to Rayleigh-Taylor-driven thermonuclear combustion in type Ia supernova is discussed. Adapting the model to other applications, possibly with different numerical techniques, is the goal of on-going research."
    },
    "0601/astro-ph0601450_arXiv.txt": {
        "abstract": "We present upper-limits on the masses of the putative central intermediate-mass black holes in two nearby Galactic globular clusters:~47Tuc (NGC104), the second brightest Galactic globular cluster, and NGC6397, a core-collapse globular cluster and, with a distance of 2.7~kpc, quite possibly the nearest globular cluster, using a technique suggested by T. Maccarone. These mass estimates have been derived from 3$\\sigma$ upper limits on the radio continuum flux at 1.4~GHz, assuming that the putative central black hole accretes the surrounding matter at a rate between 0.1~\\% and 1~\\% of the Bondi accretion rate. For 47Tuc, we find a 3$\\sigma$ upper limit of $2060 - 670 M_\\odot$, depending on the actual accretion rate of the black hole and the distance to 47Tuc. For NGC6397, which is closer to us, we derive a 3$\\sigma$ upper limit of $1290 - 390 M_\\odot$. While estimating mass upper-limits based on radio continuum observations requires making assumptions about the gas density and the accretion rate of the black hole, their derivation does not require complex and time consuming dynamical modeling. Thus, this method offers an independent way of estimating black hole masses in nearby globular clusters. If, generally, central black holes in stellar systems accrete matter faster than 0.1~\\% of the Bondi accretion rate, then these results indicate the absence of black holes in these globular clusters with masses as predicted by the extrapolated $M_\\bullet - \\sigma_{\\rm c}$ relation. ",
        "introduction": "Theory provides us with different ways of producing intermediate-mass black holes (IMBHs) in dense stellar clusters, such as globular clusters. IMBHs are defined here as having a mass larger than that of stellar black holes but smaller than that of the super-massive black holes (SMBH) found in galaxies, so $10^2\\, M_\\odot < M_\\bullet < 10^6 \\, M_\\odot$. One scenario envisages an IMBH to grow from a stellar-mass seed black hole, the remnant of a massive star that exploded as a supernova \\cite{mi02}. This seed black hole rapidly sinks to the cluster center due to dynamical friction, or, in the case of a post-core-collapse cluster, was formed near the cluster center since that is where the most massive stars can be found, and grows by accreting stars, less massive stellar black holes or interstellar gas. In a core-collapse cluster, stars near the core are packed so densely that some may collide, causing a runaway growth of the remnant's mass \\cite{pz02}. Both mechanisms seem to lead naturally to IMBHs with a mass that is about 0.1\\% of the cluster mass. The IMBH can be even more massive (up to $10^4 M_\\odot$) if it originated from multiple seeds. Another suggested way of forming IMBHs considers the radiation drag that is exerted by stars on the interstellar medium \\cite{ka03,no05}. Radiation drag is expected to funnel gas, expelled by SN{\\sc ii} and SN{\\sc i}a explosions, towards the cluster center where it is likely to form a massive central object surrounded by an accretion disk. It is expected that this central object will collapse into an IMBH if its mass exceeds a few 100 $M_\\odot$. In massive galaxies, this process may lead to central black holes that have a mass of about 0.1\\% of the bulge mass \\cite{ka03}. However, in the case of globular clusters, IMBHs formed this way are expected to be much less massive than 0.1\\% of the cluster mass (the ratio of the IMBH mass to the cluster mass is rather of the order of $10^{-5}$). In fact, clusters less massive than $\\sim 5 \\times 10^6 M_\\odot$ are not expected to contain a central black hole, although they can harbor a central object not massive enough to collapse into a black hole \\cite{no05}. Yet another, somewhat more speculative, possibility is that Population {\\sc iii} stars, with masses of a few hundred solar masses, were the seeds of IMBHs \\cite{ma01}. Nuclear energy production in Population {\\sc iii} stars more massive than $\\sim 260 M_\\odot$ cannot counteract gravity and such objects immediately collapse into black holes. Such very massive seed black holes, which could be called IMBHs in their own right, then grow further by accretion. This way of producing IMBHs, unlike the previous ones, does not require the host stellar system to be dense.  Therefore, even upper limits on the masses of central black holes in globular clusters are of great interest, since they constrain the possible formation scenarios for IMBHs and the nature of IMBH progenitors. In this paper, we present upper limits on the IMBH masses of 47Tuc and NGC6397, two of the nearest globular clusters, derived from 20~cm continuum observations, obtained with the Australia Telescope Compact Array (ATCA). In Section \\ref{detect}, we briefly discuss how IMBHs can be detected via their radio continuum emission. In Sect. \\ref{obs}, the observations and the data reduction procedures are presented, followed by a discussion of the results in Sect. \\ref{disc}. We summarize our conclusions in section \\ref{conc}. ",
        "conclusions": "\\label{conc}  Using radio-continuum observations with the ATCA radio telescope array, we estimate the 3$\\sigma$ upper-limits on the masses of the putative central IMBHs in two nearby Galactic globular clusters, 47Tuc and NGC6397, at respectively $2060 - 670 M_\\odot$ and $1290 - 390 M_\\odot$. These mass estimates have been derived from 3$\\sigma$ upper limits on the radio continuum flux. We assume that the putative central black hole accretes the surrounding matter at a rate between 0.1~\\% and 1~\\% of the Bondi accretion rate and take into account the uncertainty on the distance modulus.  Black hole masses estimated using radio-continuum observations by necessity depend on assumptions regarding the density and temperature of the interstellar medium, and on the accretion rate. Observations of a few globular clusters provide ``typical'' values for the characteristics of the interstellar medium. The black hole accretion rate can be constrained by observations of the central black hole of the Milky Way and by theoretical models of accretion disks. Still, radio-continuum estimates of black hole masses probably cannot boast the same degree of accuracy as mass estimates from e.g. dynamical models. However, while dynamical models fitted to radial velocities and proper motions of stars near the center of a globular cluster yield the most accurate mass estimates, their spatial resolution is not high enough to identify the central dark mass as a black hole. If, on the other hand, one detects radio emission coming from accreted matter falling into a steep gravitational well then one can be sure that the central object must be very compact, i.e. a black hole. A cluster of white dwarfs or neutron stars would not be able to generate such a feature.  In a way, both methods are complementary. Still, it seems telling that we obtain a 3$\\sigma$ upper-limit for the mass of the putative central black hole in 47Tuc, a nearby massive globular cluster, that places this object marginally below the extrapolation of the $M_\\bullet - \\sigma_{\\rm c}$ relation of the bright ellipticals and spirals. The data for NGC6397 unfortunately do not constrain the low-mass end of the $M_\\bullet - \\sigma_{\\rm c}$ relation.  Our results hopefully form the basis for further attempts to detect IMBHs in these globular clusters, which, in the case of a detection, would lead to a better understanding of the low-mass tail of the black-hole population and of the accretion rates of central IMBHs in globular clusters."
    },
    "0601/astro-ph0601516_arXiv.txt": {
        "abstract": "{ About 5\\,\\% of upper main sequence stars are permeated by a strong magnetic field, the origin of which is still matter of debate. } { With this work we provide observational material to study how magnetic fields change with the evolution of stars on the main sequence, and to constrain theory explaining the presence of magnetic fields in A and B-type stars. } { Using FORS1 in spectropolarimetric mode at the ESO VLT, we have carried out a survey of magnetic fields in early-type stars belonging to open clusters and associations of various ages. } { We have measured the magnetic field of 235 early-type stars with a typical uncertainty of $\\sim 100$\\,G.  In our sample, 97 stars are Ap or Bp stars. For these targets, the median error bar of our field measurements was $\\sim 80$\\,G. A field has been detected in about 41 of these stars, 37 of which were not previously known as magnetic stars.  For the 138 normal A and B-type stars, the median error bar was 136\\,G, and no field was detected in any of them. } {} ",
        "introduction": "The presence of strong ($\\sim 1$\\,kG) magnetic fields in some of the A- and B-type stars of the upper main sequence has been known for more than 50 years (Babcock \\cite{Bab47}). These fields are almost invariably associated with a suite of other unusual characteristics, which include (a) specific angular momentum of the order of 10\\,\\% or less of the typical value for normal main sequence stars of similar mass; (b) quite anomalous atmospheric chemical composition, which is to first approximation a function of effective temperature (and thus is apparently an atmospheric rather than a global feature); and (c) variation of the spectrum, light and magnetic field with the star's rotational period, that clearly indicate the presence of quite significant abundance inhomogeneities and of a magnetic field not symmetric about the rotation axis.  Although considerable progress has been made in understanding the physical processes at work in these stars, many important problems remain unsolved. Among these are two major questions.  First, although there is strong evidence (e.g., the stability of the observed fields, the lack of symptoms of Sun-like activity, and the lack of any important correlation of observed field strength with rotational angular velocity) that the observed fields are fossil fields, it is not yet clear how these fields evolve during the main sequence phase. Secondly, although it is believed that the basic mechanism leading to both chemical anomalies and to atmospheric inhomogeneities is the competition between gravitational settling, radiative levitation, and various hydrodynamic processes, the interplay of these processes is still very poorly understood.  In this situation, it is helpful to look to observations to guide physical theory. One kind of information about the magnetic Ap and Bp stars (hereafter referred to as magnetic Ap stars) that has been almost entirely lacking is the age of observed stars. Good age information would be very useful for discerning systematic evolutionary changes in field strength, chemical composition, rotation rates, etc.  The general lack of useful age information about magnetic Ap stars occurs because almost all of the bright magnetic Ap stars are field stars. Even with the accurate parallaxes now available from the Hipparcos mission, the uncertainties in luminosity and effective temperature of Ap stars are large enough that placing them in the HR diagram only suffices to determine very roughly their stage of evolution (see Sec 2.1 below).  The obvious way to obtain improved ages for magnetic Ap stars is to observe such stars in open clusters. Until recently such a study has not been possible because cluster Ap stars are mostly fainter than $V \\approx 6$ or 7, beyond the limit of accurate magnetic field measurements with available instruments. This situation has changed due to the development of a new generation of highly efficient spectropolarimeters and observing strategies, and their availability on large telescopes. In particular, the FORS1 spectropolarimeter on one of the ESO 8\\,m VLT telescopes has been shown to be a powerful tool for measuring fields in very faint Ap stars. It has recently been used to detect a field in a star of $V = 12.88$, the faintest magnetic Ap in which a field has ever been detected (Bagnulo et al.\\ \\cite{Bagetal04}).  Another very important development has been the substantial increase in the number of probable magnetic Aps identified in clusters, particularly by the systematic surveys of Maitzen and his collaborators. Furthermore, the availability of very accurate proper motions for a very large number of stars from the Hipparcos mission and the Tycho-2 catalogue has greatly facilitated the correct separation of cluster members from foreground and background stars.  The time is now clearly ripe for studying magnetic Ap stars in clusters to obtain for the first time a reasonably large sample of magnetic Ap stars of known absolute and evolutionary ages. We have started to carry out such a survey, using the FORS1 spectropolarimeter. The first stage of this survey is reported in this paper. Section~\\ref{Sect_Scope} discusses the rationale and scope of this survey, in particular what the advantages are of studying open cluster stars compared to studying field stars, why FORS1 at the VLT is an ideal instrument for this survey, and how individual targets have been selected.  Sections~\\ref{Sect_Basic} -- \\ref{Sect_Balmer_Metal} describe how the magnetic field can be determined from observations of polarized spectra in terms of basic physics, observing strategy, data reduction, and \\Bz\\ determination. These sections contain, for other FORS1 users, a detailed discussion of the optimized techniques we have developed for field measurement, and may be skipped by readers interested mainly in the observational results.  In Sect.~\\ref{Sect_Observations} and \\ref{Sect_Discussion} we present and discuss the observations obtained during this survey. Conclusions are presented in Sect.~\\ref{Sect_Conclusions}. ",
        "conclusions": ""
    },
    "0601/astro-ph0601502_arXiv.txt": {
        "abstract": "\\noindent{\\it Recently, measurements of abundances in extremely metal poor (EMP) stars have brought new constraints on stellar evolution models. Indeed, these stars are believed to have been enriched only by one or a few stars. The abundances observed in EMP stars can therefore be almost directly compared with the yields of very metal poor or metal free stars. This removes the complex filter of chemical evolution models. In an attempt to explain the origin of the abundances observed, we computed pre--supernova evolution models, explosion models and the related nucleosynthesis. In this paper, we start by presenting the pre-SN models of rotating single stars with metallicities ranging from solar metallicity down to almost metal free. We then review key processes in core-collapse and bounce, before we integrate them in a simplistic parameterization for 3D MHD models, which are well underway and allow one to follow the evolution of the magnetic fields during collapse and bounce. Finally, we present explosive nucleosynthesis results including neutrino interactions with matter, which are calculated using the outputs of the explosion models.  The main results of the pre-SN models are the following. First, primary nitrogen is produced in large amount in models with an initial metallicity $Z=10^{-8}$. This large production is due to mixing of carbon and oxygen in the H-burning shell and the main production occurs between core helium and carbon burnings. Second, at the same metallicity of $Z=10^{-8}$ and for models with an initial mass larger than about 60 $M_\\odot$, rotating models may experience heavy mass loss (up to more than half of the initial mass of the star). Several solar masses are lost during the main sequence when the star reaches break-up velocities. The largest amount of mass is lost during the red supergiant stage after rotational mixing and convection have enriched the surface in CNO elements and therefore increased its metal content. The chemical composition of these winds can qualitatively reproduce the abundance patterns observed at the surface of carbon-rich EMP stars. Third, our models predict an upturn of C/O at very low metallicities. Explosive nucleosynthesis including neutrino-matter interactions produce improved abundances for iron group elements, in particular for scandium and zinc. It also opens the way to a new neutrino and proton rich process ($\\nu$p-process) able to contribute to the nucleosynthesis of elements with A$>$64. } ",
        "introduction": "Massive stars ($M\\gtrsim 10 M_\\odot$) play an important role in astrophysics. They are the progenitors of blue supergiants (BSG), red supergiants (RSG), Wolf-Rayet (WR) and luminous blue variable (LBV) stars.  At the end of their life, they explode as type II or Ib,c supernovae (SN) and maybe also as long soft gamma-ray bursts (GRB). Their cores after collapse become neutron stars (NS) or black holes (BH). They are one of the main sites for nucleosynthesis, which takes place during both pre-SN (hydrostatic) burnings as well as during explosive burnings. A weak s-process occurs during H-burning and r-process probably occurs during the explosion. Radioactive isotopes like $^{26}$Al and $^{60}$Fe detected by the INTEGRAL satellite are produced by massive stars. Massive stars even though they are much less numerous than low mass stars contribute significantly to the integrated luminosity of galaxies.  At high redshifts ($z$, or low metallicities $Z$), their importance grows. The first stars formed are thought to be all massive or even very massive (Bromm 2005) and to be the cause of the re-ionisation of the universe. Finally, recent surveys of metal poor halo stars provide many constraints for the early chemical evolution of our Galaxy (Beers et al. 1992, Beers 1999, Israelian et al. 2004, Christlieb et al. 2004, Cayrel et al. 2004, Spite et al. 2005).  The measured abundances of many elements present a very small scatter down to very low metallicities ([Fe/H]$\\sim -4.2$, see Cayrel et al. 2004). This tends to prove that the interstellar medium (ISM) was already well mixed at an early stage. However, r-process elements show a scatter in their abundances contradicting this trend (Ryan...). There is no significant enrichment by pair-instability supernovae (PISN, see Heger and Woosley 2002, Umeda and Nomoto 2002, 2003 for more details). Production of primary nitrogen in massive stars (that died before the formation of the halo metal poor star) is necessary in order to explain the high abundance of nitrogen observed.  About a quarter of the extremely metal poor (EMP) stars show a strong enrichment of carbon (carbon-rich, CEMP), with [C/Fe]$\\sim$2--4, where $[{\\rm A/B}]= {\\rm log_{10}}(X_{\\rm A}/X_{\\rm B})- {\\rm log_{10}}(X_{\\rm A}/X_{\\rm B})_\\odot$. The CEMP stars can probably be divided into two subclasses, the Ba-rich and the Ba-normal stars (see Ryan et al. 2005 and references therein). According to Ryan et al. (2005), the Ba-rich CEMP stars have accreted enriched matter from an AGB stars in a binary system in the course of their life. On the other hand the Ba-normal CEMP stars would have formed from a cloud already enriched in carbon and neutron capture elements. The cloud itself would have been enriched by wind or supernova ejecta from massive stars. The two most metal poor halo stars found up to now, HE0107-5240 ([Fe/H]=-5.3, see Christlieb et al. 2004) and HE1327-2326 ([Fe/H]=-5.5, see Frebel et al. 2005 and Aoki et al 2005) belong to this last group, the Ba-normal CEMP stars. These two stars both present strong but different enrichment in nitrogen and oxygen.  In the light of these recent observational results and following the exploratory work of Meynet et al (2005), we explore the impact that rotation can have in massive very low metallicity stars. In Sect. 2, we present the pre-SN models, their evolution and stellar yields for different initial masses, rotation velocities and metallicities and compare them with EMP stars. In Sect. 3, we describe the recent update in the explosion models and the corresponding nucleosynthesis in 1D and the latest development of a new 3D MHD model. In Sect. 4 we give our conclusions. ",
        "conclusions": ""
    },
    "0601/astro-ph0601444_arXiv.txt": {
        "abstract": "{The OGLE-II event sc5\\_2859 was previously identified as the third longest microlensing event ever observed. Additional photometric observations from the EROS (Exp\\'erience de Recherche d'Objets Sombres) survey and spectroscopic observations of the candidate star are used to test the microlensing hypothesis.The combined OGLE and EROS data provide a high quality coverage of the light curve. The colour of the sc5\\_2859 event is seen to change with time. A spectrum taken in 2003 exhibits a strong $H_{\\alpha}$ emission line. The additionnal data show that the OGLE-II sc5\\_2859 event is actually a classical nova outburst. ",
        "introduction": "Gravitational microlensing is now a well established tool for astrophysics. Of the more than 2000 microlensing events discovered to date, the great majority are in the direction of the Galactic bulge.  The microlensing optical depth and the event timescale distribution obtained from these surveys provide unique tools for the study of Galactic structure and the mass functions of Galactic populations. Moreover, microlensing databases have proven to be an extremely rich source for the study of various classes of variable stars.  In particular, they hold out the potential to obtain nova light curves  with higher cadence and better precision than is usually the case.  The confusion between falling nova light curves and microlensing events is a potential problem for microlensing studies.  The event sc5\\_2859, discussed in this paper, stands both as a concrete example of this confusion and as a warning about its impact on calculations of the optical depth.  Nov\\ae\\ have much larger outburst amplitudes than most microlensing events. But they are also easily distinguished from microlensing events from their very different light curves on the rising side of the peak.  Hence, if these analyses excluded events with data points only over the descending part of the light curve, this potential confusion would have no practical effect on optical depth measurements. This is not the case, however, for two published optical-depth measurements \\citep{alcock97,alcock00}, which included, respectively, one and two decline-only events with long Einstein timescales $t_{\\rm E}$.  Specifically, event 124-A ($t_{\\rm E}=76\\,$days), for the first paper and events 95-BLG-d11 ($t_{\\rm E}=62\\,$days) and 96-BLG-d14 ($t_{\\rm E}=47\\,$days) for the second.  Here $t_{\\rm E}$ is given by \\begin{equation} t_{\\rm E} = {\\theta_{\\rm E}\\over \\mu_{\\rm rel}},\\qquad \\theta_{\\rm E} = \\sqrt{\\kappa M\\pi_{\\rel}}, \\label{eqn:tedef} \\end{equation} where $\\theta_{\\rm E}$ is the angular Einstein radius, $\\pi_{\\rm rel}$ and $\\mu_{\\rm rel}$ are the lens-source relative parallax and proper motion, $M$ is the lens mass, and $\\kappa\\equiv 4G/c^2\\rm AU\\sim 8.1\\,{\\rm mas}$/$M_\\odot$. Although these events do not represent the bulk of the optical depth, they contribute significantly to it, since each event contributes $\\propto t_{\\rm E}$.     Three very long events with timescales greater than one year have been discovered in the OGLE-II and MACHO data sets: OGLE-1999-BUL-32/MACHO-99-BLG-22 \\citep{Mao2002,Bennett2002,Agol2002}, OGLE-1999-BUL-19 \\citep{Smith2002} and BUL\\_SC5 244353/sc5\\_2859 \\citep{Smith2003}. The last of these, sc5\\_2859, was initially found by the OGLE-II collaboration and later identified by \\citet{Smith2003} as a long duration parallax event with one of the longest time scales ever observed, $t_\\e = 547.6^{+22.6}_{-2.8}$ days, and as one of the first examples of disc-disc microlensing.  In this paper we present the light curve of the event sc5\\_2859 including additional points from the EROS~2 and OGLE projects databases. The new light curve and complementary spectroscopic observations of the candidate reveal that the event is most probably a nova outburst. ",
        "conclusions": "\\label{sec:discussion}  The spectrum of the stellar object associated the OGLE-II event SC5\\_2859 resembles a reddened cataclysmic variable (CV) star. These systems consist of a Roche-lobe filling main sequence star and a white dwarf accreting material through an accretion disc surrounding the compact object.  The optical light is generally dominated by strong emission lines of the Balmer series of H, \\ion{He}{I}, \\ion{He}{II}, and the \\ion{C}{III}/\\ion{N}{III} Bowen blend at 4640-4640 \\AA\\ from the bright accretion disc.  A comprehensive review of the properties of CVs is given by \\cite{Warner1995}. In comparing this event with other CVs, the great strength of H$\\alpha$ emission and its breadth are both consistent with a system viewed at relatively high orbital inclination, close to 90\\degr.  The presence of a CV associated with the outburst as measured by the OGLE and EROS experiments offers a natural explanation for the nature of the observed event. CVs are the progenitors of both dwarf novae and classical novae outbursts. The spectrum of this object is certainly consistent with other CVs that undergo dwarf nova outbursts such as U Gem and SS Cyg.  The quiescent spectrum of these objects is dominated by strong Balmer emission lines accompanied by the lack of strong \\ion{He}{I} and \\ion{He}{II} emission lines relative to H$\\beta$.  In addition, the \\ion{C}{III}/\\ion{N}{III} blend is nearly always weak or absent in the quiescent spectra of dwarf novae. However, dwarf nova outbursts are relatively fast events with amplitudes of $\\simeq$2-4 magnitudes and durations of less than $\\sim$20 days. In addition, they show a recurrence time scale $\\sim$100 days or less. The light curve of OGLE SC5\\_2859 is inconsistent with the dwarf nova interpretation since the duration of the outburst is considerably longer and it is not observed to recur over the $\\sim$1300 days of monitoring.  The most likely interpretation of the sc5\\_2859 event is that it is a classical nova outburst.  Classical novae occur in CV binary systems as a result of a thermonuclear runaway on the surface of the white dwarf primary star. The time evolution of classical nov\\ae\\ outbursts can be separated into 4 phases. These phases are 1) the rise to maximum visual light, 2) a period of constant bolometric luminosity, 3) a period of cooling of the white dwarf and 4) the re-establishment of accretion from the secondary. The evolutionary track of classical nov\\ae\\ remnants is shown in Figure 5.1 of \\citet{Classical Novae}. This figure uses nova DQ Her 1934, for which the transition between phases 2 and 3 occurs at $M_{V} \\simeq 7.$ The $V$ magnitude changes by more than 10 units during phase 2. The  track of the sc5\\_2859 object in the local $(V-I,I)$ CMD is displayed in Figure \\ref{fig:diaghr}. The 1997 data may  be identified with the end of the ``constant bolometric luminosity'' period and the 1998-2000-2001 data with the \"white dwarf cooling phase\". The white dwarf cooling phase thus occurs after the turning point in the colour evolution around $JD=2450860$ days.   \\begin{figure} \\begin{center} \\mbox{\\epsfig{file=fig8.eps,width=9.5cm}} \\caption[] { $I_{\\rm OGLE}$ flux versus $V_{\\rm OGLE}$ flux during the ``white dwarf cooling phase''. The $I_{\\rm OGLE}$ flux is a linear function of the $V_{\\rm OGLE}$ flux because the $V-I$ colour of the dwarf is almost independent of temperature in the high temperature limit. } \\label{fig:vvscoul} \\end{center} \\end{figure}   We now find the distance of the nova from the evolution of colour versus time. For this, we use only OGLE colours, which are more reliable, since they are not affected by the H${\\alpha}$ line. Figure \\ref{fig:vvscoul} shows a plot of $I_{\\rm OGLE}$ flux versus $V_{\\rm OGLE}$ flux during the ``white dwarf cooling phase''. This plot is very well fitted by a straight line with slope  $s=4.2 \\pm 0.4.$ A possible interpretation is as follows: the flux comes from the blend of two objects: one with $I$ and $V$ fluxes that are constant in time and another with $I$ and $V$ fluxes decreasing with time, but with constant $V-I$ colour. Translating the value of the slope $s$ in term of $V-I$ colour gives: ${(V-I)}^{\\mbox{WD}} = 1.56.$  A high-temperature white dwarf cooling is expected to have an almost constant $V-I$ colour. This is because the spectrum of a high-temperature white dwarf cooling changes mainly in the UV. This qualitative argument is supported  by the calculations of \\citet{Bergeron95}. They find that the $V-I$ colour of a hydrogen-atmosphere white dwarf with $\\log g = 8$ and temperature $> 30000$ K should be within 0.05 mags of the asymptotic value ${(V-I)}^{H} = -0.36$ (see also their Fig.~6). Using the theoretical value   ${(V-I)}^{H} = -0.36,$ one finds for sc5\\_2859, $E(V-I)= 1.56 - (-0.36)= 1.92.$ If a helium atmosphere model is used instead, one gets $E(V-I)=1.82.$ These correspond respectively to 68.5\\% and 65\\% of the dust column to the Galactic centre.     The projected height of the field is 4.0 deg $\\times$ 8 kpc = 560 pc, while the scale height of dust is 130 pc. Hence, for the nova to lie in front of 68.5\\% the dust, a hydrogen atmosphere white dwarf would have to be at a height of $z=-130\\ {\\rm pc}\\,\\ln 0.315=150.2 {\\rm pc}$ and so at a distance of $D^{WD}= 8\\ \\mbox{kpc}(150.2\\ \\mbox{pc}/560\\ \\mbox{pc})= 2140\\ \\mbox{pc}.$ A similar calculation gives a distance of 1950 pc for a helium atmosphere dwarf. Taking the average value as the distance, we find \\begin{equation} D^{WD}= 2045\\ \\pm 95\\ \\pm 500 \\mbox{pc,} \\end{equation} where the first error represents the uncertainty for the H/He ambiguity of the WD composition and the second error represents our rough estimate in the uncertainty in estimating the distance from the relative extinction.   { Around $JD=2451704$, the apparent magnitude of the sc5\\_2859 object is $V = 21.5 \\pm 0.05.$ Recalling that the exinction to the nova is $A_V = R_{VI} E(V-I) = 4$, the absolute $V$ magnitude of the source (assuming it is a hydrogen atmosphere WD) is then \\begin{equation} M_{V} \\mbox{(JD=2451704)} \\simeq 5.8. \\end{equation} This is compatible with the minimum $M_{V}$ values given in Table 6 of \\citet{Downes2000}. However, the $M_{V}$ value of sc\\_2859 might be somewhat lower than 5.8, since it is very likely that the source is blended. } Note that the OGLE collaboration found that the source of sc5\\_2859 had a very high apparent proper motion which they interpreted as the effect of a blend \\citep{Sumi2004b}.  Since the earliest EROS $V_{E}$ point has $V \\simeq 15$, it is likely that most of the early decline part of the light curve (interpreted as ``phase 2'') was missed. The speed class of the nova can only be guessed at by means of qualitative arguments. The ``transition part'' of the V light curve is smooth and steadily declining. This favours a fast or very fast nova. Other features give only weak constraints. The slope of the $V$ light curve around  $JD=2450600$ is \\begin{equation} \\frac{dV}{dt} \\simeq 0.02\\ \\mbox{mag/day}, \\end{equation} but this is only a lower limit to the early-decline slope. A similar constraint comes from the date of the first data point of the $I_{E}$ light curve, which is observed before the outburst. These arguments favour  a fast nova outburst.  { Another useful check of the nova hypothesis is the calculation of the H${\\alpha}$ line flux. The ratio of the R band continuum flux to the I band continuum flux can be estimated from the spectrum. Transforming to magnitudes using table 2.1 of \\citet{BinneyMerrifield}, one finds that \\begin{equation} R^{cont} - I = 1.53 \\end{equation} around $JD=2452829,$ so that $R \\simeq 20.$ The $R$ absorption is $A_{R} \\simeq 3.$ The $R^{cont}$ value corresponds to a continuum flux in the R band of $2.3\\ 10^{32}$ erg/s if the sc5\\_2859 object is located at 2.05 kpc. The ratio of the H${\\alpha}$ line flux to the continuum R band flux is roughly 0.5, so the H${\\alpha}$ line flux is roughly $10^{32}$ erg/s. H${\\alpha}$ line fluxes of respectively 1,\\ 3 and 4~$10^{32}$ erg/s are expected for very fast, fast and moderately fast, 6.5 year old classical nova outbursts, according to \\citet{Downes2001}. The observed H${\\alpha}$ line flux is thus in agreement with the hypothesis of a fast or very fast nova. }  The paragraphs above summarize the arguments in favour of the classical nova hypothesis. It is also possible in principle that the sc5\\_2859 event is related to the subclass of recurrent novae such as T CrB, RS Oph, V3890 Sgr, and V745 Sco. The lack of obvious spectral features arising from a late-type giant star in the red weakens this association, however, in spite of the similarity of the light curve to some other recurrent novae. Nevertheless, a search for previous outbursts in archival plate material might prove worthwhile.   \\bigskip"
    },
    "0601/astro-ph0601358_arXiv.txt": {
        "abstract": "We carried out global three-dimensional resistive magnetohydrodynamic simulations of galactic gaseous disks to investigate how the galactic magnetic fields are amplified and maintained. We adopt a steady axisymmetric gravitational potential given by \\citet*{Miy75} and \\cite*{Miy80}. As the initial condition, we assume a warm ($T\\sim10^{5}\\,\\, \\mathrm{K}$) rotating gas torus centered at $\\varpi=10\\,\\, \\mathrm{kpc}$ threaded by weak azimuthal magnetic fields. Numerical results indicate that in differentially rotating galactic gaseous disks, magnetic fields are amplified due to magneto-rotational instability and magnetic turbulence develops. After the amplification of magnetic energy saturates, the disk stays in a quasi-steady state. The mean azimuthal magnetic field increases with time and shows reversals with period of 1Gyr (2Gyr for a full cycle). The amplitude of $B_{\\varphi}$ near the equatorial plane is $B_{\\varphi}\\sim1.5\\mu\\,\\, \\mathrm{G}$ at $\\varpi=5\\,\\, \\mathrm{kpc}$. The magnetic fields show large fluctuations whose standard deviation is comparable to the mean field. The mean azimuthal magnetic field in the disk corona has direction opposite to the mean magnetic field inside the disk. The mass accretion rate driven by the Maxwell stress is $\\sim10^{-3}M_{\\odot}/\\,\\, \\mathrm{yr}$ at $\\varpi=2.5\\,\\, \\mathrm{kpc}$ when the mass of the initial torus is $\\sim5\\times10^{8}M_{\\odot}$. When we adopt an absorbing boundary condition at $r=0.8\\,\\, \\mathrm{kpc}$, the rotation curve obtained by numerical simulations almost coincides with the rotation curve of the stars and the dark matter. Thus even when magnetic fields are not negligible for gas dynamics of a spiral galaxy, galactic gravitational potential can be derived from observation of rotation curves using gas component of the disk. ",
        "introduction": "Spiral galaxies have magnetic fields. In our Galaxy, the average magnetic field strength is a few $\\mu$G. \\citep*[e.g.,][]{Sof86,Bec96}. These magnetic fields have been explained by the dynamo action operating in galactic gas disks \\citep*[e.g.,][]{Par71}. In conventional theory of galactic dynamos, the following induction equation is solved; \\begin{equation} \\frac{\\partial\\bar{\\mbox{\\boldmath $B$}}}{\\partial t} = \\nabla\\times(\\mbox{\\boldmath $v$}\\times\\bar{\\mbox{\\boldmath $B$}}) +\\eta \\nabla^{2} \\bar{\\mbox{\\boldmath $B$}} + \\nabla\\times(\\alpha \\bar{\\mbox{\\boldmath$B$}}), \\label{eqn1} \\end{equation} where $\\eta$ denotes the turbulent magnetic diffusivity, \\mbox{\\boldmath $v$} is the mean velocity, and \\mbox{\\boldmath $\\bar{B}$} is the mean magnetic field. The last term on the right hand side of equation (\\ref{eqn1}) is the dynamo term which represents the re-generation of mean magnetic fields ($\\alpha$ effect). When velocity fields \\mbox{\\boldmath $v$} are given, equation (\\ref{eqn1}) is linear with \\mbox{\\boldmath $\\bar{B}$}. It can be solved with appropriate boundary conditions. This approach is referred to as the {\\it kinematic dynamo}. The effects of nonlinearities such as the quenching of the $\\alpha$-effect for strong magnetic fields and the loss of magnetic flux due to the magnetic buoyancy have also been studied in the framework of the kinematic dynamo theory \\citep*[e.g.,][]{Sch89,Bra89,Bra92}.  In 1990s, it became clear that nonlinear dynamical processes are essential for the evolution of magnetic fields in differentially rotating disks. \\citet*{Bal91} showed that in the presence of weak magnetic fields, differentially rotating gas disks subject to the robust magneto-rotational instability (MRI). This instability is caused by the back-reaction of magnetic fields to the fluid motion. Thus, in order to study the evolution of magnetic fields in differentially rotating disks, we should solve the momentum equation including Lorentz force simultaneously with the induction equation. This approach is called {\\it dynamical dynamo}. The developments of computational magnetohydrodynamics (MHD) enabled us to study this dynamical dynamo process by direct three-dimensional (3D) global simulations.  In the context of accretion disks, the nonlinear growth of MRI was first studied by local 3D MHD simulations \\citep*[e.g.,][]{Haw95,Mat95,Bra95}. They showed that MRI drives magnetic turbulence in accretion disks and amplifies magnetic fields. Subsequently, global 3D MHD simulations of initially weakly magnetized rotating torus were carried out \\citep*[e.g.,][]{Mat99,Haw00,Mac00}. They showed that the magnetic field amplification saturates when $\\beta=P_{gas}/P_{mag} \\sim 10$. The magnetic fields maintained in the quasi-steady state are mainly toroidal but have turbulent poloidal components. They also showed that in the innermost region of black hole accretion disks where radial infall of disk gas becomes important, spiral magnetic fields which accompany radial field reversals are created \\citep*[e.g.,][]{Mac03}.  In numerical studies of accretion disks, the gravitational field is assumed to be that of the central point mass. It is straightforward to extend this gravitational potential to more general potentials such as those of spiral galaxies. Since the galactic gaseous disk is also a differentially rotating disk, it subjects to the MRI. The growth of MRI and generation of MHD turbulence in galactic HI disks were discussed by \\citet*{Sel99}. \\citet*{Kim03} carried out three-dimensional local MHD simulations of galactic gaseous disks and showed that MRI drives interstellar turbulence and helps forming the giant molecular clouds. The coupling of the thermal instability and MRI in the interstellar space has been studied by \\citet*{Pio04}.  \\citet*{Dzi03} reported the results of global three-dimensional MHD simulations of the galactic gaseous disks including the dynamo $\\alpha$ term. More recently, \\citet*{Dzi04} carried out global three-dimensional MHD simulations of galactic gaseous disks without assuming the dynamo $\\alpha$ term starting from vertical magnetic fields and showed that mean magnetic fields and interstellar turbulence are generated. However, their simulation area was limited to 5kpc in the radial direction and 1kpc in the vertical direction. \\citet*{Kit04} carried out linear but global analysis for MRI in a disk geometry and showed that MRI can amplify very weak seed magnetic fields in proto-galaxies.  In this paper, we report the results of global 3D MHD simulations of galactic gaseous disks by using the axisymmetric gravitational potential created by stars and the dark matter. We solve the nonlinear MHD equations without introducing the phenomenological dynamo $\\alpha$ parameter. The initial magnetic field is assumed to be toroidal. This initial condition enables us to start the simulation without much disturbing the interface between the disk and the disk corona. When we start from vertical magnetic fields, efficient angular momentum loss near the surface of the initial torus drives an avalanche-like surface accretion, which deforms magnetic field lines into an hourglass shape and produces magnetocentrifugally driven outflows \\citep*[e.g.,][]{Mat96}. For simplicity we neglect the multi-phase nature of galactic gaseous disk, non-axisymmetric spiral gravitational potential, self gravity and supernova explosions. These processes will be incorporated in subsequent papers.  In section 2, we present numerical methods. Numerical results are given in section 3. Section 4 is devoted for summary and discussion. ",
        "conclusions": "We carried out global three-dimensional magnetohydrodynamic (MHD) simulations of galactic gaseous disks to investigate how the galactic magnetic fields are amplified and maintained. We showed the dependence of numerical results on physical conditions in the innermost region, gravitational potential, and strength of initial magnetic fields.  In all models, the magnetorotational instability (MRI) grows in the disk and the magnetic energy is amplified and maintained for more than $5 \\,\\, \\mathrm{Gyrs}$. The saturation level of magnetic fields, however, is smaller than that expected from observations ($B \\sim$ several $\\mu G$) at $\\varpi = 10 \\,\\, \\mathrm{kpc}$. The multiphase nature of interstellar space, supernova explosions, cosmic rays  and/or non-axisymmetric spiral potential may be important to further amplify magnetic fields.  Numerical results indicate that mean azimuthal field inside the galactic disk is amplified and that its direction changes quasi-periodically. The amplitude of the oscillation of the azimuthal field increases with time. The interval of reversal of equatorial azimuthal field is $\\sim 1 \\,\\, \\mathrm{Gyr}$. This timescale is comparable to the timescale of the buoyant rise of azimuthal magnetic flux from the disk to the corona. The amplitude of the fluctuating field is comparable to the mean magnetic field. The strength of the mean field only weakly depends on the initial strength of magnetic fields. In the coronal region, the mean azimuthal field has direction opposite to the mean field of the disk.  Mean magnetic fields obtained from numerical simulations also show reversals in the radial direction near the equatorial plane around $\\varpi = 5 \\,\\, \\mathrm{kpc}$ at $t = 3.8 \\,\\, \\mathrm{Gyr}$. The radius of the field reversal changes with time because the spiral magnetic channels which produce the reversal of azimuthal fields in the radial direction move inward as the disk gas accretes toward the galactic center.  In the Galactic plane, magnetic field reversals are observed by Rotation Measure (RMs) of pulsars \\citep*[e.g.,][]{Han02}. The reversals of mean magnetic fields are observed in our Galaxy between the local Orion arm and the inner Sagittarius arm \\citep*{Sim79}. More recent analysis imply axisymmetric field with two reversals (\\citealt*{Ran89,Ran94}; see also \\citealt*{Bec96}). Simulation results of model III (Figure \\ref{fig:f13}) show reversals of mean magnetic fields near the equatorial plane around $\\varpi = 5 \\,\\, \\mathrm{kpc}$. Note that RMs measure the mean magnetic field in some specific direction from the Earth. When we draw lines from a point in Figure \\ref{fig:f13}, we can see field reversals in several directions. Since the magnetic fields inside the galactic disk is turbulent, more and more field reversals will be observed as angular resolution of RMs increase.  In conventional theories of $\\alpha \\omega$-dynamos, various kinds of spatial and temporal oscillations appear depending on the dynamo number $D = C_{\\alpha} C_{\\Omega}$ where $C_\\Omega = H^{2} \\Omega/ \\eta$ ($H$ is the thickness of the disk, $\\Omega$ is the rotation angular speed and $\\eta$ is the magnetic diffusivity) and $C_\\alpha = H \\alpha/\\eta$ \\citep*[e.g.,][]{Ste88}. In our global MHD simulations, we do not need to introduce the dynamo $\\alpha$ parameter. The amplification of magnetic fields due to MRI and the reproduction of toroidal fields through gas motions are self-consistently incorporated in our simulations. We introduced anomalous resistivity $\\eta$ to handle magnetic reconnection which takes place in local region where current sheets are formed, but the turbulent diffusivity is not explicitly assumed. We showed by direct 3D simulations that mean toroidal magnetic fields show reversals both in space and time.  \\citet*{Tou92} constructed phenomenological models of disk dynamos by taking into account the growth of magnetic fields due to MRI, escape of magnetic flux by the Parker instability, and the dissipation of magnetic fields by magnetic reconnection. They showed that the magnetic fields in the disk oscillate quasi-periodically with period comparable to the time scale of the buoyant escape of the magnetic flux. This timescale is consistent with the timescale of field reversals obtained from our MHD simulations.  Let us discuss the similarities and dissimilarities of the magnetic activity of galactic gas disks produced by our numerical simulations and those in the Sun. Both objects are differentially rotating plasma with high conductivity in which magnetic fields are almost frozen to the rotating plasma. Thus the magnetic fields can be deformed and amplified by the plasma motion.  The Sun is a slow rotator in which the centrifugal force is much smaller than the gravity, thus its atmosphere is confined by the gravity. The convective motions and/or the differential rotation in the convection region may be the source of magnetic field amplification in the Sun. However, since the magnetic flux tube in the convection region buoyantly rises with a timescale of the order of a month \\citep[e. g., ][]{Par75}, much shorter than the solar magnetic cycle ($\\sim 22 \\,\\, \\mathrm{year}$), solar dynamo activity is considered to be operating in the interface between the convection zone and the radiative zone, where magnetic buoyancy is small and differential rotation exists \\citep*[e. g., ][]{Spi80,Sch83}. On the other hand, galactic gas disks are supported by the rotation and subject to differential rotation, which induces the growth of MRI in the whole disk. In galactic gas disks, since the gas disk is convectively stable, the buoyant escape of the magnetic flux takes place with timescale $t_{buoy} \\sim 10 H /v_{A} \\sim 10 (1/\\Omega)(\\beta/2)^{1/2}$. This timescale is longer than the growth time of MRI $t_{MRI} \\sim 1/\\Omega$. Thus the magnetic fields can be amplified in the disk before escaping to the corona.  Numerical results indicate that the total azimuthal magnetic flux is nearly conserved. When turbulent diffusivity exists, the azimuthal magnetic flux is not necessarily conserved. In our simulations, however, since magnetic diffusion only exists in the local current sheet, the magnetic flux is almost conserved. We found that the toroidal magnetic flux in the disk buoyantly rises into the corona. This result indicates that the buoyant escape of the magnetic flux from the disk enables the amplification of mean magnetic fields in the galactic gas disk. When the magnetic diffusion is not negligible, the total azimuthal magnetic flux is not conserved but the magnetic helicity is conserved. \\citet*{Bla03} pointed out the importance of helicity conservation and the role of coronal mass ejections, which help sustaining the solar dynamo cycle. In our simulations, buoyantly rising magnetic loops carry helicity as well as the magnetic flux threading the disk. This escape of magnetic flux and helicity enable the disk to amplify the azimuthal magnetic flux in the equatorial region.   We showed that when we remove the central absorber at $r = 0.8 \\,\\, \\mathrm{kpc}$, formation of dense gas bulge forces the rotation curve to deviate from that expected from the gravitational potential (model II). This indicates that in the central region of the galaxy, the gas absorption or the conversion of the accumulated gas to stars are essential to explain the observed rotation curve. The mass accretion rate to the central region is $\\sim 10^{-3} M_{\\odot}/ \\,\\, \\mathrm{yr}$ when we adopt the absorbing boundary condition. This accretion rate depends on the initial density of the torus, which can be formed by the infall of intergalactic matter or by the supernova explosions after bursts of star formation in the certain radius of the disk. When the galaxy was more gas rich in the early stage of its evolution, the accretion rate should be much higher.   We also showed that the gas rotation curve approximately coincides with the rotation curve for stars and dark matter even when magnetic fields are dynamically important. This justifies us to use the gas rotation curve to estimate the distribution of the dark matter."
    },
    "0601/astro-ph0601391_arXiv.txt": {
        "abstract": "Gravitational lensing assists in the detection of quasar hosts by amplifying and distorting the host light away from the unresolved quasar core images.  We present the results of {\\it HST} observations of 30 quasar hosts at redshifts $1 < z < 4.5$.  The hosts are small in size ($r_e \\lesssim 6$ kpc), and span a range of morphologies consistent with early-types (though smaller in mass) to disky/late-type.  The ratio of the black hole mass ($M_{BH}$, from the virial technique) to the bulge mass ($M_{bulge}$, from the stellar luminosity) at $1 \\lesssim z \\lesssim 1.7$ is broadly consistent with the local value; while $M_{BH}/M_{bulge}$ at $z \\gtrsim 1.7$ is a factor of 3--6 higher than the local value.  But, depending on the stellar content the ratio may decline at $z\\gtrsim4$ (if E/S0-like), flatten off to 6--10 times the local value (if Sbc-like), or continue to rise (if Im-like).  We infer that galaxy bulge masses must have grown by a factor of 3--6 over the redshift range $3\\gtrsim z \\gtrsim 1$, and then changed little since $z \\sim 1$.  This suggests that the peak epoch of galaxy formation for massive galaxies is above $z \\sim 1$.  We also estimate the duty cycle of luminous AGNs at $z \\gtrsim 1$ to be $\\sim 1\\%$, or $10^7$ yrs, with sizable scatter. ",
        "introduction": "Elsewhere in these proceedings the reader can find summaries of previous work on quasar hosts. We concentrate here on the benefits and challenges of using gravitational lensing as a technique for measuring the host galaxy, especially in regimes where traditional direct imaging methods are difficult -- at high redshift, or when the host is either sub-luminous or compact.  The context for this work is the coevolution of galaxies and supermassive black holes.  In the local universe, a tight relation is observed between bulge mass (\\mbulge) and black hole mass (\\mbh) measured with stellar kinematics \\citep{Geb00, Fer00}. We now extend the relation to $z \\gtrsim 1$ using the virial technique to estimate $M_{BH}$ (e.g. Kaspi et al., 2000; Vestergaard \\& Peterson, 2006), and the stellar luminosity to estimate \\mbulge\\ \\citep{Kor95,Mag98}.  The existence, slope, and scatter of a \\mbh/\\mbulge\\ relation at high redshift can be used to analyze the relative growth rates of galaxies and their (presumably ubiquitous) central engines.   The CfA-Arizona Space Telescope Lens Survey (CASTLES) is a project to image all known, multiply-imaged, quasars in a homogeneous set of optical and near infrared passbands using the {\\it Hubble Space Telescope} ({\\it HST}). With the number of lens systems now near 100, data are in hand for 80 targets.  The observations are shallow, 1-2 orbits per filter, but the excellent surface brightness sensitivity of the {\\it HST} leads to a host detection in most cases.  The overall CASTLES project is described by \\citet{Fal01}; early results in the lensed host search are given by \\citet{Rix01} in the same conference proceedings; and detections and models of the hosts of several individual quasars have been published \\citep{Imp98,Koc00,Kee00}.  Gravitational lensing is a large and growing field of astrophysics so only the rudiments can be given here; a number of excellent book and reviews are available for the full formalism and diverse applications \\citep{Bla92,Sch92,Cla02}. About one in 500 quasars has a sight-line passing close enough to the central potential of a massive galaxy for multiple image formation. It has taken surveys of tens of thousands of radio and optically selected quasars to yield the sample of $\\sim$100 objects (see the CASTLES web site\\footnote{http://www.cfa.harvard.edu/castles} and that of Li\\'ege group\\footnote{http://vela.astro.ulg.ac.be/themes/extragal/gravlens}). Even though adaptive optics techniques from the ground are improving, stable point spread functions, obtained with {\\it HST}, are still essential for reliable modeling.   In gravitational lensing, the AGN is magnified into multiple images, but remains unresolved, whereas the extended light from the host galaxy maps into arcs or Einstein rings (ER).  A lens model is needed to extract the full information content of the lensed host light.  In principle, a typical ER having a radius $\\sim$ 1 arcsec means that {\\it HST} imaging potentially yields 50-100 resolution elements in the host galaxy in the deep images from the survey.  \\begin {figure} \\hskip 1.2cm \\vskip -0.1in \\vbox{ \\hbox{ \\hskip +1.0in \\psfig{file=pg1115.ps,height=3.5truein,angle=0} } } \\vskip -0.1in  \\caption {An example of the lens modeling technique.  (a) Original NIC2 image of PG~1115+080 ($z_{\\mbox QSO} = 1.72$).  (b) The host galaxy Einstein ring, after removing the best fit lensing galaxy and quasar point sources.  (c) The best fitting residuals.  (d) The parametric model of the host galaxy in the source plane.} \\label{fig:pg1115}  \\end {figure} ",
        "conclusions": "Detailed modeling of 30 well-observed systems from a total sample of 80 lensed quasars has provided new insights into the properties of host galaxies at $1 < z < 4.5$.  About half have S\\'ersic model fits indicative of early type galaxies.  However, combined with their small sizes of $r_e < 6$ kpc, luminosities, and \\mbh, it appears that luminous, fully-formed, ellipticals are in a minority as hosts of luminous quasars at $z\\gtrsim 2$. No difference is seen between the luminosities of radio-loud and radio-quiet quasars in the sample, with a caveat on sample selection.  Even at $z \\gtrsim 2$, the host galaxies follow nearly the same relationship between \\mbh\\ and luminosity as at low redshifts, but the bulges must gain in mass by a factor of 3-6 between $1\\lesssim z \\lesssim 4.5$. However, by $z\\approx1$, the mass deficit is mostly gone.  Thus massive bulges at $z\\approx 1$ may be consistent with being passively evolving, or may still grow by at most a factor of 1.  Our estimate of the AGN duty cycle is $\\approx 1\\%$, or $10^7$ years.  Ongoing work includes using color information to constrain the host star formation histories, obtaining sub-millimeter data to measure star formation rate, and characterizing detailed host morphology with deeper {\\it HST} imaging."
    },
    "0601/astro-ph0601672_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec:1}  The case for alternate gravity is easily made. The best that can be done from observing cosmic motions is to infer the metric $g_{\\mu\\nu}$ in some coordinate system. From this one can reconstruct the Einstein tensor and then ask whether or not general relativity predicts it in terms of the observed sources of stress-energy, \\begin{equation} \\Bigl(R_{\\mu\\nu} - \\frac12 g_{\\mu\\nu} R\\Bigr)_{\\rm rec} = 8 \\pi G \\Bigl( T_{\\mu\\nu}\\Bigr)_{\\rm obs} \\; ? \\end{equation} One way of explaining any disagreement is by positing the existence of an unobserved, ``dark'' component of the stress-energy tensor, \\begin{equation} \\Bigl(T_{\\mu\\nu}\\Bigr)_{\\rm dark} \\equiv \\frac1{8\\pi G} \\Bigl(R_{\\mu\\nu} - \\frac12 g_{\\mu\\nu} R\\Bigr)_{\\rm rec} - \\Bigl(T_{\\mu\\nu}\\Bigr)_{\\rm obs} \\; . \\end{equation} This always works, but recent observations make it seem epicyclic.  The theory of nucleosynthesis implies that no more than about 4\\% of the energy density currently required to make general relativity agree with all observations can consist of any material with which we are presently familiar \\cite{BBN} --- and only a fraction of this 4\\% is observed. Just to make general relativity agree with the observed motions of galaxies and galactic clusters we must posit that {\\it six times} the mass of ordinary matter comes in the form of nonbaryonic, cold dark matter \\cite{CDM}. Although there are some plausible candidates for what this might be, no Earth-bound laboratory has yet succeeded in detecting it.  I belong to the minority of physicists who feel that this factor of six already strains credulity. Easing that strain is what led Milgrom to propose MOND \\cite{MM}, which can be viewed as a phenomenological modification of gravity in the regime of very small accelerations. There is an impressive amount of observational data in favor of this modification \\cite{SM} --- although see \\cite{GSKVK}. Bekenstein has recently constructed a fully relativistic field theory \\cite{JDB} which reproduces MOND, and a preliminary analysis of the resulting cosmology works better than many experts thought possible \\cite{SMFB}.  However, the worst problem for conventional gravity comes on the largest scales. To make general relativity agree with the Hubble plots of distant Type Ia supernovae \\cite{SNCP,SNST,SNLS}, with the power spectrum of anisotropies in the cosmic microwave background \\cite{CMB} and with large scale structure surveys \\cite{LSS}, one must accept an additional component of ``dark energy'' that is about {\\it eighteen times} larger than that of ordinary matter. This would mean that 96\\% of the current universe's energy exists in forms which have so far only been detected gravitationally! Even people who believe passionately in dark matter (and hence accept the factor of six) find this factor of $6 \\!+\\! 18 \\!=\\! 24$ difficult to swallow. That is why there has been so much recent interest in modifying gravity to make it predict observed cosmic phenomena without the need for dark energy, and sometimes even without the need for dark matter.  I want to stress that the issue is one of plausibility. There is no problem inventing field theories which give the required amount of dark energy. The simplest way of doing it is with a minimally coupled scalar \\cite{CW,PR}, \\begin{equation} \\mathcal{L} = -\\frac12 \\partial_{\\mu} \\varphi \\partial_{\\nu} g^{\\mu\\nu} \\sqrt{-g} - V(\\varphi) \\sqrt{-g} \\; . \\label{quint} \\end{equation} The usual procedure is to begin with a scalar potential $V(\\varphi)$ and work out the cosmology, but it is easy to start with whatever cosmological evolution is desired and {\\it construct} the potential which would support it. I will go through the construction here, both to make the point and so that it can be used later.  On the largest scales the geometry of the universe can be described in terms of a single function of time known as the scale factor $a(t)$, \\begin{equation} ds^2 = -dt^2 + a^2(t) d\\vec{x} \\cdot d\\vec{x} \\; . \\end{equation} The logarithmic time derivative of this quantity gives the Hubble parameter, \\begin{equation} H(t) \\equiv \\frac{\\dot{a}}{a} \\; . \\end{equation} If we specialize to a solution $\\varphi_0(t)$ of the scalar field equations which depends only upon time, the two nontrivial Einstein equations are, \\begin{eqnarray} 3 H^2 & = & 8 \\pi G \\Bigl(\\frac12 \\dot{\\varphi}_0^2 + V(\\varphi_0)\\Bigr) \\; , \\label{E1} \\\\ -2 \\dot{H} - 3 H^2 & = & 8 \\pi G \\Bigl(\\frac12 \\dot{\\varphi}_0^2 - V(\\varphi_0)\\Bigr) \\; . \\label{E2} \\end{eqnarray} Let us assume $a(t)$ is known as an explicit function of time, and construct $\\varphi_0(t)$ and $V(\\varphi)$. By adding (\\ref{E1}) and (\\ref{E2}) we obtain, \\begin{equation} -2 \\dot{H} = 8 \\pi G \\dot{\\varphi}_0^2 \\; . \\label{twoeqns} \\end{equation} The weak energy condition implies $\\dot{H}(t) \\leq 0$ so we can take the square root and integrate to solve for $\\varphi_0(t)$, \\begin{equation} \\varphi_0(t) = \\varphi_I \\pm \\int_{t_I}^t dt' \\sqrt{\\frac{-2 \\dot{H}(t')}{ 8 \\pi G}} \\; . \\label{phi} \\end{equation} One can choose $\\varphi_I$ and the sign freely.  Because the integrand in (\\ref{phi}) is always positive, the function $\\varphi_0(t)$ is monotonic. This means we can invert to solve for time as a function of $\\varphi_0$. Let us call the inverse function $T(\\varphi)$, \\begin{equation} \\psi = \\varphi_0\\Bigl(T(\\psi)\\Bigr) \\; . \\label{inv} \\end{equation} By subtracting (\\ref{E2}) from (\\ref{E1}) we obtain a relation for the scalar potential as a function of time, \\begin{equation} V = \\frac1{8\\pi G} \\Bigl( \\dot{H}(t) + 3 H^2(t)\\Bigr) \\; . \\end{equation} The potential is determined as a function of the scalar by substituting the inverse function (\\ref{inv}), \\begin{equation} V(\\varphi) = \\frac1{8\\pi G} \\Biggl\\{ \\dot{H}\\Bigl(T(\\varphi)\\Bigr) + 3 H^2\\Bigl(T(\\varphi)\\Bigr) \\Biggr\\} \\; . \\end{equation}  This construction gives a scalar which supports any evolution $a(t)$ (with $\\dot{H}(t) < 0$) all by itself. Should you wish to include some other, known component of the stress-energy, simply add the energy density and pressure of this component to the Einstein equations, \\begin{eqnarray} 3 H^2 & = & 8 \\pi G \\Bigl(\\frac12 \\dot{\\varphi}_0^2 + V(\\varphi_0) + \\rho_{\\rm known}\\Bigr) \\; , \\\\ -2 \\dot{H} - 3 H^2 & = & 8 \\pi G \\Bigl(\\frac12 \\dot{\\varphi}_0^2 - V(\\varphi_0) + p_{\\rm known}\\Bigr) \\; . \\end{eqnarray} Provided $\\rho_{\\rm known}$ and $p_{\\rm known}$ are known functions of either time or the scale factor, the construction goes through as before.\\footnote{ This construction seems to be due to Ratra and Peebles \\cite{PR}. Recent examples of its use include \\cite{TW2,SRSS,NO0}.}  Using this method one can devise a new field $\\varphi(x)$ which will support {\\it any} cosmology with $\\dot{H}(t) < 0$. However, the introduction of such a ``quintessence'' field raises a number of questions: \\begin{enumerate} \\item{Where does $\\varphi$ reside in fundamental theory?} \\item{Why can't $\\varphi$ couple to fields other than the metric? And if it does couple to other fields, why haven't we detected its influence in Earth-bound laboratories?} \\item{Why did $\\varphi$ come to dominate the stress-energy of the universe so recently in cosmological time?} \\item{Why is the $\\varphi$ field so homogeneous?} \\end{enumerate} When a phenomenological fix raises more questions than it answers people are naturally drawn to investigate other fixes. One possibility is that general relativity is not the correct theory of gravity on cosmological scales.  In this talk I shall review gravitational Lagrangians of the form, \\begin{equation} \\mathcal{L} = \\frac1{16 \\pi G} \\Bigl(R + \\Delta R[g]\\Bigr) \\sqrt{-g} \\; , \\label{ansatz} \\end{equation} where $\\Delta R[g]$ is some local scalar constructed from the curvature tensor and possibly its covariant derivatives. Examples of such scalars are, \\begin{equation} \\frac1{\\mu^2} R^{\\alpha\\beta} R_{\\alpha\\beta} \\qquad , \\qquad \\frac1{\\mu^4} g^{\\mu\\nu} R_{,\\mu} R_{,\\nu} \\qquad , \\qquad \\mu^2 \\sin\\Bigl(\\frac1{\\mu^4} R^{\\alpha\\beta\\rho\\sigma} R_{\\alpha\\beta\\rho\\sigma}\\Bigr) \\; . \\end{equation} I begin by reviewing a powerful no-go theorem which pervades and constrains fundamental theory so completely that most people assume its consequence without thinking. This is the theorem of Ostrogradski \\cite{MO}, who essentially showed why Newton was right to suppose that the laws of physics involve no more than two time derivatives of the fundamental dynamical variables. The key consequence for our purposes is that the only viable form for the functional $\\Delta R[g]$ in (\\ref{ansatz}) is an algebraic function of the undifferentiated Ricci scalar, \\begin{equation} \\Delta R[g] = f(R) \\; . \\end{equation} I review the Ostrogradski result in section 2, and hopefully immunize you against some common misconceptions about it in section 3. In section 4 I explain why $f(R)$ theories do not contradict Ostrogradski's result. I also demonstrate that, in the absence of matter, $f(R)$ theories are equivalent to ordinary gravity, with $f(R) = 0$, plus a minimally coupled scalar of the form (\\ref{quint}). Then I use the construction given above to show how one can choose $f(R)$ to enforce an arbitrary cosmology. This establishes that an $f(R)$ can be found to support any desired cosmology. In section 5 I discuss problems associated with the particular choice function $f(R) = -\\frac{\\mu^4}{R}$. Section 6 presents conclusions. ",
        "conclusions": "\\label{sec:6}  The potential of a quintessence scalar can be chosen to support any cosmology, but the epicyclic nature of this construction suggests we consider modifications of gravity. Ostrogradski's theorem \\cite{MO} limits local modifications of gravity to just algebraic functions of the Ricci scalar. Models of this form can give a late phase of cosmic acceleration such as we are currently experiencing. However, they can be tuned to give anything else as well. They seem every bit as epicyclic as scalar quintessence. Further, the $f(r) = -\\mu^4/R$ model is problematic, both inside and outside matter sources.\\footnote{Observations also rule out the somewhat different version of this model that results from regarding the connection and the metric as independent, fundamental variables in the Palatini formalism \\cite{AEMM}.}  An interesting and largely overlooked possibility for modifying gravity is the fully nonlocal effective action that results from quantum gravitational corrections. In weak field perturbation theory it has long been known that the most cosmologically significant one loop corrections are not of the $R^2$ form usually studied but rather of the form $R \\ln(\\square) R$ \\cite{EMV}. More potentially interesting is the possibility of very strong infrared effects from the epoch of primordial inflation \\cite{TW3,RBM}.  It can be shown that quantum gravitational corrections to the inflationary expansion rate grow with time like powers of $\\ln(a)$. Although suppressed by very small coupling constants, the exponential growth in $a(t)$ during inflation must eventually cause the effect to become nonperturbatively strong \\cite{TW4,TW5}. Similar secular growth occurs as well for minimally coupled scalar field theories \\cite{OW1,OW2}, in which context Starobinski\\u{\\i} has developed a technique for summing the leading powers of $\\ln(a)$ at each loop order \\cite{AAS,SY}. If Starobinski\\u{\\i}'s technique can be generalized to quantum gravity \\cite{RPW5,TW6} it might result in a nonlocal effective gravity theory for late time cosmology in which a large, bare cosmological constant is almost completely screened by a nonperturbative quantum gravitational effect. In such a formalism the current phase of acceleration might result from a very slight mismatch between the bare cosmological constant and the quantum effect which screens it. It is even conceivable that one could reproduce the phenomenological successes of MOND \\cite{MM,SM} with such a nonlocal metric theory, although it would have to unstable against decay into galaxy-scale gravitational waves \\cite{SW3}.  \\vskip 1cm \\centerline{Acknowledgements} It is a pleasure to acknowledge conversations and correspondence on this subject with S. Deser, A.D. Dolgov, D.A. Eliezer, S. Odintsov, M.E. Soussa, A. Strominger and M. Trodden. This work was partially supported by NSF grant PHY-244714 and by the Institute for Fundamental Theory at the University of Florida."
    },
    "0601/astro-ph0601322_arXiv.txt": {
        "abstract": "We demonstrate that the rapid spectral variability of prompt GRBs is an inherent property of radiation emitted from shock-generated, highly anisotropic small-scale magnetic fields. We interpret the hard-to-soft evolution and the correlation of the soft index $\\alpha$ with the photon flux observed in GRBs as a combined effect of temporal variation of the shock viewing angle and relativistic aberration of an individual thin, instantaneously illuminated shell. The model predicts that about a quarter of time-resolved spectra should have hard spectra, violating the synchrotron $\\alpha=-2/3$ limit. The model also naturally explains why the peak of the distribution of $\\alpha$ is at $\\alpha\\sim-1$. The presence of a low-energy break in the jitter spectrum at oblique angles also explains the appearance of a soft X-ray component in some GRBs and their paucity. We emphasize that our theory is based solely on the first principles and contains no ad hoc (phenomenological) assumptions. ",
        "introduction": "Rapid spectral variability is a remarkable, yet unexplained feature of the prompt GRB emission. The variation of the hardness of the spectrum and the hard-to-soft evolution are the most acknowledged features \\citep{Bhat+94,Crider+97,RP02}. A quite remarkable ``tracking'' behavior, when the low-energy spectral index $\\alpha$ follows (or correlates with) the photon flux \\citep{Crider+97} is particularly intriguing. An example of such a trend is shown in Fig. \\ref{f:1} (we used data from \\citealp{Preece+00}). ",
        "conclusions": ""
    },
    "0601/astro-ph0601052_arXiv.txt": {
        "abstract": "The structure of the dark energy equation of state phase plane holds important information on the nature of the physics. We explain the bounds of the freezing and thawing models of scalar field dark energy in terms of the tension between the steepness of the potential vs.\\ the Hubble drag.  Additionally, we extend the phase plane structure to modified gravity theories, examine trajectories of models with certain properties, and categorize regions in terms of scalar field hierarchical parameters, showing that dark energy is generically not a slow roll phenomenon. ",
        "introduction": "\\label{sec.intro}  The discovery of the acceleration of the cosmic expansion has thrown physics and astronomy research into a ferment of activity, from a search for fundamental theories to investigation of predictions relating models and the cosmological dynamics, to development of astrophysical surveys yielding improved measurements.  The acceleration, or more generally the expansion history of the scale factor evolution with time, $a(t)$, can be equivalently treated by a dark energy pressure to density, or equation of state, ratio $w(a)$ \\cite{linjen}.  One model, Einstein's cosmological constant, predicts $w=-1$ at all times, but generically the dark energy phenomenon has dynamics, a time varying $w(a)$.  It is important to note that the current epoch of accelerated expansion is very different from the early epoch of inflation.  We know a priori that dark energy does not completely dominate the universe now and we do not know a priori that dark energy obeys a slow roll approximation (in fact we will see it is unlikely to).  In these senses, dark energy is a more challenging phenomenon than inflation.  We are faced with three ``Goldilocks'' problems: 1) dark energy is dynamically important, but not fully dominant, with $\\sim75\\%$ of the total energy density today, 2) the universe is accelerating, as from a component with $w\\lesssim-0.8$, but the field responsible may not be slowly rolling (unless fine tuned) as it would if nearly completely potential dominated, and 3) if it's not the cosmological constant, what happened to the cosmological constant?  To discover whether the physics is the cosmological constant, and to distinguish between alternate theories, requires measurement of the possible dynamics.  Caldwell \\& Linder \\cite{caldlin} (Paper 1) investigated the dynamics in the phase plane of $w$-$w'$, where $w'=dw/d\\ln a$, for canonical scalar field dark energy, or quintessence, finding that this reveals important clues to the nature of the new physics.  In such a plane, the time or scale factor variable is a parameter along the paths of dynamics.  They found distinct structure, categorizing fields into those that at early times are locked by Hubble friction into a cosmological constant like state, and then move away from this (thaw) as the dark energy dominates, and those that initially roll due to the steepness of the potential but later approach the cosmological constant (freeze). In this article some of these results are put on a firmer footing, examined in greater detail, and extended to models beyond canonical scalar fields, including modified gravity theories. ",
        "conclusions": "\\label{sec.concl}  We have deepened and elaborated the understanding of the role that the dark energy dynamics, through the $w-w'$ phase plane, can play in leading our understanding of the nature of dark energy.  This includes the foundations of the null line, coasting line, constant pressure line, and phantom line dividing the phase plane into distinct, physical regions.  We also elucidate the uppermost and lowermost boundaries of the thawing and freezing regions.  The physical structure has been extended beyond canonical scalar fields, including specific instances of modified gravity scenarios such as scalar-tensor, braneworld, and $H^\\alpha$ models, and barotropic and polytropic generalizations of the Friedmann equation. We outlined similarities and differences with the scalar field case, showing that many act as freezing fields, and that we should be able to clearly distinguish certain models that do not possess a deSitter future.  Mocker models, implementing a unification of dark matter and dark energy, were shown to have difficulties purely from dynamical considerations, in addition to their problems in structure formation.  Dark energy is demonstrated to be generically not amenable to a slow roll description -- a major difference from early universe inflation -- as one of its ``Goldilocks'' conundra.  This makes the dark energy problem in some sense even more challenging than the early universe. However, it also opens the possibility that if some physical bound can be placed on the flatness of the potential, e.g.\\ due to quantum corrections, then this implies a barrier around the cosmological constant $\\Lambda$ model.  This would offer hope, possibly accessible to next generation experiments, that dark energy could definitely be distinguished from $\\Lambda$, if it is not $\\Lambda$.  That would be exciting!  The dynamics of the dark energy, in the form of the equation of state ratio $w$ and its time variation $w'$, provides powerful insight into the new physics behind cosmic acceleration.   Spatial inhomogeneities in the dark energy are seen to be much weaker and less forthcoming, unless one entertains direct couplings.  While we are not guaranteed to zero in on the physics -- there is a dead zone of minimal dynamics and possibly a ``confusion'' zone near the cosmological constant -- any highly precise and accurate result would be an enormous success in enlightening us on the dark universe.  $\\,$ \\\\"
    },
    "0601/astro-ph0601578_arXiv.txt": {
        "abstract": "{} {We aim at showing that the broad-band SED characteristics of our sample of post-AGB stars are best interpreted, assuming the circumstellar dust is stored in Keplerian rotating passive discs.} {We present a homogeneous and systematic study of the Spectral Energy Distributions (SEDs) of a sample of 51 post-AGB objects. The selection criteria to define the whole sample were tuned to cover the broad-band characteristics of known binary post-AGB stars. The whole sample includes 20 dusty RV\\,Tauri stars from the General Catalogue of Variable Stars (GCVS). We supplemented our own Geneva optical photometry with literature data to cover a broad range of fluxes from the UV to the far-IR.} {All the SEDs display very similar characteristics: a large IR\\,excess with a dust excess starting near the sublimation temperature, irrespective of the effective temperature of the central star. Moreover, when available, the long wavelength fluxes show a black-body slope indicative of the presence of a component of large mm sized grains.} {We argue that in all systems, gravitationally bound dusty discs are present. The discs must be puffed-up to cover a large opening angle for the central star and we argue that the discs have some similarity with the passive discs detected around young stellar objects. We interpret the presence of a disc to be a signature for binarity of the central object, but this will need confirmation by long-term monitoring of the radial velocities. We argue that dusty RV\\,Tauri stars are those binaries which happen to be in the Population\\,{\\rm II}\\,instability strip.} ",
        "introduction": "Post-AGB stars are low and intermediate initial mass ($\\leq\\,8$-$9$\\,M$_{\\sun}$) stars which have suffered a large and dusty mass-loss phase at the end of the Asymptotic Giant Branch (AGB) during which almost the whole stellar envelope was expelled. They are evolving at constant luminosity on a fast evolutionary track in which the central star crosses the HR-diagram from a cool AGB photosphere to the ionizing temperature of the central star of a Planetary Nebula (PN). Given the small evolutionary timescale of about $10^4$\\,years, post-AGB stars are rare and not many are known \\citep[e.g][]{01Szczerba,03VanWinckel}.  Discussions on the morphology of PNe usually start with a display of the aesthetic pictures of the Hubble Space Telescope (HST) showing the complex geometry and structure of the nebulae, immediately followed by the puzzling and contradictory finding that, on the AGB, the mass-loss is found to be spherically symmetric. During the transition time, the star and circumstellar envelope must undergo fundamental and rapid changes in structure, mass-loss mode and geometry which are still badly understood. The debate on which physical mechanisms are driving the morphology changes gained even more impetus from the finding that also resolved cooler post-AGB stars or Proto-Planetary Nebulae (PPNe) display a surprisingly wide variety in shapes and structures, very early in their post-AGB evolution \\citep[][and references therein]{02Balick}.  In a survey of 27 PPNe, 21 were found to be resolved \\citep{00Ueta}. Moreover, the degree of asymmetry could be linked to the pole-to-equator density contrast, as determined on the basis of high spatial resolution mid-infrared images \\citep{99Meixner}. Recently also \\citet{05Gledhill} found evidence for axi-symmetry in the dust density in his polarimetric imaging survey of candidate post-AGB stars. The detached shells correspond to stars with an optically thin expanding circumstellar envelope whereas the bipolar and unresolved targets have optically thick dust structures, probably in the form of discs. It is suggested once again that this bifurcation in morphology is rooted in the presence or absence of a binary companion, which determines whether or not a disc forms.  Stunning kinematic information resulted from the extensive CO survey of \\citet{01Bujarrabal}: it appears to be a fundamental property of the omnipresent fast molecular outflow in PPNe, that it carries a huge amount of linear momentum, up to $1000$ times the momentum available for a radiation driven wind. Clearly, other momentum sources have to be explored. Some molecular jets are resolved by high spatial resolution imaging \\citep{04Sahai}. The formation process of the strongly collimated jets is, however, still badly understood. An intriguing suggestion is that the processes similar to the jet formation in low-mass young stellar objects operate and that the jets are born in accretion discs. This mechanism requires a significant amount of mass orbiting the post-AGB star. Such a disc could be present, but likely only in binary stars. Testing of such an hypothesis is severely hampered by the lack of observational information on binarity in PNe and PPNe but also on our poor theoretical understanding of AGB evolution in binary systems. Note that most objects which were detected in CO are strongly embedded and the sample is probably biased towards more massive PPNe.  Probing binarity in PNe and embedded PPNe directly with radial velocity monitoring is not easy. For optically bright post-AGB stars, this is different and some famous examples exist in which it is clear that the binary nature of the central object must have played a key role in the evolution of the system.  The most famous example is certainly HD\\,44179, the central star of the Red Rectangle nebula. It displays a huge and broad IR\\,excess. The IR luminosity is $33$ times stronger than the optical luminosity \\citep{89Leinert,96Waelkens}. \\citet{95VanWinckel} have shown that HD\\,44179 is a spectroscopic binary with an orbital period of $298$ days. The eccentricity is a remarkably high $e=0.37$. Since its discovery by \\citet{75Cohen}, HD\\,44179 has often been used as an archetypical example of a C-rich post-AGB object, but it is now generally accepted that many of the remarkable phenomena and the peculiar morphology of the nebula \\citep[for an overview see][]{04Cohen} are closely related to the presence of a stable circumbinary disc around the binary central star. The longevity of the disc was dramatically confirmed by the detection of cool O-rich crystalline silicate dust grains in the disc \\citep{98Waters}. The mixed chemistry is best explained assuming the formation of the O-rich disc predated the more recent C-rich transition of the central object. Recent Spitzer data show, however, that also far from the central star, the nebula appears to have a dust component which is O-rich \\citep{05Markwick-Kemper}. The (chemical) history of the binary, disc and nebula is therefore far from understood. The disc is resolved in ground-based high spatial-resolution imaging at optical and near-IR wavelengths \\citep[][and references therein]{02Menshchikov} as well as in HST optical images \\citep{04Cohen}. The disc was also resolved in interferometric CO(2-1) maps, and the Keplerian kinematics of the disc were directly detected \\citep{03Bujarrabal,05Bujarrabal}. HD\\,44179 shows a considerable amount of dust processing in the disc with indications of the presence of very large grains \\citep{97Jura} and possibly even macro-structures \\citep{98Jura}.  Another remarkable evolved object with a long-lived disc is HR\\,4049. It is a binary with an orbital period of $430$ days with a remarkably high eccentricity of $e=0.30$. Also in this object, the circumstellar material shows a mixed chemistry with both carbon rich and oxygen rich features. The SED of the dust is also very peculiar as it can be fitted with a single black-body of about 1150\\,K, from $1\\,\\mu$m down to $850\\,\\mu$m. These SED characteristics are very constraining and the best model for the circumstellar material is, that the dust is trapped in a very opaque dust torus at Keplerian rotation \\citep{03Dominik}. It is clear that also in this object the dusty disc plays a lead role in the (future) evolution. Other examples exist and there is substantial observational evidence that the systems are all likely surrounded by a circumstellar orbiting disc \\citep{03VanWinckel}. Note that so far only for the Red Rectangle the Keplerian disc is spatio-kinematically resolved by CO interferometric maps \\citep{05Bujarrabal}.   To gain insight in the evolution of binary systems and their circumstellar material, we report in this contribution on a homogeneous and systematic study of a sample of objects with similar IR\\,characteristics as the known binary post-AGB stars. The main aim is to study the broad-band SEDs in a systematic way, in order to gain insight in the possible evolutionary link between the different objects.  Our sample and the selection criteria are presented in Sect.~\\ref{section:stars}. In Sect.~\\ref{section:data} we present our photometric data complemented with the literature values and the results of the large all sky surveys. In Sect.~\\ref{section:SED} we present the detailed Spectral Energy Distribution (SED) construction of all individual objects. Distances to the objects are estimated in Sect.~\\ref{section:D}. We end this extensive study with a detailed discussion in Sect.~\\ref{section:discussion}. Conclusions are summarized in Sect.~\\ref{section:conclusions}. ",
        "conclusions": "\\label{section:conclusions} A homogeneous and systematic study of the Spectral Energy Distributions (SEDs) of a sample of 51 post-AGB objects was presented. The selection criteria to define the sample were tuned to cover the broad-band characteristics of known binary post-AGB stars. Our whole sample included 20 dusty RV\\,Tauri stars from the General Catalogue of Variable Stars.  We conclude that the broad-band SED characteristics of the sample are best interpreted, assuming the circumstellar dust is stored in Keplerian rotating passive discs. First-order estimates indicate that the discs are small in the radial direction and, given the large distances of these stars, they are likely only resolvable with interferometric observations.  The actual structure of those discs, let alone their formation, stability and evolution are not well understood. In Table~\\ref{tab:Porbit}, we list the orbital elements of $15$ objects of our total sample for which the orbital elements are covered in the literature. These systems are now not in contact (assuming typical post-AGB luminosities), but the orbits are all too small to accomodate a full grown AGB star. It is clear that those systems did not evolve on single star evolutionary tracks. They must have suffered a phase of strong interaction while the primary was at giant dimension. It is well known that those orbital elements are not very well understood yet, since most objects display a significant eccentricity as well \\citep[][and references therein]{03VanWinckel}. One of the models argues that it is the feedback of the disc on the orbital dynamics, which will induce a large eccentricity \\citep{91Artymowicz}.  We postulate that the same formation mechanism could apply to the whole sample, which means that the objects must all be binaries. Direct detection of orbital motion and, even better, determination of the orbital elements are clearly needed but poses a non-trivial observational challenge and certainly requires a long-term project.  As our sample contains a significant fraction of all known post-AGB stars, this would lead to the conclusion that binarity and binary interaction are a widespread phenomenon among \\mbox{(post-)} AGB stars. Dusty RV\\,Tauri stars are those evolved binaries which happen to be in the Population\\,{\\rm II}\\,Cepheid instability strip \\citep{99LloydEvans}.  Since the orbital periods detected till today, span already a wide range, from $115$ to about $2600$ days, we conclude that the formation of a stable orbiting disc is common and appears more and more to be a key ingredient in understanding the final evolution of a very significant fraction of binary post-AGB stars."
    },
    "0601/astro-ph0601473.txt": {
        "abstract": "{ The time delays between the components of a lensed quasar are basic tools to analyze the expansion of the Universe and the structure of the main lens galaxy halo. In this paper, we focus on the variability and time delay of the double system SBS 0909+532A,B as well as the time behaviour of the field stars. We use $VR$ optical observations of SBS 0909+532A,B and the field stars in 2003. The frames were taken at Calar Alto, Maidanak and Wise observatories, and the $VR$ light curves of the field stars and quasar components are derived from aperture and point--spread function fitting methods. We measure the $R$--band time delay of the system from the $\\chi^2$ and dispersion techniques and 1000 synthetic light curves based on the observed records. One nearby field star (SBS 0909+532c) is found to be variable, and the other two nearby field stars are non--variable sources. With respect to the quasar components, the $R$--band records seem more reliable and are more densely populated than the $V$--band ones. The observed $R$--band fluctuations permit a pre--conditioned measurement of the time delay. From the $\\chi^2$ minimization, if we assume that the quasar emission is observed first in B and afterwards in A (in agreement with basic observations of the system and the corresponding predictions), we obtain $\\Delta \\tau_{BA}$ = $-$ 45 $^{+ 1}_{-11}$ days (95\\% confidence interval). The dispersion technique leads to a similar delay range. A by--product of the analysis is the determination of a totally corrected flux ratio in the $R$ band (corrected by the time delay and the contamination due to the galaxy light). Our 95\\% measurement $\\Delta m_{BA}$ = $m_B(t + \\Delta \\tau_{BA}) - m_A(t)$ = 0.575 $\\pm$ 0.014 mag is in excellent agreement with previous results from contaminated fluxes at the same time of observation. ",
        "introduction": "The system SBS 0909+532 was discovered by Stepanyan et al. (1991). Some years later, a collaboration between the Hamburger Sternwarte and the Harvard--Smithsonian Center for Astrophysics resolved the system into a pair of quasars (A and B) with a direct $R$--band flux ratio (at the same time of observation) $\\Delta m$ = $m_B - m_A$ = 0.58 mag and a separation of about 1\\farcs1 (Kochanek et al. 1997). The direct $R$--band flux ratio was not consistent with the direct flux ratios at other wavelengths: $\\Delta m$ = 0.31 mag in the $I$ band and  $\\Delta m$  = 1.29 mag in the $B$ band. From observations with the 4.2 m William Herschel Telescope, a Spanish collaboration got spectra for each component of the system. The data showed that the system consists of two quasars with the same redshift  ($z_s$ = 1.377) and identical spectral distribution, supporting the gravitational lens interpretation of SBS 0909+532 (Oscoz et al. 1997). Oscoz et al. (1997) detected a \\ion{Mg}{ii} doublet in absorption at the same redshift ($z_{abs}$ = 0.83) in both components, and they suggested that the absorption features were associated with the photometrically unidentified lensing galaxy. Through a singular isothermal sphere (SIS) lens model, the authors also inferred the first constraint on the time delay between the components: $|\\Delta \\tau_{BA}| \\leq$ 140 days, where $\\Delta \\tau_{BA}$ is the delay of B with respect to A and the Hubble constant is assumed to be $H_0$ = 70 km s$^{-1}$ Mpc$^{-1}$.  In recent years, Lubin et al. (2000) indicated the possible nature of the main deflector (early--type galaxy) and confirmed its redshift ($z_d$ = 0.830). Leh\\'ar et al. (2000) reported on a program including Hubble Space Telescope (HST) observations of SBS 0909+532. They discovered the main lens galaxy between the components, which has a large effective radius, with a correspondingly low surface brightness. This lens galaxy is closer to the brightest component (A), which is not in contradiction with SIS--like lens models when the farther and fainter component (B) is stronger affected by dust extinction (see below). The colors of the lens are consistent with those of an early--type galaxy at redshift 0.83. Assuming a singular isothermal ellipsoid (SIE) model, Leh\\'ar et al. predicted a time delay $\\Delta \\tau_{BA}$ in the range [$-$ 10, $-$ 87] days ($H_0$ = 70 km s$^{-1}$ Mpc$^{-1}$). At a given emission time, the sign \"$-$\" means that the corresponding signal is observed first in B and later in A. The COSMOGRAIL collaboration provided the distribution of predicted time delays of the system (Saha et al. 2005). In their histogram (Fig. 10 of Saha et al.), there are two features: the main feature is an asymmetric peak around $-$ 80 days and the secondary one is another asymmetric peak around $-$ 45 days. Therefore, if the COSMOGRAIL predictions are right, the time delay is very probably of 2--3 months (component B leading component A), but we cannot rule out a delay of about one and a half months. On the other hand, the flux ratio anomaly pointed by Kochanek et al. (1997) was confirmed and accurately studied by Motta et al. (2002) and Mediavilla et al. (2005), who reported the existence of differential extinction in the main lens galaxy. Chartas (2000) and Page et al. (2004) also studied the system in the X--ray domain.  Time delays are basic tools to discuss the present expansion rate of the Universe and the structure of the main lens galaxy haloes (e.g., Refsdal 1964; Kochanek, Schneider \\& Wambsganss 2004), so that variability studies are crucial. While some time delays have been measured from radio light curves (PKS 1830$-$211: Lovell et al. 1998; Q0957+561: Haarsma et al. 1999; B0218+357: Biggs et al. 1999; B1600+434: Koopmans et al. 2000; B1422+231: Patnaik \\& Narasimha 2001; B1608+656: Fassnacht et al. 2002) or X--ray variability (e.g., Q2237+0305: Dai et al. 2003), an important set of delays are based on optical monitoring of gravitationally lensed quasars. Optical frames taken at Apache Point Observatory, Fred Lawrence Whipple Observatory and Teide Observatory were used to estimate a 14--month delay for the double system Q0957+561 (e.g., Pelt et al. 1996; Kundi\\'c et al. 1997; Serra--Ricart et al. 1999; Ovaldsen et al. 2003). Although the time delay of this first multiple quasar has been confirmed through independent observations, the measurement is only 5\\% accurate, or equivalently, there is an uncertainty of about 20 days (Goicoechea 2002). The Tel--Aviv University (TAU) group have recently determined the time delay between the two components of HE 1104$-$1805 (Ofek \\& Maoz 2003). The TAU delay of HE 1104$-$1805 disagrees with the earlier estimation by Gil--Merino, Wisotzki \\& Wambsganss (2002), but it is in excellent agreement with the determination by Wyrzykowski et al. (2003). Schechter et al. (1997) measured two delays for the quadruply imaged quasar PG 1115+080. The Belgian--Nordic collaboration carried out a very intense activity during the past five years. They participated in several monitoring projects and measured several time delays at optical wavelengths: B1600+434 (Burud et al. 2000), HE 2149$-$2745 (Burud et al. 2002a), RXJ 0911.4+0551 (Hjorth et al. 2002), SBS 1520+530 (Burud et al. 2002b) and FBQ 0951+2635 (Jakobsson et al. 2005). The formal accuracies of these 5 estimations range from 5 to 25\\% (the 1$\\sigma$ error bars vary from 4 to 24 days). Kochanek et al. (2005) also measured the time delays between the components of the quadruple quasar HE 0435$-$1223.  The aim of this paper is to present $VR$ observations of SBS 0909+532 in 2003 conducted by the University of Cantabria (UC, Spain), the Institute of Astronomy of Kharkov National University (IAKhNU, Ukraine), the Sternberg Astronomical Institute (SAI, Russia) and the Ulug Beg Astronomical Institute of Uzbek Academy of Science (UBAI, Uzbekistan). We also present TAU observations of the field stars in 2003, which have been kindly made available to us. This new optical monitoring campaign was carried out at the Calar Alto Observatory (Spain), the Maidanak Observatory (Uzbekistan) and the Wise Observatory (Israel), and the frames were taken with the 1.5 m Spanish telescope, the 1.5 m AZT$-$22 telescope at Mt. Maidanak and the Wise Observatory 1 m telescope (Section 2). In Section 3, we describe the methodology to obtain the fluxes of the quasar components and the field stars. The $VR$ light curves are also shown in Section 3. Section 4 is devoted to the time delay estimation from the light curves of A and B (quasar components) in the $R$ band. Finally, in Section 5 we summarize our conclusions and discuss the feasibility of an accurate determination of the cosmic expansion rate and the surface density in the main lensing galaxy. ",
        "conclusions": "Nowadays several groups are trying to coordinate the rich but scattered research potential in the field of gravitationally lensed quasar monitoring. The goals are to rationalize the astronomical work and to catalyze big scientific collaborations so that the astrophysics community can get a significant progress in the understanding of the central engine in lensed quasars, the structure of the lensing galaxies and the physical properties of the Universe as a whole. Some examples about that are the Astrophysics Network for Galaxy LEnsing Studies (ANGLES, http://www.angles.eu.org/), the Cosmic Lens All-Sky Survey (CLASS, http://www.aoc.nrao.edu/$\\sim$smyers/class.html) and the COSmological MOnitoring of GRAvItational Lenses (COSMOGRAIL, http://www.cosmograil.org/). The University of Cantabria group (Spain), three groups of the former Soviet Union (Institute of Astronomy of Kharkov National University, Ukraine, Sternberg Astronomical Institute, Russia, and Ulug Beg Astronomical Institute of Uzbek Academy of Science, Uzbekistan) and the Tel--Aviv University group (Israel) are also carrying out a series of initiatives to better exploit the recent individual monitoring campaigns as well as to solidify some future common project. In this paper we present the first collaborative programme on the variability of the double quasar SBS 0909+532A,B. The $VR$ observations of the system and the field stars were made with three modern ground--based telescopes in the year 2003.  The SBS 0909+532c star (N23210036195 in the GSC2.2 Catalogue) at ($\\alpha$, $\\delta$) = (09:12:53.59, +52:59:39.82) in J2000 coordinates is found to be variable, with two different levels of flux. The $VR$ gap between the low--state and the high--state is of 80--100 mmag, and the low--state lasts about one month. In the high--state the star also seems to vary, but these small--amplitude variations are not so significant as the gap between states. We want to remark the variability of this nearby star (\"c\" star), and to encourage colleagues to follow-up its fluctuations and identify the kind of variable source. The \"c\" star cannot be used as the reference object (differential photometry), because it introduces a zero--lag global correlation between the light curves of the quasar components A and B. However, the \"a--b\" nearby stars are non--variable sources, and we choose the \"b\" star as the reference candle. On the other hand, the \"s1\" and \"s2\" stars are relatively far objects, which were proposed as good references in a previous analysis (Nakos et al. 2003). However, the new $R$--band light curve $m_{s2} - m_{s1}$ reveals the variability of one (\"s1\" or \"s2\") or both stars. This variability could be either a very rare phenomenon or a consequence of doing more refined measurements (aperture or PSF fitting on several frames per night). We warn about the possible problems with this pair of stars and think it merits more attention. The point--spread function (PSF) fitting methods permit to resolve the two components of the quasar and to derive the $VR$ light curves of each component. These new $VR$ light curves represent the first resolved brightness records of SBS 0909+532. Although the $V$--band curves are interesting, the $R$--band records seem more reliable and are more densely populated. The $R$--band curves show a moderate variability through 2003, and the observed fluctuations are promising for different kinds of future studies.  To estimate the time delay between the components of SBS 0909+532, we use an 120--day piece of the $R$\u0096-band brightness records, and $\\chi^2$ and dispersion ($D^2$) techniques. The cross--correlation of the two light curves (A and B) leads to complex $\\chi^2$ spectra. However, assuming that the quasar emission is observed first in B and afterwards in A, or in other words, $\\Delta \\tau_{BA} <$ 0 (in agreement with basic observations of the system), 95\\% measurements $\\Delta \\tau_{BA}$ = $-$ 45 $^{+ 1}_{-11}$ days and $\\Delta m_{BA}$ = 0.590 $\\pm$ 0.014 mag are inferred from 1000 repetitions of the experiment (synthetic light curves based on the observed records). From the $D^2$ minimization (Pelt et al. 1996) and 1000 repetitions, we also obtain 90\\% measurements $\\Delta \\tau_{BA}$ = $-$ 48 $^{+ 7}_{-6}$ days and $\\Delta m_{BA}$ = 0.585 $\\pm$ 0.020 mag. The $D^2$ uncertainties are derived under the already mentioned assumption that $\\Delta \\tau_{BA}$ is negative. There is a clear agreement between the results from both techniques, so a delay value of about one and a half months is strongly favoured. Our light curves rule out a delay close to three months, which has been claimed in a recent analysis (Saha et al. 2005). When we measure the time delay of the system, we simultaneously derive the time--delay--corrected flux ratio (at the same emission time) in the $R$ band. This quantity, $\\Delta m_{BA}$ = $m_B(t + \\Delta \\tau_{BA}) - m_A(t)$, is contaminated by light of the lens galaxy, and taking into account the weak contaminations of A and B (see the end of subsection 3.2), the totally corrected $R$--band flux ratio is 0.575 $\\pm$ 0.014 mag. We remark that our final $R$ flux ratio is in total agreement with the rough (uncorrected by the time delay and the contamination by galaxy light) measurement by Kochanek et al. (1997): 0.58 $\\pm$ 0.01 mag. To properly determine a flux ratio, one must use clean fluxes at the same emission time, i.e., fluxes at different observation times and without contamination (Goicoechea, Gil--Merino \\& Ull\\'an 2005). Only for particular cases (e.g., faint lens galaxy, short delay and moderate variability), it may be reasonable to use direct fluxes.  In order to get a reasonably good value of $\\chi^2$, we do not need to introduce a time dependent magnitude offset or a complex iterative procedure (e.g., Burud et al. 2000; Hjorth et al. 2002), i.e., only a delay and a constant offset are fitted. This is a strong point of the analysis. The agreement between the results from different techniques is another strong point. However, the new measurements have some weak points that we want to comment here. The weakest point is the relatively poor overlap between the A and B records, when the A light curve is shifted by the best solutions of the time delay and the magnitude offset (e.g., see the top panel of Fig. 13). Moreover, we carry out pre--conditioned measurements, since a negative interval [$-$ 90, 0] days is considered in the estimation of uncertainties (component B leading component A). This second weak point is related to the presence of 10/20--day gaps and the moderate variability of the components, which does not permit to fairly rule out positive delays. We nevertheless remark that the negative interval is in good agreement with the predictions by Leh\\'ar et al. (2000) and Saha et al. (2005), and we find $\\chi^2$ and $D^2$ minima around $-$ 45 days when the observed data and both negative and positive lags are taken into account (see Figs. 11 and 14). Of course, as any another first determination of a time delay, the 1.5--month value should be confirmed from future studies.  Forty years ago, Refsdal (1964) suggested the possibility of determining the current expansion rate of the Universe (Hubble constant) and the masses of the galaxies from the time delays associated with extragalactic gravitational mirages. More recently, for a singular isothermal ellipsoid (SIE), Koopmans, de Bruyn \\& Jackson (1998) found that the time delay can be cast in a very simple form, depending on basic cosmological parameters, redshifts and image positions. The relevant image positions are the positions with respect to the centre of the main lens galaxy, and the SIE delay is similar to the delay for a singular isothermal sphere (SIS). In principle, a singular density distribution is justified because a small core radius changes the time delay negligibly, and only a small core radius seems to be consistent with the absence of a faint central image (e.g., Kochanek 1996). Moreover, individual lenses and lens statistics are usually consistent with isothermal models (e.g., Witt, Mao \\& Keeton 2000 and references therein), so it is common to adopt an isothermal profile. Witt, Mao \\& Keeton (2000) showed that an external shear changes the simple SIS time delay in proportion to the shear strength. For two--image lenses that have a small shear and images at different distances from the centre of the lens, the shear should have a small effect on the time delay. Thus, when one has accurate measurements of image positions, redshifts and time delay, it is viable an accurate estimation of $H_0$ (using complementary information on the matter/energy content of the Universe).  Very recently, Kochanek (2002) also presented a new elegant approach to the subject. He modelled the surface density locally as a circular power law, with a mean surface density $<\\kappa>$ in the annulus between the images. Expanding the time delay as a series in the ratio of the thickness of the annulus to its average radius, it is derived a delay that is proportional to the SIS time delay. The zero--order expansion term consists of the SIS delay and a multiplicative factor $2(1 - <\\kappa>)$. Kochanek also incorporated the quadrupoles of an internal shear (ellipsoid) and an external shear. However, for two-image lenses where the images lie on opposite sides of the lens, the delay depends little on the quadrupoles. This novel perspective is useful to infer $<\\kappa>$ from observations of the lens system (time delay, image positions and redshifts) and complementary cosmological data (expansion and matter/energy content of the Universe).  For SBS 0909+532, although the redshifts are very accurately known and the time delay is now tightly constrained (or at least there is a first accurate estimation to be independently confirmed), the inaccurate position of the main lens galaxy does not permit an accurate measurement the cosmic expansion rate and the surface density of the main deflector. We have $H_0 \\propto \\theta_B^2 - \\theta_A^2$ and $1 - <\\kappa> \\propto (\\theta_B^2 - \\theta_A^2)^{-1}$, where $\\theta_A$ and $\\theta_B$ are the image angular positions with respect to the centre of the main lens galaxy. On the other hand, using the astrometry in Table 3 of Leh\\'ar et al. (2000), it is easy to obtain $\\theta_B^2 - \\theta_A^2$ = 0.4 $\\pm$ 0.2. Thus we conclude that the accuracy in $\\theta_B^2 - \\theta_A^2$ is only 50\\%, indicating the necessity of new accurate astrometry of SBS 0909+532."
    },
    "0601/astro-ph0601374_arXiv.txt": {
        "abstract": "{We report the discovery of the faint ($V\\simeq 21.7$) quasar QSO~03027$-$0010 at $z=2.808$ in the vicinity of Q~0302$-$003, one of the few quasars observed with STIS to study intergalactic \\ion{He}{ii} absorption. Together with another newly discovered QSO at $z=2.29$, there are now 6 QSOs known near the line of sight towards Q~0302$-$003, of which 4 are located within the redshift region $2.76 \\la z \\la 3.28$ covered by the STIS spectrum. We correlated the opacity variations in the \\ion{H}{i} and \\ion{He}{ii} Lyman forest spectra with the locations of known quasars. There is no significant proximity effect in the \\ion{H}{i} Ly$\\alpha$ forest for any of the QSOs, except for the well-known line of sight effect for Q~0302$-$003 itself. By comparing the absorption properties in \\ion{H}{i} and \\ion{He}{ii}, we estimated the fluctuating hardness of the extragalactic UV radiation field along this line of sight. We find that close to each foreground quasar, the ionizing background is considerably harder than on average. In particular, our newly discovered QSO~03027$-$0010 shows such a hardness increase despite being associated with an overdensity in the \\ion{H}{i} Lyman forest. We argue that the spectral hardness is a sensitive physical measure to reveal the influence of QSOs onto the UV background even over scales of several Mpc, and that it breaks the density degeneracy hampering the traditional transverse proximity effect analysis. We infer from our sample that there is no need for significantly anisotropic UV radiation from the QSOs. From the transverse proximity effect detected in the sample we obtain minimum quasar lifetimes in the range $\\sim$10--30~Myr. ",
        "introduction": "Observations of high-redshift quasars enable us to study the intergalactic medium (IGM) along their lines of sight via the absorption of quasar radiation by various chemical elements in different ionization stages. Hydrogen and helium are by far the most abundant elements in the universe and the Ly$\\alpha$ transitions of \\ion{H}{i} and \\ion{He}{ii} in an incompletely ionized medium cause a Gunn-Peterson trough at redshifts smaller than the emission redshift of the observed quasar \\citep{gunn65}. If the IGM is highly ionized, a plethora of discrete absorption lines stemming from the remaining neutral fraction is visible both in \\ion{H}{i} Ly$\\alpha$ and \\ion{He}{ii} Ly$\\alpha$, giving rise to the name Ly$\\alpha$ forest.  The intensity of the metagalactic UV radiation field at a characteristic frequency (typically the ionization energy of a given element) varies in time due to the temporal evolution of the source population. Spatial fluctuations induced by the discreteness of the source population have been examined for randomly distributed sources and absorbers by \\citet{zuo92} and \\citet{fardal93}. They are also treated in recent numerical simulations trying to quantify their impact on the Ly$\\alpha$ flux power spectrum \\citep{croft99,meiksin04,croft04,mcdonald05}. The spatial fluctuations of the UV radiation field computed from these simulations along random lines of sight are generally gentle (few per cent around the mean) and occur on large scales ($\\ga 100$~Mpc comoving) at $z<4$.  This simple picture with only mild large-scale fluctuations is expected to change considerably for lines of sight that pass close to sources of the UV background, such as luminous quasars \\citep{fardal93,croft04,mcdonald05}. The source flux acts as a local enhancement of the UV radiation field. As a consequence, the absorbers in the region affected by this excess flux will be statistically more ionized than the rest of the Ly$\\alpha$ forest along the line of sight, resulting in a statistically increased transmission, a radiation-induced `void' in the \\ion{H}{i} Ly$\\alpha$ forest in the vicinity of the UV source. This so-called proximity effect has been detected with high statistical significance in lines of sight towards luminous quasars \\citep[e.g.][]{bajtlik88,scott00}. However, a transverse proximity effect created by foreground ionizing sources nearby the line of sight has not been clearly detected in the \\ion{H}{i} Ly$\\alpha$ forest. The full range of possible results extends from large voids claimed to be due to the transverse proximity effect by \\citet[][however see \\citealt{dobrzycki91b}]{dobrzycki91} and \\citet{srianand97}, over marginal detections \\citep{fernandez-soto95,liske01} to non-detections \\citep{crotts89,moller92,crotts98,schirber04}. \\citet{croft04} measured the average transverse Ly$\\alpha$ transmission from all projected quasar pairs in the SDSS DR1 and even found excess absorption near foreground quasars instead of the expected excess transmission caused by the transverse proximity effect.  The detectability of the transverse proximity effect is hampered by several systematic effects. Anisotropic radiation of quasars has been invoked to explain redshift offsets between the void and the foreground quasar \\citep{dobrzycki91} or to explain the lack of the transverse proximity effect \\citep{crotts89,moller92,schirber04}. Also quasar variability affects the detection of the proximity effect \\citep{schirber04}. Finally, the possible gravitational clustering around quasars that are assumed to reside in the densest environments may dilute the proximity effect \\citep{loeb95,schirber04}. According to \\citet{schirber04} only a combination of these systematic effects may explain the apparent absence of the transverse proximity effect.  \\ion{He}{ii} Ly$\\alpha$ $303.78$~\\AA~ absorption can be studied only towards a few lines of sight to date because most of the quasars at $z>2$ have intervening optically thick Lyman limit systems that truncate the flux in the observable \\ion{He}{ii} Ly$\\alpha$ wavelength range in the far UV. The observations of the lines of sight towards Q~0302$-$003 at $z=3.285$ \\citep{jakobsen94,hogan97,heap00}, \\object{PKS~1935$-$692} at $z=3.18$ \\citep{anderson99} and recently \\object{SDSS~J2346$-$0016} at $z=3.50$ \\citep{zheng04b} show in most parts of their \\ion{He}{ii} absorption spectra very strong absorption at $z>3$ that is consistent with a Gunn-Peterson trough ($\\tau_\\ion{He}{ii}>3$). In contrast, the three lines of sight at $z<3$ probed so far towards \\object{HS~1700$+$6416} at $z=2.72$ \\citep{davidsen96,reimers04}, \\object{HE~2347$-$4342} at $z=2.885$ \\citep{reimers97,kriss01,smette02,shull04,zheng04} and recently \\object{QSO~1157$+$3143} at $z\\simeq 3$ \\citep{reimers05} display patchy intergalactic \\ion{He}{ii} absorption with voids ($\\tau_\\ion{He}{ii}\\lesssim 1$) and troughs ($\\tau_\\ion{He}{ii}>3$) that evolves to a \\ion{He}{ii} Ly$\\alpha$ forest at $z<2.7$ resolved with FUSE \\citep{kriss01,shull04,zheng04,reimers04}. In combination with the observed evolution of \\ion{H}{i} line widths \\citep{schaye00,ricotti00,theuns02a} these observations point to a late \\ion{He}{ii} reionization between $z\\sim 2.7$ and $z\\sim 3$.  By comparing the \\ion{H}{i} absorption with the corresponding \\ion{He}{ii} absorption one can estimate the hardness of the ionizing radiation field that penetrates the IGM, since \\ion{H}{i} is ionized at $h\\nu>13.6$~eV, whereas \\ion{He}{ii} is ionized at $h\\nu>54.4$~eV. The amount of \\ion{He}{ii} compared to \\ion{H}{i} gives a measure of the spectral hardness. Already low-resolution \\ion{He}{ii} observations obtained with HST indicated a fluctuating radiation field in the voids (hard) and the troughs (soft) \\citep{reimers97,heap00,smette02}. The recent high-resolution FUSE observations of the \\ion{He}{ii} Ly$\\alpha$ forest reveal large fluctuations on very small scales of $\\Delta z\\sim 10^{-3}$ \\citep{kriss01,shull04,reimers04}. Due to the hard ionizing field required in the \\ion{He}{ii} voids that is consistent with the integrated radiation of a surrounding quasar population, these \\ion{He}{ii} voids have been interpreted as the onset of \\ion{He}{ii} reionization in Str\\\"{o}mgren spheres around hard \\ion{He}{ii} photoionizing sources along or near the line of sight \\citep{reimers97,heap00,smette02}. A subsequent survey for putative quasars that may cause the prominent \\ion{He}{ii} void at $z=3.05$ towards Q~0302$-$003 yielded the quasar \\object{QSO~03020$-$0014} located $6\\farcm5$ away on the sky that coincides with this \\ion{He}{ii} void \\citep{jakobsen03}. Thus, the \\ion{He}{ii} void in Q~0302$-$003 is the first clear case of a transverse proximity effect due to a luminous quasar.  In this paper we report on results from a slitless spectroscopic quasar survey that resulted in the discovery of another foreground quasar in the vicinity of Q~0302$-$003. The structure of the paper is as follows. Sect.~\\ref{observations} describes the observation and the supplementary data employed for the paper. In Sect.~\\ref{intensity_h1} and \\ref{intensity_he2} we examine the evidence for opacity variations in the \\ion{H}{i} and \\ion{He}{ii} Ly$\\alpha$ forest regions caused by intensity fluctuations in the UV radiation field towards Q~0302$-$003 at 1~ryd and 4~ryd, respectively.  In Sect.~\\ref{hardness} we consider the possibility to detect the transverse proximity effect via spectral hardness diagnostics of the ionizing radiation. We argue that this is actually the most sensitive method, and we demonstrate that essentially each quasar near the line of sight is associated with a local hardening of the radiation field. Finally, we use these observations to estimate a lower limit to the quasar lifetime (Sect.~\\ref{lifetime}). We present our conclusions in Sect~\\ref{conclusions}. Throughout the paper we adopt a flat cosmological model with $\\Omega_\\mathrm{m}=0.3$, $\\Omega_\\Lambda=0.7$ and $H_0=70$~$\\mathrm{km}\\,\\mathrm{s}^{-1}\\,\\mathrm{Mpc}^{-1}$. ",
        "conclusions": "\\label{conclusions}  The transverse proximity effect in quasar spectra has so far mostly been associated with the notion of voids in the \\ion{H}{i} Lyman forest, created by the overionized zones around quasars near the line of sight. It has always been clear, however, that true voids due to large scale structure and apparent voids due to the proximity effect would be essentially indistinguishable on an individual basis. This problem holds also for the \\ion{He}{ii} Lyman forest that has now become accessible for a small number of high-redshift QSOs.  We argue in this study that spectral hardness indicators provide an efficient way to discriminate between large scale structure and proximity effect.  The line of sight proximity effect towards Q~0302$-$003 and the transverse proximity effect due to QSO~03020$-$0014 (dubbed `QSO~A' in this paper) were already discussed by \\citet{heap00} and \\citet{jakobsen03} and are, since the discovery of QSO~A, clear-cut cases where both the optical depth ratio and the hardness parameter $\\eta$ show pronounced minima close to the quasars. We investigated the \\ion{H}{i} and \\ion{He}{ii} absorption properties at redshifts near our newly discovered QSO~03027$-$0010 (`QSO~B'), and also near the other known QSOs close to the line of sight.  In each single case we find evidence for a link between spectral hardness and the presence of a nearby quasar.  Thus, the relative UV hardness is a sensitive \\emph{physical} quantity to search for individual sources of the metagalactic UV radiation field beyond the simple detection of associations between quasars and voids. Quasars and voids may be unrelated, whereas the spectral hardness is related to the spectral energy distributions of the ionizing sources. In particular, a void will not occur if the quasar overionizes an intrinsically overdense region. The spectral hardness may break this density degeneracy that affects studies of the proximity effect \\citep{loeb95,schirber04}. At least for QSO~B, this case of an overdense region is apparently applicable.  The excellent match between the measured QSO redshifts and the inferred minima of $\\eta(z)$ imply that there is essentially no need to invoke strongly anisotropic radiation for any of the quasars in question. This is in stark contrast to the situation for the \\ion{H}{i} proximity effect where anisotropy or even beaming has repeatedly been claimed \\citep{dobrzycki91,crotts89,moller92,schirber04}. It is likely that true large scale structure in the Lyman forest has contributed to mask the intrinsic transverse proximity effect in several of these cases.  The redshift range of $2.7 \\la z < 3.2$ studied in the line of sight towards Q~0302$-$003 is of very high interest in the context of studying the putative reionization epoch of \\ion{He}{ii}. We have shown that one can determine the sources that ionize this line of sight, even at low resolution. Near the low-redshift end of the covered spectral range, there are clear indications for a transition from \\ion{He}{iii} bubbles in the IGM towards a more or less fully ionized \\ion{He}{iii} IGM where also the overall UV background is dominated by QSO radiation. Our estimates of $\\eta$ in this redshift range are fully compatible with this concept. Even then, local sources of very hard radiation such as QSO~B can significantly modify the UV background over distances of several Mpc. It would be highly desirable to be able to study these processes at higher spectral resolution. Unfortunately, Q~0302$-$003 is too faint for observations with the \\emph{Far Ultraviolet Spectroscopic Explorer}, and the only perspective for significantly improved data lies with the Cosmic Origins Spectrograph to be installed on HST."
    },
    "0601/astro-ph0601697_arXiv.txt": {
        "abstract": "We examine the outer Galactic HI disk for deviations from the $b=0\\degr$ plane by constructing maps of disk surface density, mean height, and thickness. We find that the Galactic warp is well described by a vertical offset plus two Fourier modes of frequency 1 and 2, all of which grow with Galactocentric radius. Adding the $m=2$ mode  accounts for the large asymmetry between the northern and southern warps. We use a Morlet wavelet transform to investigate the spatial and frequency localization of higher frequency modes; these modes are often referred to as ``scalloping.'' We find that the $m=$ 10 and 15 scalloping modes are well above the noise, but localized; this suggests that the scalloping does not pervade the whole disk, but only local regions. ",
        "introduction": "Although the topography of the gas disk of the Milky Way has previously been  mapped \\citep[and many others]{W1957, HJK1982,BT1986}, its complete Fourier structure has never been quantitatively described, though \\citet{BM1998} have approximated its three lowest frequency terms. Which spatial oscillation frequencies are most important? Is the scalloping a local or global effect? This paper constitutes an in depth analysis of the shape of the outer HI disk, as well as the first quantitative analysis of the scalloping.  The large-scale warp in the gas disk of the Milky Way has been known since 1957 \\citep{B1957,K1957,W1957,KHC1957}. The warp has a large amplitude, rising to a height greater than 4 kpc  at a Galactocentric radius of 25 kpc in the northern data. It is also asymmetric; in the south the gas falls about 1 kpc below the plane before rising back to it. By observing other galaxies, it is possible to develop a general understanding of how warps behave. \\citet{B1991} found that at minimum half of all galaxies are warped, and that galaxies with smaller dark matter halo core radii are less likely to be warped. \\citet{B1990} claimed that a warp's line of nodes starts out straight, and at a transition begins to advance in the direction of rotation, with some exceptions. Another survey has shown that warps are common in galaxies with HI disks that are extended compared to their optical components, and are often asymmetric \\citep{GSK2002}. In the Milky Way, many components other than HI also participate in the warp. A partial list includes  dust \\citep{FBD1994}, CO \\citep{WBB1990}, solar neighborhood stars \\citep{D1998}, and IRAS point sources \\citep{DS1989}.  Many efforts have been directed toward understanding the warp  on a theoretical basis. Bending modes have long been suspected as the mechanism creating and maintaining the warp. Early work studying the evolution of  bending mode oscillations was hampered by a lack of knowledge regarding galactic halos, but showed that the shape of the density falloff near the edge of the disk  plays an important role in the stability of bending modes \\citep{HT1969,T1983}. The distribution of matter in the halo controls the ability of the disk to sustain long-lived bending wave warps, so studying the properties of the warp will  tell us about the shape of the halo  \\citep{B1978,S1984}. \\citet{SC1988} argued that bending mode oscillations are plausible when the halo and the self-gravity of the disk are taken into account, but \\citet{BJD1998} showed that in this situation the warp will wind up within a few dynamical times. Previous work has largely been concerned with the existence and behavior of $m=1$ warps only, though \\citet{S1995} demonstrated the stability of $m=0$ modes in an axisymmetric halo.  Several other mechanisms have been suggested as possibilities for creating and maintaining a warp \\citep{KG2000}. Gravitational interaction with satellites such as the Magellanic Clouds is a   promising candidate \\citep{B1957,W1998,WB2005}, but there is a longstanding debate of whether tidal effects are strong enough to produce the observed effect \\citep{K1957}. Indeed, even in galaxies with companions similar to the LMC, tidal amplification may not be strong enough to account for the size of the warp \\citep{GKD2002}. Accretion of matter onto the halo is another plausible cause \\citep{JB1999}, or matter can accrete directly onto onto the disk and torque the gas orbits \\citep{LBB2002,S2004}.  Intergalactic magnetic fields can act on a slightly ionized gas disk to produce a warp \\citep{BFS1990}, or the intergalactic medium can excite a warp through a wind \\citep{KW1959}.   Throughout this paper, we will refer to shorter wavelength (5--25 kpc) oscillations in azimuth as scalloping. \\citet{GKW1960} first noticed a ``waviness'' in the gas layer of the inner Galaxy;  observational evidence  for scalloping turns out to be prominent in the outer Galaxy \\citep{HJK1982,KBH1982}. \\citet{F1983} investigated the possibility that the Milky Way scalloping results from the Kelvin-Helmholz instability, and \\citet{S1995} suggested that the scalloping should only be present in the outer parts of Galactic disks. \\citet{SF1986} find azimuthal corrugations in HI and other components along the spiral arms after removing the warp. In an N-Body simulation, \\citet{EE1997} found evidence for spiral corrugations due to gravitational interaction with a satellite galaxy. In this paper we will only investigate oscillation in Galactocentric azimuth, and not in radius.  In \\S 2 we describe the method of transforming the Heliocentric data cube into Galactocentric coordinates, and present maps of the surface density, average height off the $b=0\\degr$ plane, and vertical thickness of the outer Galaxy. While these maps are constructed using new data, they are not significantly different from previous work.  In \\S 3 we perform  global and local analyses on these maps to better understand the warp and the scalloping. ",
        "conclusions": "We fit the global shape of the warp on our grid of concentric rings using a vertical offset and two sinusoidal modes. Outside of $R\\approx20$ kpc, each of these three modes has  more power than any of the higher frequencies we look at. The amplitude increases with radius over our entire radius range for the 1 mode, and starting from around 15 kpc for the 0 and 2 modes. The growth of the 0 and 2 modes results in  asymmetry in the warp; this growth begins near where the stellar disk ends. The line of maxima of the $m=1$ mode is essentially coincident with one of the lines of maxima of the $m=2$ mode. There is little evidence for precession or winding of these two modes.  A global analysis with the Lomb periodogram shows that each $m$ mode evolves differently with radius. The most interesting include $m=3-6,10$, and 15, each of which start at small radii below the 95\\% significance level and then cross it further out in the disk.  An analysis combining the global Lomb periodogram and the local wavelet transform shows that none of the modes with $7\\le m\\le15$  have strength over a full ring of the disk. Using a wavelet transform, we show that the scalloping observed by previous authors near $\\ell\\sim 310\\degr$ is real. We therefore conclude that the scalloping is a local effect. Lower frequency modes proved impossible to study with wavelets due to the presence of the excluded regions."
    },
    "0601/astro-ph0601467.txt": {
        "abstract": "We report the results of a spectroscopic search for debris disks surrounding 41 nearby solar type stars, including 8 planet-bearing stars, using the {\\it Spitzer Space Telescope}. With accurate relative photometry using the  Infrared Spectrometer (IRS) between 7-34 $\\micron$ we are able to look for excesses as small as $\\sim$2\\% of photospheric levels with particular sensitivity to weak spectral features. For stars with no excess, the $3\\sigma$ upper limit in a band at 30-34 $\\mu$m corresponds to $\\sim$ 75 times the brightness of our zodiacal dust cloud. Comparable limits at 8.5-13 $\\mu$m correspond  to $\\sim$ 1,400 times the brightness of our zodiacal dust cloud. These limits correspond to material located within  the $<$1 to $\\sim$5 AU region that, in our solar system, originates from debris associated with the asteroid belt.  We find excess emission longward of $\\sim$25 $\\mu$m from five stars of which  four also show excess emission at 70 $\\mu$m. This emitting dust must be located around 5-10 AU. One star has 70 micron emission but no IRS excess.  In this case, the emitting region must begin outside 10 AU; this star has a known radial velocity planet. Only two stars of the five show emission shortward of 25 $\\micron$ where spectral features reveal the presence of a population of small, hot dust grains emitting in the 7-20 $\\mu$m band. The data presented here strengthen the results of previous studies to  show that excesses at 25 $\\micron$ and shorter are rare: only 1 star out of 40 stars older than 1 Gyr or $\\sim 2.5$\\% shows an excess.  Asteroid belts 10-30 times more massive than our own appear are rare among mature, solar-type stars. ",
        "introduction": "The infrared excess associated with main sequence stars \\citep{aumann1984} is produced by dust orbiting at distances ranging from less than 1 to more than 100 AU. In discussing these debris disks we often use the analogy of our Solar System, where dust is generated by collisions between larger bodies in the asteroid and  Kuiper belts, as well as outgassing comets. As in the Solar System, we distinguish between: a) hot (200-500 K) debris located relatively close to the star ($<$1 AU out to $\\sim$ 5 AU) and arising in a region dominated by refractory material, e.g. debris associated with an asteroid belt \\citep{dermott02};  and b) cooler debris ($<30-200$ K) located beyond 5 AU that may be associated with more volatile, cometary material, i.e. the analog of our Kuiper Belt. Observed dust luminosities range from $\\ld \\simeq 10^{-5}$ to greater than $10^{-3}$ for main-sequence stars, compared to $\\ld \\simeq 10^{-7}-10^{-6}$ inferred for our own Kuiper belt  \\citep{stern1996,  backman1995} and $\\sim 10^{-7}$ measured for our asteroid belt \\citep{bp1993, dermott02}.  The FGK survey of solar-type stars is a \\spit\\ GTO program designed to search for excesses around a broad sample of 148 nearby, F5-K5 main-sequence field stars,   sampling wavelengths from 7-34 \\um\\ with the Infrared Spectrograph (IRS; Houck et al. 2004) and 24 and \\70um with  the Multiband Imaging Photometer for Spitzer (MIPS; Rieke et al. 2004). In the context of our  paradigm,  IRS is well suited to searching for and characterizing hot dust in an asteroid belt analogue while MIPS is better suited to studying cooler material in a Kuiper Belt analogue. The sample consists of two components --- stars with planets and those without. The MIPS observations of planet-bearing stars are discussed in Beichman et al. (2005a) and of non planet-bearing stars in Bryden et al. (2006).  This paper presents the results of the IRS survey of a sub-sample of 39 stars, 7 with planets, and the remainder without planets based on current knowledge from radial velocity programs. In addition, the paper includes IRS follow-up observations of 2 stars (one of which is planet-bearing) that were not originally scheduled for spectroscopy,  but which were found to have a 70 $\\mu$m excess. Results for one star with a notable IRS spectrum, HD 69830, have been presented elsewhere \\citep{beichman2005b} but will be summarized below in the context of the complete survey.  The IRS portion of the program was designed with four goals in mind: 1) to extend the wavelength coverage beyond that possible with MIPS to look for dust  emitting either shortward or longward of MIPS's $24\\, \\mu$m band; 2) to take advantage of IRS's broad wavelength coverage to reduce systematic errors and thus, possibly, identify smaller excesses than previously  possible; 3) to constrain the spatial distribution of dust grains by assessing the range of temperatures present in a given debris disk; and 4) to search for mineralogical features in the spectra of any excesses. In this paper we describe the sample of stars ($\\S2$);  present our calibration procedure and give the reduced spectra for each star ($\\S3$ and Appendix-1); discuss detections of or limits to IR excesses ($\\S4$); present models of the excesses in the 4 stars with IRS and MIPS 70 um excesses and discuss the nature of their debris disks; discuss overall results and conclusions ($\\S5$ and $\\S6$). ",
        "conclusions": "We have used the IRS spectrometer on the Spitzer Space Telescope  to look for weak excesses around main sequence, solar type stars. After careful calibration, the IRS spectra allow us to achieve relatively low values of $\\ld$ compared to previous studies. At 8-13 $\\micron$ the IRS data provide an average 3$\\sigma$ limit of $\\ld =14 \\pm 6 \\times 10^{-5}$ or roughly  1,400 times the level of our solar system zodiacal cloud for material in the 1 AU region of the target stars.  Two stars, HD 72905 and HD 69830, have detectable emission at these wavelengths and in both cases there is evidence for emission by small dust grains. At 30-34 $\\micron$, we reach a more  sensitive, limit on $\\ld$ of $0.74 \\pm 0.31 \\times 10^{-5}$, or roughly  74 times the level of our solar system zodiacal cloud for material in the $5-10$ AU  region which may represent the inner edge of the Kuiper Belt. Four stars (HD 7570, HD 72905, HD 76151, HD 206860) showed evidence for long wavelength IRS emission that appears to link smoothly to a 70 $\\micron$ MIPS excess. HD 69830 shows no evidence for 70 $\\micron$ excess and thus for colder, more distant material.  The lack of IRS excess emission toward planet-bearing stars may be due to the small size of the sample observed to date, to the effects of planets removing dust from the orbital locations probed at IRS wavelengths, or to more general processes that remove dust from regions interior to a few AU \\citep{dominik2003}. We suggest that asteroid belts located between 1-10 AU and as massive as 8-35 times our own asteroid belt are rare around mature, solar type stars."
    },
    "0601/astro-ph0601004_arXiv.txt": {
        "abstract": "A sterile neutrino with mass of several keV can account for cosmological dark matter, as well as explain the observed velocities of pulsars. We show that X-rays produced by the decays of these relic sterile neutrinos can boost the production of molecular hydrogen, which can speed up the cooling of gas and the early star formation, which can, in turn, lead to a reionization of the universe at a high enough redshift to be consistent with the WMAP results. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601232_arXiv.txt": {
        "abstract": "Galactic nuclei and globular clusters act as laboratories in which nature experiments with normal stars, neutron stars and black holes, through collisions and through the formation of bound states, in the form of binaries.  The main difference with the usual Earth-based laboratories is that we cannot control the experiments.  Instead, we have no choice but to create virtual laboratories on Earth, in order to simulate all the relevant physics in large-scale computational experiments.  This implies a realistic treatment of stellar dynamics, stellar evolution, and stellar hydrodynamics.  Each of these three fields has its own legacy codes, workhorses that are routinely used to simulate star clusters, stars, and stellar collisions, respectively.  I outline the main steps that need to be taken in order to embed and where needed transform these legacy codes in order to produce a far more modular and robust environment for modeling dense stellar systems.  The time is right to do so: within a few years computers will reach the required speed, in the Petaflops range, to follow a star cluster with a million stars for ten billion years, while resolving the internal binary and multiple star motions.  By that time simulation software will be the main bottleneck in our ability to analyze dense stellar systems.  Only through full-scale simulations will we be able to critically test our understanding of the `microphysics' of stellar collisions and their aftermath, in a direct comparison with observations. ",
        "introduction": "In experimental high-energy physics, there are two ways to probe elementary particles.  One way is by studying the properties of their bound states, which may be more stable or more easily accessible than unbound particles.  The other way is by smashing particles together, and to study the remnants emerging from such collisions.  In this way, we have learned an enormous amount in the last century about the electroweak and strong interactions.  Gravity, however, has remained almost totally elusive.  The problem is related to the weakness of the gravitational force, which implies that we have to add quite a number of particles before gravity can dominate.  A self-gravitating bound state of nuclear matter, a neutron star, contains some $10^{57}$ particles, in a ball with a diameter of order 10 km, and a mass like that of the sun.  And the only known way to produce a purely gravitational object, a black hole, of moderate size, is to let a star implode to form an object of several solar masses with a horizon size of the order of ten kilometers or so.  Alas, we have no laboratories in which to create such objects.  Fortunately, nature has been kind enough to provide us with labs in the sky.  And we don't have to look far away, either.  We are accustomed to observe quasars and gamma ray bursts at distances measured in Gigaparsecs, but our local gravitational laboratories are a million times closer than that.  There are dozens of them in our very own Milky Way galaxy, close enough to have a good peek to see what is going on.  At a distance of several kiloparsecs, a number of globular clusters have a high enough central density to let neutron stars interact and collide with other stars, often forming exotic binaries as byproducts. And at a distance of less than ten kiloparsecs, our galactic nucleus contains the mother of all gravitational laboratories, where a central massive black hole of a few million solar masses is surrounded by swarms of neutron stars, black holes and all kind of stars, some of them very unusual looking and all of them prone to collisions.  So here we are, with great gravitational laboratories, just around the corner of where we live, cosmologically speaking, and we are accumulating a treasure trove of high-precision observational data. The main problem in analyzing the data is that we came in late. Most of the experimental runs that we are currently watching were started at least millions and sometimes billions of years ago.  In order to interpret the data correctly, we have to reconstruct what has happened during the time that the experiments have been underway.  For some purposes, we can make back-of-the-envelope estimates to describe the essence of some of the main physical processes involved. For more detailed investigations, however, we have no choice but to conduct computer simulations in which we reconstruct the history of the stellar system under consideration.  In many astrophysical simulations, this task neatly breaks up into the task of simulating the individual elements, such as stars and binaries on the one hand, and the star system on the other.  In the case of dense stellar systems, however, by definition such a clear separation is not possible.  Dense stellar systems are characterized by an ecological network, where everything influences everything.  In a globular cluster, for example, dynamical interactions between passing stars can form new binaries and modify the properties and even the membership of existing binaries.  At the same time, internal changes in binaries, through mass overflow or coalescence, feed back into the energy budget of the star cluster as a whole.  In fact, most globular clusters probably have more gravitational binding energy locked up in internal binary degrees of freedom than in the bulk binding energy of the cluster as a whole.  Even a partial access to those internal `microscopic' degrees of freedom can greatly influence the `macroscopic' behavior of a star cluster as a whole.  Progress in our understanding of dense stellar systems is therefore an extreme form of two-step process.  We can gain a considerable insight in the basic processes that are at work, through dimensional analysis and back-of-the-envelope calculations, as was done very successfully in the seventies and to some extent in the eighties.  Having identified the main processes, we then suddenly faced a wall: in order to make significant further progress, we had no choice but to model the whole ecological network.  This has proved daunting: we are still only in the initial phases of living up to this challenge.  The current paper provides an outline of how we will plunge into this problem fully during the next ten years. ",
        "conclusions": ""
    },
    "0601/astro-ph0601718_arXiv.txt": {
        "abstract": "{}{To search for evidence of triggered star formation within four bright-rimmed clouds, SFO~58, SFO~68, SFO~75 and SFO~76.}{We present the results of radio continuum and molecular line observations conducted using the Mopra millimetre-wave telescope and Australia Telescope Compact Array. We use the \\mbox{$J$=1--0} transitions of $^{12}$CO, $^{13}$CO and C$^{18}$O to trace the distribution of molecular material and to study its kinematics.}{These observations reveal the presence of a dense core ($n_{\\rm{H}_2}>10^4$~cm$^{-3}$) embedded within each cloud, and the presence of a layer of hot ionised gas coincided with their bright-rims. The ionised gas has electron densities significantly higher than the critical density ($>$~25 cm$^{-3}$) above which an ionised boundary layer can form and be maintained, strongly supporting the hypothesis that these clouds are being photoionised by the nearby OB star(s). Using a simple pressure-based argument, photoionisation is shown to have a profound effect on the stability of these cores, leaving SFO~58 and SFO~68 on the edge of gravitational stability, and is also likely to have rendered SFO~75 and SFO~76 unstable to gravitational collapse. From an evaluation of the pressure balance between the ionised and molecular gas, SFO~58 and SFO~68 are identified as being in a post-pressure balance state, while SFO~75 and SFO~76 are more likely to be in a pre-pressure balance state. We find secondary evidence for the presence of ongoing star formation within SFO~58 and SFO~68, such as molecular outflows, OH, H$_2$O and methanol masers, and identify a potential embedded UC HII region, but find no evidence for any ongoing star formation within SFO~75 and SFO~76.}{Our results are consistent with the star formation within SFO~58 and SFO~68 having been triggered by the radiatively driven implosion of these clouds.} ",
        "introduction": "From the moment they turn on OB stars begin to drive an ionisation front into the surrounding molecular material, photo-evaporating and dissipating the molecular cloud from which they have formed. The rapidly expanding HII region leads to the formation of a dense shell of neutral gas, swept up in front of the ionisation front. These dense shells, and dense neutral clumps of material, surrounding evolved HII regions have long been suspected to be regions where star formation could have been triggered \\mbox{(\\citealt{elmegreen1977,sandford1982})}. Bright-Rimmed Clouds (BRCs) are small molecular clouds located on the edges of evolved HII regions and are considered, due to their relatively simple geometry and isolation within HII regions, to be ideal laboratories in which to study the physical processes involved in triggered star formation.  The photoionisation of the BRCs surface layers by UV photons from nearby OB stars leads to the formation of a layer of hot ionised gas, known as an \\emph{Ionised Boundary Layer} (IBL), which surrounds the rim of the molecular cloud. The hot ionised gas streams off the surface of the cloud into the low density HII regions, resulting in a continuous mass loss by the cloud, known as a \\emph{photo-evaporative flow} (\\citealt{megeath1997}). Within the IBL the incoming ionising photon flux is balanced by recombination, with only a small fraction of ionising photons penetrating the IBL to ionise new material (\\citealt{lefloch1994}), which replenishes the ionised material within the IBL lost to the photo-evaporative flow. If the IBL is over-pressured with respect to the molecular gas within the BRC, shocks are driven into the molecular gas, resulting in the compression of the cloud, and can lead to the formation of dense cores which are then triggered to collapse by the same (or a subsequent) shock front (\\citealt{elmegreen1992}). The propagating shock front may also serve to trigger the collapse of pre-existing dense cores, thus leading to the creation of a new generation of stars. This method of triggered star formation is known as \\emph{Radiative--Driven Implosion} (RDI) and may be responsible for the production of hundreds of stars in each HII region (\\citealt{ogura2002}), and perhaps even contributing up to $\\sim 15$ \\% of the low-to-intermediate mass IMF (\\citealt{sugitani2000}).  Shocks continue to be driven into the cloud, compressing the molecular material until the internal density and pressure is balanced with the pressure of the IBL; after which the shock fronts dissipate and the cloud is considered to be in a quasi-steady state known as the cometary stage (\\citealt{bertoldi1990,lefloch1994}). Once equilibrium is reached the ionisation front is unable to have any further influence on the internal dynamics of the cloud, but continues to propagate into the cloud photo-evaporating the molecular gas, which streams into the HII region, eroding the cloud and accelerating it radially away from the ionising star via the \\emph{rocket effect} (\\citealt{oort1955}). The mass loss resulting from photoionisation ultimately leads to the destruction of the cloud on a timescale of several million years.  A search of the IRAS point source catalogue, correlated with the Sharpless HII region catalogue (\\citealt{sharpless1959}) and the ESO(R) Southern Hemisphere Atlas resulted in a total of 89 BRCs being identified with an associated IRAS point source, 44 in the northern and 45 in the southern sky (\\citealt{sugitani1991, sugitani1994}; collectively known as the SFO catalogue). The association of an IRAS point source located within these BRCs suggests that these clouds might contain embedded protostars. Several comprehensive studies of individual BRCs have been reported (e.g. \\citealt{lefloch1997, megeath1997,yamaguchi1999,codella2001, devries2002,dobashi2002,thompson2004a}), all of which have confirmed their association with protostellar cores. However, the question of star formation occurring more widely in these objects still remains unclear, and evidence for triggered star formation within the small number of BRCs so far observed remains circumstantial and inconclusive. To address these issues we are currently conducting a complete census of the SFO catalogue; here we report the results of a detailed study of four southern BRCs.  In an earlier paper (Thompson, Urquhart \\& White, 2004b; hereafter Paper~I) we reported the results of a relatively low angular resolution radio continuum survey of the 45 southern BRCs taken from the SFO catalogue. In that survey we detected radio continuum emission toward eighteen BRCs. In each case the radio emission was found to correlate extremely well with both the morphology of the optical bright rim seen in the DSS~(R band) image and the Photo-Dominated Region (PDR) traced by the Midcourse Space eXperiment (MSX)\\footnote{Available from the NASA/IPAC Infrared Processing and Analysis Center and NASA/IPAC Infrared Space Archive both held at http://www.ipac.caltech.edu.} 8~$\\mu$m image (\\citealt{price2001}), consistent with the hypothesis  that these eighteen BRCs are being photoionised by the nearby OB star(s).   In this paper we present high resolution molecular line and radio continuum observations toward four clouds (i.e. SFO~58, SFO~68, SFO~75 and SFO~76) from the eighteen photoionised clouds identified in Paper~I which displayed the best correlation between the optical, MIR and radio emission and appeared to possess the simplest geometry.  These clouds are thus considered ideal candidates for further investigation into the applicability of the RDI star-formation mechanism. From these observations we investigate the internal and external structure of these clouds and calculate their physical parameters, which when combined with archival data will provide a comprehensive picture of star formation within these BRCs and allow us to determine whether or not it could have been triggered.  The structure of this paper is as follows: in Section~2 we briefly describe the morphologies of the BRCs and the HII regions in which they are located, including a summary of any previous observations that have been reported toward them. In Section~3 we describe our observation strategy and the molecular line and radio continuum observations, followed in Section~4 by the observational results and analyses. In Section~5 we  discuss the impact the ionisation front has had on the stability, morphology and future evolution of these clouds and try to evaluate the current state of star formation within each cloud and whether it could have been triggered. We  present a summary and our conclusions in Section~6. ",
        "conclusions": "In this paper the results of a detailed investigation of four BRCs (SFO~58, SFO~68, SFO~75 and SFO~76) are presented, including high resolution radio molecular line and continuum observations obtained with the Mopra millimetre telescope and the ATCA. The main aim is to distinguish between pre- and post-pressure balance clouds, and to evaluate to what extent the star formation within these BRCs has been influenced by the photoionisation from the nearby OB star(s). Each of the BRCs was mapped using the $^{12}$CO, $^{13}$CO and C$^{18}$O rotational transitions using the Mopra telescope. To complement the molecular line observations, high resolution radio continuum maps of all four BRCs were obtained using the ATCA.  The CO observations reveal the presence of a dense molecular core within every BRC, located behind the bright rim and, in most cases, coincident with the position of the IRAS point sources. The distribution of the $^{13}$CO emission maps suggest that the cores have a spherical structure and are centrally condensed, consistent with the presence of an embedded gravitationally bound object, such as a YSO. The H$_2$ number densities and cores masses range between 3--$8\\times10^4$~cm$^{-3}$ and 10--180~\\msun~respectively. The core temperatures ($\\sim$ 30~K) are significantly higher than expected for starless cores ($\\sim$ 10~K, \\citealt{clemens1991}), supporting evidence for the presence of an internal heating source, such as a protostar.  The high angular resolution radio observations have confirmed the presence of an IBL surrounding the rim of all four clouds, with the detected emission displaying excellent correlation with the morphology of the cloud rim seen in the optical images. The increased resolution and improved \\emph{u-v} coverage of the radio observations  has resulted in significantly higher flux densities being measured, which in turn has led to higher values for the ionising fluxes, which are now more in line with the predicted fluxes than the low angular resolution observations presented in Paper I. The electron densities are all significantly higher than the critical density of 25~cm$^{-3}$ (\\citealt{lefloch1994}) above which an IBL can form and be maintained. All these facts strongly support the hypothesis that these clouds are being photoionised by the nearby OB star(s).  CO and radio continuum data are used to evaluate the  pressure balance at the HII/molecular region interface. Comparing these values the clouds are found to fall into two categories: pre- and post-pressure balance states; SFO~75 and SFO~76 are identified as being in a pre-pressure balance state, and SFO~58 and SFO~68 are identified as being in a post-pressure balance state (taking account of the errors, see Section~\\ref{sect:pressure_balance}). We draw the following conclusions from our observations:  \\begin{enumerate}  \\item Analysis has revealed clear morphological and evolutionary differences between the pre- and post-pressure balance clouds:  \\begin{itemize} \\item The two clouds identified in this survey as being in a post-pressure balance state are also the same two identified as having a type A rim morphology and show strong evidence of ongoing high- to intermediate-mass star formation (e.g. UC HII region, masers and molecular outflows).  \\item The two clouds identified as being in a pre-pressure balance state have all been classified as type 0 clouds, and show no evidence of recent or ongoing star formation. \\end{itemize}  \\item The two classifications of rim morphologies, type 0 and type A, correspond to the 0.036 Myr and 0.126 Myr snapshots from the \\citet{lefloch1994} RDI model, where the time indicates the exposure time of a cloud to an ionising front. This is consistent with our conclusion that SFO~75 and SFO~76 have only recently been exposed to the ionisation front and that SFO~58 and SFO~68 have been exposed for a significantly longer period of time. Moreover, the morphological age predicted by the RDI models is similar to the estimated age of the star formation within SFO~58 and SFO~68, support the possibility that the star formation has been triggered.  \\item Using a simple pressure-based argument, exposure to the FUV radiation field within the HII regions is shown to have a profound effect on the stability of these cores. All of the cores were stable whilst embedded in their natal molecular clouds, however, in all cases exposure to the FUV radiation field of the HII region reduces the stability of the cores by more than a factor of two (in the case of SFO~76 by almost a factor of seven). The reduced stability leaves two clouds on the edge of being unstable to gravitational collapse, and two clouds that have masses at least a factor of two greater than the pressurised virial masses. Analysis of the pre-exposure stability indicates that it is possible that the core embedded within SFO~75 was unstable to gravitational collapse prior to being exposed to the ionisation front. From this analysis we conclude that the cores within SFO~75 and SFO~76 are both unstable to gravitational collapse.  \\item The radio continuum observations toward SFO~58 reveal the presence of an embedded compact radio source within the optical boundary of the bright rim. The radio source is offset by 1\\arcmin~from the position of the IRAS point source, but correlates extremely well with the position of the peak of the molecular core, both of which are located at the focus of the bright rim, suggesting that the radio source is associated with the cloud. The size ($<$~0.06 pc) and integrated radio flux are all consistent with the presence of an UC HII region embedded within the molecular core. The radio flux of this source is consistent with the presence of an UC HII region excited by a single ZAMS B2--B3 star. Inspection of the core-averaged $^{12}$CO spectrum reveals evidence of a substantial blue wing, possibly indicating the the presence of a molecular outflow. This is the first tentative evidence for ongoing star formation within SFO~58.  \\item The physical sizes and masses of the molecular cores are typically larger than a single star might be expected to form from. The IRAS luminosities are generally much higher than the bolometric luminosities typical for individual Class 0 and Class I stars (\\citealt{andre1993,chandler2000}). Additionally,  evidence that two BRCs are active high-mass star forming regions is presented in this paper, which form exclusively in clusters. It is therefore highly likely that the presence of the IRAS point source indicates the presence of multiple protostellar systems rather than a single protostar. Higher resolution molecular line  observations are required to investigate the possible multiplicity of these sources.  \\end{enumerate}"
    },
    "0601/astro-ph0601554_arXiv.txt": {
        "abstract": "Continuous-time, random-walk Monte Carlo simulations of H$_{2}$ formation on grains have been performed for surfaces that are stochastically heated by photons. We have assumed diffuse cloud conditions and used a variety of grains of varying roughness and size based on olivine. The simulations were performed at different optical depths.  We  confirmed that small grains (\\(r\\leq 0.02\\) \\(\\mu\\)m) have low modal temperatures with strong fluctuations, which have a large effect on the efficiency of the formation of molecular hydrogen. The grain size distribution highly favours small grains and therefore H$_2$ formation on these particles makes a large contribution to the overall formation rate for all but the roughest surfaces. We find that at $A_{V}=0$ only the roughest surfaces can produce the required amount of molecular hydrogen, but by $A_{V}=1$, smoother surfaces are possible alternatives. {Use of a larger value for the evaporation energy of atomic hydrogen, but one still consistent with experiment, allows smoother surfaces to produce more H$_{2}$.} ",
        "introduction": "Although molecular hydrogen is abundant in diffuse and dense regions of the neutral interstellar medium (ISM), the detailed mechanism of its formation is still not completely clear. At the low gas-phase temperatures typical of the interstellar medium, a gas-phase formation is not possible, and H$_{2}$ is therefore formed on the surfaces of grains. An interpretation of experimental studies \\citep{Pirronello:1997,Pirronello:1997a,Pirronello:1999} indicates that the surface temperature range over which efficient H$_{2}$ formation occurs is very small for olivine (6-10 K) and for amorphous carbon (13-17 K) \\citep{Katz:1999} under diffuse-cloud conditions. The grain temperature for midsize grains in unshielded regions is however around 20 K. This means that molecular hydrogen cannot be formed efficiently on such grains assuming that the olivine surfaces used in the laboratory are a realistic representation of interstellar grains. The analysis of \\cite{Katz:1999} considered, however, only a single barrier for diffusion between binding sites and a single low binding energy between physisorbed H atoms and the surface. In previous papers \\citep{Chang:2005, Cuppen:H2}, we showed that the introduction of sites of different binding energy leads to an increase of the temperature range for efficient H$_2$ formation. The different sites are the result of a difference in local environment either due to the amorphous character of the grain or to the  topological structure of the grain; i.e., surface roughness. These results are in agreement with the findings of \\cite{Cazaux:2004}, who introduced extra chemisorption sites and made a distinction between H and D atoms. They found a larger temperature range depending on the width and the height of the barrier between sites of physisorption and chemisorption.  In this paper, we will study the formation of molecular hydrogen using similar  continuous-time, random-walk Monte Carlo simulations as in our previous papers \\citep{Chang:2005, Cuppen:H2} while grains are heated by the interstellar radiation field. The photons hitting a grain give it a short heat impulse, resulting in temperature fluctuations especially for small grains. These temperature fluctuations have received considerable attention in the literature \\citep{Greenberg:1974,Purcell:1976,Aannestad:1979,Tabak:1987,Draine:2001}, where they are discussed with a varying degree of detail. While most of the studies simply use the heat capacity to calculate the temperature changes, \\cite{Draine:2001} looked at the heating and cooling of small silicon clusters and PAHs by considering the infrared emission of the grains, taking into account the infrared bands of the grain material. Since our main goal is to study molecular hydrogen formation, we will use a relatively simple model for the heating and cooling. As we will show, this model is capable of reproducing the detailed method to a large extent.  Since both the heating and the recombination of H atoms into H$_{2}$ are stochastic processes, the Monte Carlo technique is an ideal method to study how both processes influence each other. During the simulations, the individual H atoms are followed as they land on the grain surface, undergo their random walk, and evaporate. If two atoms arrive at the same position on the grain, they react to form molecular hydrogen. As these processes occur, photons are hitting the grain, leading to a pulsed temperature increase and subsequent decrease for the smaller grains. The radiation field and the absorption and emission coefficients used in the simulations are explained in Section \\ref{rad field}. The times at which the possible events - hopping, evaporation, and deposition of H atoms, and heating of the grain - occur are randomly determined using the rates of the processes and a random number. The details of the Monte Carlo procedure are discussed in \\cite{Chang:2005}. We  briefly summarise them in Section \\ref{Model}, which also explains the heating and cooling of the grain in more detail. Section \\ref{tempfluc} contains our results for the fluctuations in temperature, while Section \\ref{temp vs eff} shows their effect on the formation efficiency of H$_2$ as well as the effects of grain temperature vs. grain size and immediate grain desorption following formation of H$_2$. Our results are put into the context of the ISM in Section \\ref{ISM}. Finally, our conclusions are discussed in Section \\ref{conclusions}.   \\begin{figure*} \\begin{center} \\includegraphics[width=0.95\\textwidth]{dist.eps} \\end{center} \\caption{(left) The absorbed photon flux in photons per second per wavelength interval and (right) the relation between the random number and the photon energy for different grain sizes, \\(r=0.01\\) and 0.1 $\\mu$m, and at different visual extinctions, \\(A_V=0\\), 0.5 and 1 (see legend, which applies to both panels).} \\label{dist} \\end{figure*} ",
        "conclusions": "\\label{conclusions} Previous calculations of the rate of molecular hydrogen formation on interstellar grain surfaces in diffuse clouds have ignored the subtle role played by radiation in determining the surface temperature and its fluctuations as functions of grain size.  As shown by earlier investigators (e.g. \\cite{Draine:2001}) and confirmed here, the temperature of grains fluctuates strongly for smaller grains ($r < 0.02$ $\\mu$m) but not for larger ones, and the modal temperature peaks for grains in the middle range of size (0.02 - 0.05 $\\mu$m). We have shown that the efficiency of formation of H$_2$ depends both on the modal surface temperature and its fluctuations.  According to the type of material and its smoothness or roughness, each surface has a range of temperature where the efficiency is high.  For olivine grains, which are considered here, a smooth surface has only a very small range of surface temperature for high efficiency  6-9 K - while a very rough surface, with half of its binding sites different from the norm, allows efficient H$_2$ production up to about 20 K.  For small grains, photons pulse the temperature above the range of efficiency, thus reducing it, although the low modal temperature aids H$_2$ production, especially for flat and moderately rough olivine grains. The net result for these two types of surfaces is an enhancement in the efficiency for smaller vs. larger grains although the overall efficiency remains low at all sizes.  The relatively high efficiency for small grains is augmented by the high overall surface area given the distribution of grain sizes, so that the production of H$_2$ is dominated by the smaller grains.  For very rough grains, on the other hand, the efficiency is greater for larger grains since fluctuations are minimal and the modal temperature is still within the range of high efficiency, while it is above this range for smoother surfaces.  Here the countervailing larger surface area for smaller grains results in a peak H$_2$ production for grains of intermediate sizes.  A calculation of the rate coefficient $R$ for hydrogen formation in diffuse clouds for which grains from a minimum radius of 0.005 $\\mu$m to a maximum radius of 0.25 $\\mu$m are included shows that in order to equal or exceed the value of $R$ deduced by \\cite{Jura:1974} to reproduce the H$_2$/H balance in diffuse clouds, one needs a very rough surface of olivine if $A_V=0$.  At a slightly higher visual extinction of 1.0, flat and moderately rough surfaces lead to values of $R$ a factor of {six} lower than the \\cite{Jura:1974} value.  Since the latter visual extinction is quite reasonable for internal portions of  diffuse interstellar clouds, it appears that the rate of formation of H$_2$ on olivine grains of widely differing topologies can possibly account for the production of this species although the rougher surfaces are better candidates.  This conclusion rests ultimately on the energy needed for evaporation of H atoms deduced by \\cite{Katz:1999} for an olivine surface; a somewhat higher value, also consistent with their data, improves the suitability of the smoother surfaces for efficient hydrogen production, especially if the extinction is greater than zero. If the effects of the radiation field are removed, on the other hand, the difference between the smoother and rougher surfaces remains large.  In our calculations, we have only considered grains {larger than 5 nm in radius} consisting of olivine,  and we have only treated a standard radiation field.  {The inclusion of smaller silicate particles is an obvious extension, as is} the inclusion of carbonaceous grains, which would also allow us to consider much smaller particles, including PAHs.  For these very small particles, temperature fluctuations will be severe, and strong H atom-surface bonds, known as chemisorption \\citep{Cazaux:2004}, must probably be considered to allow H atoms to remain on the grains at the peak temperatures. It is possible, if barriers to chemisorption are low, that the inclusion of { small silicate particles} and PAHs will increase the calculated values of the rate coefficient $R$ for diffuse clouds, and we intend to investigate this effect in the near future. On the other hand, an enhanced radiation field, as associated with photon-dominated regions, will increase the frequency of fluctuations in temperature for small grains, and likely decrease the overall rate of H$_2$ formation although the effect has yet to be studied."
    },
    "0601/astro-ph0601281_arXiv.txt": {
        "abstract": "For the evolution of density fluctuation in nonlinear cosmological dynamics, adhesion approximation (AA) is proposed as a phenomenological model, which is especially useful for describing nonlinear evolution. However, the origin of the artificial viscosity in AA is not clarified. Recently, Buchert and Dom\\'{\\i}nguez report if the velocity dispersion of the dust fluid is regarded as isotropic, it works on a principle similar to viscosity or effective pressure, and they consider isotropic velocity dispersion as the origin of the artificial viscosity in AA. They name their model the Euler-Jeans-Newton (EJN) model. In this paper, we focus on the velocity distribution in AA and the EJN model and examine the time evolution in both models. We find the behavior of AA differs from that of the EJN model, i.e., although the peculiar velocity in the EJN model oscillates, that in AA is monotonically decelerated due to viscosity without oscillation. Therefore it is hard to regard viscosity in AA as effective pressure in the EJN model. ",
        "introduction": "\\label{sec:intro}  The Lagrangian description for the cosmological fluid can be usefully applied to the structure formation scenario. This description provides a relatively accurate model even in a quasi-linear regime. Zel'dovich~\\cite{zel} proposed a linear Lagrangian approximation for dust fluid. This approximation is called the Zel'dovich approximation (ZA)~\\cite{zel,Arnold82,Shandarin89,buchert89,Paddy93,coles,saco,Jones05,Tatekawa05R,Paddy05}. ZA describes the evolution of density fluctuation better than the Eulerian approximation \\cite{munshi,sahsha,Yoshisato98,Yoshisato05}. Although ZA gives an accurate description until the quasi-linear stage, ZA cannot describe the model after the formation of caustics. In ZA, even after the formation of caustics, the fluid elements keep moving in the direction set up by the initial condition.   In order to proceed with a hydrodynamical description in which caustics do not form, a qualitative pressure gradient \\cite{Zeldovich82} and thermal velocity scatter~\\cite{Kotok87, Shandarin89} in a collisionless matter have been discussed. Similarly, adhesion approximation (AA)~\\cite{Gurbatov89} has been proposed, which is a model based on a nonlinear diffusion equation (i.e., Burgers's equation~\\cite{Burgers40}). In AA, an artificial viscosity term is added to ZA; thus we can avoid caustics formation. The problem of structure formation has been discussed from the standpoint of AA~\\cite{Shandarin89,weinberg90,nusser,Kofman92,msw}, where it is shown that the density divergence does not occur in AA and that a density distribution close to the N-body simulation can be produced. However, the origin of the viscosity in AA has not yet been clarified.  Buchert and Dom\\'{\\i}nguez~\\cite{budo} discussed the effect of velocity dispersion using the collisionless Boltzmann equation~\\cite{BT}. They argued that models of a large-scale structure should be constructed by a flow describing the average motion of a multi-stream system. Then they showed that when the velocity dispersion is regarded as small and isotropic, it produces effective pressure or viscosity terms, and they consider that the isotropic velocity dispersion corresponds to the origin of the artificial viscosity in AA. Furthermore, in consideration of kinematic theory, they derived the relation between mass density $\\rho$ and pressure $P$, i.e., an equation of state. Buchert {\\it et al.}~\\cite{bdp} also showed how the viscosity term or the effective pressure of a fluid is generated, assuming that the peculiar acceleration is parallel to the peculiar velocity. Recently Buchert and Dom\\'{\\i}nguez~\\cite{Buchert05} provided an evaluation of the current status of adhesive gravitational clustering, which includes the above discussion, and they tried to improve past models. In their paper, they named their approach the Euler-Jeans-Newton (EJN) model. On the other hand, Dom\\'{\\i}nguez~\\cite{domi00,domi0106} proposed another approach. In these papers he clarified that a hydrodynamic formulation is obtained via a spatial coarse-graining in a many-body gravitating system, and that the viscosity term in AA can be derived by the expansion of coarse-grained equations. This model is named the Small-Size Expansion (SSE) model ~\\cite{Buchert05}.  So far, with respect to the correspondence of the viscosity term with the effective pressure, and with regard to the extension of the Lagrangian description to various matter, the EJN model has been considered. Actually, Adler and Buchert~\\cite{adler} have formulated the Lagrangian perturbation theory for a barotropic fluid. Morita and Tatekawa~\\cite{moritate}, Tatekawa {\\it et al.} \\cite{Tatekawa02}, and Tatekawa~\\cite{Tatekawa05A, Tatekawa05B} solved the Lagrangian perturbation equations for a polytropic fluid. However, it is still a open problem whether AA could realize the behavior of the EJN model in the proper way. It is known that the viscosity in AA decelerates the peculiar velocity in a dense region and avoids the formation of the caustic, while the effective pressure in the EJN model also decelerates the motion of the fluid. But the fluid in the EJN model not only would decelerate but also might bounce in a dense region due to the effective pressure, which is known as the Jeans instability for cosmological fluid. In other words, we consider that AA would not realize the behavior of the EJN model.  In this paper, to compare AA and the EJN model, we especially analyze the peculiar velocity in a cylindrical collapse model. Although we already compared AA with the EJN model with respect to density fluctuation in a past paper~\\cite{Tatekawa04}, because we used an explicit method for solving partial differential equations, the accuracy of the numerical calculation in the dense region was not good. Therefore we could not observe the Jeans instability but, rather, the numerical instability. In this paper, instead of the explicit method, we apply the iterated Crank-Nicholson method~\\cite{Teukolsky}. This method resolves the difficulty of solving the resulting implicit algebraic equations in the original Crank-Nicholson method and preserves the good stability properties.  As our analyses show, although the peculiar velocity in AA decelerates due to the viscosity, the velocity does not oscillate. On the other hand, as a preliminary expectation, the peculiar velocity in the EJN model decelerates and also oscillates. Although the tendency of the evolution of the peculiar velocity would depend on the viscous parameter or the Jeans length, the behavior of AA is obviously different from that of the EJN model. We notice that while both AA and the EJN model certainly describe the quasi-nonlinear evolution well, the detail of the evolution is different. When we take a large value for the Jeans length, the fluid bounces to the outside. On the other hand, when we take a small value for the Jeans length, although the collapse of the cylindrical matter decelerates, a caustic might finally form at the center. Therefore it is problematic that the viscous term in AA is explained as an effect similar to the effective pressure term in the EJN model. For an explanation of the viscous term in AA, we will have to analyze more Lagrangian models.  This paper is organized as follows. In Sec.~\\ref{sec:Lagrangian}, we present Lagrangian perturbative solutions in the Einstein-de Sitter (E-dS) universe. First, we show linear perturbative solutions for dust fluid in Sec.~\\ref{subsec:dust}, and in Sec.~\\ref{subsec:adhesion} we mention the problem of ZA and show the solution of AA. Then we explain the EJN model in Sec.~\\ref{subsec:pressure}. In Sec.~\\ref{sec:compari} we compare the evolution of the peculiar velocity in AA with that in the EJN model. Finally in Sec.~\\ref{sec:discuss} we discuss our results and state our conclusions. ",
        "conclusions": "\\label{sec:discuss}  With respect to the distribution of the ``peculiar velocity,'' we examine the correspondence between AA and the EJN model with cylindrical symmetry. In this analysis, even if we consider linear perturbation, it is hard to describe the solution with explicit form. Therefore we carried out numerical calculation, where to avoid numerical instability, we adopted the iterated Crank-Nicholson method. From our calculation, when we take small value for the Jeans wavenumber, the ``peculiar velocity'' in the EJN model oscillates due to the pressure. For the cosmological fluid, such oscillation is known as the Jeans instability. On the other hand, the ``peculiar velocity'' in AA is decelerated due to the viscosity, where there is no oscillation. Thus we can see that the behavior of AA is different from that of the EJN model.  Furthermore, when we take a large value for the Jeans wavenumber in the EJN model, the evolution of the ``peculiar velocity'' is similar to that in AA (see Figure~\\ref{fig:Vr-g43A}). However, it is well known that in Cartesian coordinates the linear perturbative solutions for $\\gamma \\ge 4/3$ in the EJN model have a growing solution. Therefore the perturbation in the EJN model would eventually diverge. In other words, we can predict that even if we take a large value for the Jeans wavenumber in the EJN model, the radial motion of the fluid does not stop and finally collapses. Thus there also exists an essential difference between AA and the EJN  model on this point, because it is known that the formation of caustic does not occur in AA. Hence AA cannot express the EJN model with propriety, and in order to clarify the origin of the viscosity term in AA, we should consider other effects besides isotropic velocity dispersion or isotropic effective pressure.  Recently, Buchert and Dom\\'{\\i}nguez~\\cite{Buchert05} discussed adhesive gravitational clustering and tried to provide a clear explanation of the assumuptions for adhesion approximation. They applied the Eulerian and Lagrangian expansions to the non-perturbative regime and proposed a new non-perturbative approximation. When we analyze this new approach with both analytic and numerical methods, we may be able to explicate the origin of the artificial viscosity in AA. Also, as a future work, we would describe the evolution of the density fluctuation in a nonlinear regime with a semi-analytic method."
    },
    "0601/astro-ph0601138_arXiv.txt": {
        "abstract": "This paper explores resonance-driven secular evolution between a bar and dark-matter halo using N-body simulations.  We make direct comparisons to our analytic theory \\citep{Weinberg.Katz:05a} to demonstrate the great difficulty that an N-body simulation has representing these dynamics for realistic astronomical interactions. In a dark-matter halo, the bar's angular momentum is coupled to the central density cusp (if present) by the Inner Lindblad Resonance. Owing to this angular momentum transfer and self-consistent re-equilibration, strong realistic bars {\\it WILL} modify the cusp profile, lowering the central densities within about 30\\% of the bar radius in a few bar orbits.  Past results to the contrary \\citep{Sellwood:06,McMillan.Dehnen:05} may be the result of weak bars or numerical artifacts.  The magnitude depends on many factors and we illustrate the sensitivity of the response to the dark-matter profile, the bar shape and mass, and the galaxy's evolutionary history.  For example, if the bar length is comparable to the size of a central dark-matter core, the bar may exchange angular momentum without changing its pattern speed significantly.  We emphasise that this apparently simple example of secular evolution is remarkably subtle in detail and conclude that an N-body exploration of any astronomical scenario requires a deep investigation into the underlying dynamical mechanisms for that particular problem to set the necessary requirements for the simulation parameters and method (e.g. particle number and Poisson solver).  Simply put, N-body simulations do not divinely reveal truth and hence their results are not infallible.  They are unlikely to provide useful insight on their own, particularly for the study of even more complex secular processes such as the production of pseudo-bulges and disk heating. ",
        "introduction": "\\label{sec:intro}  Current theoretical work in galaxy evolution attempts to relate the epoch of galaxy formation to the state of galaxies today.  This requires understanding an approximately ten gigayear interval over which disk galaxies are in a slowly changing near equilibrium state. Although this equilibrium will be punctuated by minor mergers, the formation of bars, and the excitation of spiral structure, the disk's existence tells us that these perturbations must be relatively mild.  As in Nature itself, cosmological simulations often produce barred disks.  This has led to a cacophony of predictions about the importance and implications of bars to the overall evolution of galaxies in the presence of cuspy dark matter haloes.  To summarise, theory predicts that bars can transfer angular momentum to haloes through resonances and most groups agree that the bar does slow \\citep{Sellwood:81,Weinberg:85,Hernquist.Weinberg:92,Debattista.Sellwood:00, Sellwood:03, Athanassoula:03,Holley-Bockelmann.Weinberg.ea:05} although \\citet[hereafter VK]{Valenzuela.Klypin:03} find a more modest slow down. \\citet{Debattista.Sellwood:00} use this and the observational evidence that bars are not slow rotators to constrain the density of the present day dark-matter halo.  \\citet[hereafter WK02]{Weinberg.Katz:02} argue that the angular momentum deposited in the halo will change the halo density profile.  This latter work was criticised by \\citet{Sellwood:06} and \\citet{McMillan.Dehnen:05} who argue that the results of WK02 do not occur in there simulations.  The goal of this paper and the previous companion paper, \\citet[hereafter Paper I]{Weinberg.Katz:05a}, is an exposition of the underlying dynamical principles that govern secular evolution and their application to N-body simulations.  In many cases of interest, the secular evolution is mediated by resonances and the implicit time dependence of an evolving system affects the subsequent evolution.  We use a numerical technique in Paper I for solving the perturbation theory that allows arbitrary time dependence to be included.  In other words, the history of the system does matter and time-asymptotic results (e.g.  Landau damping or the Lynden-Bell \\& Kalnajs 1972 theory, hereafter LBK\\nocite{Lynden-Bell.Kalnajs:72}) do not give accurate results \\citep{Weinberg:04}.  Guided by the numerical requirements derived in Paper I, we perform a series of simulations to illustrate the special features of the bar-halo resonant interaction, its dependence on the properties of the halo and the bar, and the implication for N-body simulations to properly model this interaction. We hope that our results will help to clarify and reconcile some of the apparently disparate findings of other research groups (op. cit.).  Most importantly, we show that the details of the torque depend on all aspects of the interaction, the halo profile, the bar shape, bar strength, bar pattern speed and history and cannot simply be predicted by an application of the Chandrasekhar dynamical friction or LBK formula.  For example, dynamical theory predicts that a bar inside of a homogeneous core will not slow appreciably, which contradicts the naive dynamical friction analogy; we also demonstrate this using an N-body simulation below.  An N-body simulation adds two additional complications to the dynamics of secular evolution.  First, an individual particle passing through a resonance receives a perturbation that depends sensitively on its initial position in phase-space.  The correct secular evolution is the net average of many such trajectories.  In Nature, a dark matter halo most likely has $N\\rightarrow\\infty$ from the perspective of any simulation and, therefore, has no difficultly averaging over all of phase-space.  The simulation, however, must have a sufficient number of particles in the vicinity of the resonance to obtain the correct net torque.  Paper I calls the resulting requirement on the number of particles the \\emph{coverage criterion}.  Second, representation of the dark-matter and stellar components by an unnaturally small number of particles leads to fluctuations in the gravitational potential. For modern simulations, the magnitude of these fluctuations yields a very long relaxation time but the interaction region for a resonance has a very small phase-space volume.  Paper I shows that the noise is sufficient to cause orbits to random walk through resonances.  Of course, if some orbits walk out of the resonance, others walk in. However for some resonances, ILR in particular, orbits should {\\em linger} near the resonance for many rotation periods. This increases the amplitude and changes the dependence of the torque on the phase-space distribution.  The fluctuation noise prevents this lingering and in so doing changes both the amplitude of the net torque and the location of the orbits in phase space receiving the torque. Paper I shows that natural noise sources like satellites and subhaloes will not destroy the lingering orbits.  The existence of multiple regimes underlines the necessity of understanding the dynamical mechanisms \\emph{before} fully trusting the results of a simulation constructed to investigate unknown dynamics. Paper I develops two criteria, a {\\em small-scale noise} particle number criterion that treats the scale typical of gravitationally softened particles and a {\\em large-scale noise} criterion that describes scales typical of basis expansions.  For simulations in this paper, we will use a basis expansion code (also known by the Hernquist \\& Ostriker 1992 moniker \\emph{self-consistent field (SCF) code}\\nocite{Hernquist.Ostriker:92}) to solve the Poisson equation.  Our variant of this method is reviewed in \\S\\ref{sec:method} along with the details of our initial conditions and the bar perturbation.  We choose the expansion technique for three reasons: 1) it is fast, scaling as ${\\cal O}(N)$ with small overhead; 2) it restricts spatial sensitivity to the scales of interest; and 3) it facilitates direct comparison to perturbation theory.  Paper I shows that the coverage criterion and the small-scale noise criterion dominates for the bar problem considered here.  The expansion technique eliminates the small-scale noise and, therefore, in this paper we are only able to address the breakdown of the coverage criterion.  We encourage groups with particle-particle and particle-mesh codes to test their codes as outlined in Paper I for the effects of small-scale noise.  We investigate centring in \\S\\ref{sec:dipole} and concur with \\citet{Sellwood:06} and \\citet{McMillan.Dehnen:05} that a rotating quadrupole fixed to be centred on one position leads to an $m=1$ artifact in the halo.  This can be remedied by giving the bar a monopole component, i.e. a mass, which allows it to establish its own centre naturally by conserving linear momentum, removing the $m=1$ artifact.  In \\S\\ref{sec:fid}, we simulate a bar's slow down, angular momentum transfer, and the subsequent evolution of the dark matter halo using a large bar.  We show that a strong Inner Lindblad Resonance (ILR) exists in a cuspy dark-matter halo \\citep[e.g.][hereafter NFW]{Navarro.Frenk.ea:97}. The coupling at the ILR together with the self gravity of the affected orbits drives the evolution of the inner cusp profile.  A large bar decreases the particle number requirements derived in Paper I and, because our inner dark-matter halo has a scale-free (power law) profile, the same results should obtain with little dependence on the bar size for an appropriately scaled bar mass.  We use this to investigate and corroborate the predictions by decreasing and the size and mass of the bar in \\S\\ref{sec:barsize}.  Then, in \\S\\ref{sec:barshape} we investigate the dependence on bar shape and pattern speed, and in \\S\\ref{sec:profiles} investigate the dependence on the dark matter halo profile.  We compare with other published findings in \\S\\ref{sec:others} and end with a discussion and summary in \\S\\ref{sec:sum}. ",
        "conclusions": "\\label{sec:sum}  This paper, together with Paper I, has three goals: 1) to explain and clarify the dynamics of resonances in secular evolution; 2) to derive the criteria necessary for a N-body simulation to obtain these resonance dynamics; and 3) to demonstrate the physics of resonances and the applicability of these criteria in N-body simulations.  This paper emphasises the bar--halo interaction as a test case for resonant-driven secular evolution but the same techniques can be used generally.  We confirm our prediction \\citep[WK]{Weinberg.Katz:02} that angular momentum transport can drive substantial cusp evolution within $\\approx 30$\\% of the bar radius, mediated by the inner Lindblad resonance.  Our N-body simulations corroborate the particle-number predictions of Paper I and illustrate the failure of cusp evolution when the particle criteria are not met. A small number of low-order resonances dominate the angular momentum transport.  Therefore, the features and magnitude of the evolution may vary considerably with the model depending on which resonances exist and can couple effectively.  More generally, this paper underscores the importance of a prior understanding of the underlying dynamics to ensure a successful simulation.  The importance of resonances is easily motivated.  Weak but large-scale perturbations from satellites, self-excited spiral arms, and bars break the axisymmetry of a near-equilibrium galaxy.  These perturbations require resonances to redistribute angular momentum which drive the secular evolution.  Conversely, without resonances, the torque would vanish and no evolution could ever be driven by these perturbations.  For example, the existence of the bar forces an oscillation in the orbits but many of these orbits remain adiabatically invariant to the bar forcing.  However, because the bar is slowing or changing at a finite rate, a measurable number of orbits end up with broken invariants and cause a net angular momentum change. However, the change in angular momentum of these orbits varies from positive to negative with initial phase.  These larger variations cancel to leave the net positive contribution from the bar. Therefore, a simulation must have sufficient particles to pick up the net effect of this cancellation and sufficiently low noise that an orbit remains coherent over many orbital periods during the resonant interaction.  We have discussed this and several other key points to keep in mind when considering resonant interactions.  First, Paper I further elaborates how and why resonances govern the long-term evolution near-equilibrium of slowly evolving galaxies.  Without resonances, evolution would depend on astronomical sources of scattering and scattering has an {\\em extremely} long time scale for a quiescent galaxy.  Second, resonances are broadened in frequency space by both the finite life time of the bar and the evolution of the bar. Therefore, no bar resonance is ever infinitely thin in frequency space, even one with a constant pattern speed.  To first approximation in the ratio of the evolution time scale to the characteristic period, the width of the frequency broadening has no effect on the angular momentum exchange.  We demonstrate this empirically in \\S\\ref{sec:barsize} using slowing bars of different sizes and masses. This result can motivated by an oscillator analogy.  The resonance acts as a forcing whose frequency is changing with time. The rate of change for this forcing corresponds to the resonance sweeping through phase space.  An orbit near resonance has the perturbation imposed on the natural orbital motion with a beat frequency.  As the orbit approaches the resonance, this beat frequency decreases and the perturbation amplitude increases.  The phase of the beat pattern as the beat frequency passes through zero then determines the torque.  If the time to pass through the resonance is short, many orbits will be left with a torque from the bar although the change from each will be small.  If the time to pass through is long, fewer orbits will receive a net torque but those that do will receive proportionately more angular momentum.  In second-order perturbation theory, these two opposing trends precisely cancel and the torque is the same independent of waiting time.  In other words, it is the integral under the frequency ``line'' that matters, not its width.  For example, the torque transferred in an idealised perturbation with a constant pattern speed \\citep[e.g.][]{Weinberg.Katz:02} will have the same torque as those with a varying pattern speed over a few dynamical times.  There are two additional conditions necessary for this to work: the angular momentum exchange per orbit must remain linear, $|\\Delta J|/J\\ll1$, and the rate of change in pattern speed must be larger than the squared libration frequency for the resonance (the {\\em fast limit} from TW).  This simple explanation is complicated by the non-linear dynamics that plays a role when the amplitude of the perturbation is large or the bar evolves slowly.  This is the {\\em slow limit} from TW and the resulting perturbation theory is no longer second-order.  Paper I shows that this limit is important for the ILR, in particular, and angular momentum exchange in this regime is very susceptible to small-scale noise (see Paper I for details).  For strong bars, many of the low-order resonances are close to the transition between these two regimes.  These considerations motivated particle-number criteria for N-body simulations to accurately represent the resonance dynamics.  The first, the {\\em coverage} criterion, ensures that phase space density be sufficiently high that the sum of interactions with different phase accurately yields the net torque.  The second and third criteria ensures that the gravitational potential fluctuations are sufficiently small that slow-limit resonances can still occur in the simulation. We artificially divide the noise into two regimes: 1) {\\em small-scale} noise refers to fluctuations on interparticle scales or larger.  This noise is typical of particle-particle codes; 2) {\\em large-scale} noise refers to fluctuations on the scale of the inhomogeneity of the galaxy equilibrium.  This noise is typical of the basis expansion code used in this study.  Different combinations of these three criteria dominate for different astronomical scenarios and for different codes.  We find here and in Paper I that the coverage and small-scale noise criteria dominate. Since our work here uses the basis expansion, our simulations are only affected by the coverage criterion.  We show that a scale-length bar (similar to the Milky Way) in an NFW halo with $c=15$ requires $N\\gta 2\\times10^9$ particles to recover the predicted evolutionary features. We accomplish this large value of $N$ with a multimass phase-space distribution (\\S\\ref{sec:expansion}) which yields a large effective particle number where it is needed.  We strongly encourage others to investigate the importance of these criteria and the noise criteria in particular for their favourite Poisson solver.  It is possible, for example, that tree and grid codes might introduce sources of noise not considered here and in Paper I.  The shape and mass of the bar affects these estimates.  A bar with the same size and shape as our fiducial model but a smaller mass will have a smaller resonance region and, therefore, require more particles to satisfy the coverage criterion.  Paper I shows that this criterion scales approximately as $M_b^{-1/2}$ and our simulations here are in good agreement with this scaling.  Similarly, a bar with a 5:1 in-plane axis ratio and shallow or flat surface density (similar to observed strong bars) can provide more than an order of magnitude stronger coupling that a weak oval with a falling surface density.  We have illustrated these features with a suite of simulations for bars with differing axis ratios and different particle numbers.  We find that the coupling to the inner cusp appears at the bar strength predicted by our particle number criteria.  Resonance-driven secular evolution is related to dynamical friction, which is also a secular process.  Traditional Chandrasekhar dynamical friction is non-periodic and works according to the classical scattering explanation.  For quasi-periodic systems like galaxies, the correct physical picture is angular momentum transfer near resonances, as described in TW and Paper I at length.  Technically, these dynamics can be understood as the superposition of many second-order secular Hamiltonian perturbation theory problems, one for each resonance.  In the limit of a small scale perturbation, say a tiny satellite in orbit in a dark matter halo, one can show that these two views give identical results \\citep{Weinberg:86}.  So dynamical friction operates in quasi-periodic systems because of resonances.  This does not imply that one can replace the dynamics described in Paper I with Chandrasekhar's dynamical friction formula.  For example, suppose that we removed the band of phase space around some of the low-order resonances in Figure 5 in Paper I and filled it in with unresponsive particles.  The torque would vanish but the torque predicted by Chandrasekhar's formula would only drop by a negligible amount.  Our numerical experiment of a rotating bar inside of a King-model core (\\S\\ref{sec:profiles}) shows that the torque is negligible even though the dark-matter density is the same as the rapidly slowing bar in an NFW model, because too few low-order resonances are available to accept torque.  A larger bar in the same model may be able to slow through the outer resonances but, without an ILR, the profile evolution would be quite modest.  Finally, all halo profiles affect the magnitude and location of the bar--halo torque through the existence of various resonances.  For example, by eliminating the ILR and DRR resonances with a small bar-sized core in the dark-matter halo, we can all but eliminate subsequent bar slowing.  If, for example, an early bar could flatten an initially cuspy profile, the CDM bar slowing \\emph{catastrophe} \\citep{Debattista.Sellwood:00} would be eliminated.  Paper I derives analytically and demonstrates numerically that noise from the finite number of particles leads to fluctuations in the gravitational potential that can change the nature of the resonant interaction that would obtain in the limit $N\\rightarrow\\infty$.  In Nature, galaxies have a variety of substructures and the sum total influence of these may be described as {\\em noise}.  However, all noise does not have the same spatial and temporal scales as particle noise and, therefore, will have a different impact on the resonances.  For example, consider a single satellite orbiting in the halo.  We can consider its influence on the galaxy using the perturbation theory described in Paper I and simulations by the methods described in this paper.  If we add a few more satellites, the net effect can be treated in the same way as one satellite.  As long as the influence of each satellite is coherent over particle orbital times, the same resonant dynamics obtains.  As we begin to increase the number of satellites, the combined effect will appear as rapid changes in the force in space and time.  As these fluctuations in the correlation length of the gravitational field approach the size of the resonance potential and if the correlation times become smaller than the orbital times, the net effect on the resonances is the same as with small-scale noise. However, if we restrict the population of satellites to large radii, the spatial scales of the fluctuations will be large and the temporal correlations will be long; therefore, this population only weakly effects the dynamics of the inner galaxy.  As another example, molecular clouds in a disk, because they are confined to the disk, only contribute noise on scales that are small relative to the dominant halo resonances considered in Paper I.  Noise owing to the finite number of simulation particles generates noise at all scales, and because this noise source is generated by the orbits themselves, it naturally couples to the spatial and temporal noise scales in the most effective combination to affect the resonances.  In summary, all noise is {\\em not} alike; the spatial and temporal properties of the noise are just as important as its total power.  The most obvious sources of astronomical noise---molecular clouds in the disk and substructure in the halo---are not likely to have the same effect as particle noise.  The power on various scales for some particular noise source may be computed using the method described in \\citet{Weinberg:01b}.  Finally these stringent particle number requirements does not bode well for {\\it ab initio} galaxy formation simulations (e.g. Steinmetz.Navarro:XX and Governato.etal:05) to correctly capture the correct resonant dynamics.  Without doing so, results about disk sizes and rotation curves, some of the most controversial topics, are all suspect. For example, a typical such simulation that focuses on the formation of one galaxy requires about five times more particles than those required within the virial radius to place the galaxy in its proper cosmological context.  Such simulations also use equal mass particles within the virial radius.  This means that such a simulation would require at least 10 billion particles to correctly model the disk dynamics! This is just to model the bar-halo interaction. Other processes would likely require even more particles.  Such large simulations seem unlikely with our lifetime and hence astronomers will have to invent more clever approaches than pure brute force to make progress in this area of study.  This study motivates the importance of combining simulations with analytic theory to understand galaxy evolution.  To recapitulate the introduction, the bar--halo problem was chosen for its overall structural simplicity and, therefore, its amenability to both perturbation theory and N-body simulation.  To our surprise, the problem exhibited subtlety at every turn.  From perturbation theory, we find that two separate regimes apply: one that scales as the square of the perturbation strength and one that scales as the square root of the perturbation strength.  Both regimes apply to problems of astronomical interest.  From kinetic theory, we find that these regimes can be affected by relaxation for particle numbers currently used by simulations. We further show that the quantity of torque and the qualitative behaviour of the subsequent evolution depends on most aspects of the model including its history.  Because of these subtleties, N-body studies have produced varying results leading to controversy in the literature.  Moreover, researchers typically rely on N-body simulation to study astronomical scenarios that appear trickier than this one.  For example, the heating of a disk by a satellite may require proper representation of bending modes and subsequent damping through resonant coupling.  The range of frequencies and scales in this latter problem is much broader than the bar--halo interaction considered here and, therefore, likely to present similar sorts of surprises that only a combined N-body--perturbation theory comparison will reveal.  Additional multiscale problems include the satellite--halo interaction and resonant (pseudo-)bulge formation, and these are likely to require similarly intense levels of dynamical scrutiny."
    },
    "0601/astro-ph0601412_arXiv.txt": {
        "abstract": "S\\'ersic parameters characterising the density profiles of remnants formed in collision-less disc galaxy mergers are obtained; no bulge is included in our simulations. For the luminous component we find that the S\\'ersic index is  $n\\!\\in\\!(1.5,5.3)$ with $\\langle n \\rangle\\approx 3 \\pm 1$ and an effective radius of $R_{\\rm e}\\!\\in\\!(1.6,12.9)\\,$kpc with $\\langle R_{\\rm e} \\rangle \\approx 5\\pm 3\\,$kpc. The mean values of these quantities increases as the radial interval of fitting is reduced. A strong  correlation of $n$ with the central projected density $I_0$ is found ($n\\propto I_0^{-0.14}$) which is consistent with observations. No positive linear correlation between the size ($R_{\\rm e}$) and structure ($n$) of our remnants is found; we do not advocate the existence of this. The photometric plane (PHP)  of the luminous component ($n\\propto R_{\\rm e}^{0.05}I_0^{0.15}$) agrees well, within the uncertainties and the assumption of a constant mass-to-light ratio, with those observationally determined for ellipticals in  the $B$-band ($n\\propto  R_{\\rm e}^{0.09}I_0^{0.15}$) and for $K$-band remnants ($n\\propto  R_{\\rm e}^{0.11}I_0^{0.14}$). We found that the surface defined by S\\'ersic parameters $\\{n,R_{\\rm e},\\mu_0\\}$ in log-space is not a true plane, but a pseudo-plane with a small curvature at low values of $n$ owed to intrinsic properties of the S\\'ersic model. The dark haloes of the  remnants have a 3-dimensional S\\'ersic index of $\\langle n \\rangle \\approx 4 \\pm 0.5$ that are smaller than the ones obtained for dark haloes in $\\Lambda$CDM cosmologies $n \\approx 6\\pm 1$. A tight dark S\\'ersic ``plane'' (DSP) is also defined by the parameters of the remnants haloes with $n\\propto r_{\\rm e}^{0.07} \\rho_0^{0.10}$. We conclude that collision-less merger remnants of {\\it pure} disc galaxies have S\\'ersic properties and correlations consistent with those of observed in early-type galaxies and local remnants. It seems that a ``primordial'' bulge in spirals is not a necessary condition to form bona fide ellipticals on grounds of the  S\\'ersic properties of remnants. ",
        "introduction": "Hierarchical galaxy formation theory (e.g. Cole~et~al.~2000, De~Lucia~et~al.~2005, Bower~et~al.~2005) considers that early-type galaxies have an accretion/merger origin, as was originally suggested by Toomre (1977). Observational  (e.g. Schweizer~1998, Struck~2005, Rothberg \\& Joseph~2006, Kaviraj~et~al.~2006)  and theoretical (e.g.  Naab \\& Burkert 2003, Meza~et~al.~2005, Naab~et~al.~2005)  evidence supports this picture although several topics remain unsolved (e.g. Peebles~2002, Tantalo \\& Chiosi~2004).   Early-type galaxies show several correlations among their colours, luminosities, velocity dispersions, effective radii and surface brightness  (e.g. Baum~1990, Faber~\\&~Jackson~1976, Kormendy~1977, Djorgovski~\\& Davis~1987, Dressler~et~al.~1987, Bernardi~et~al.~2003). These correlations provide constraints to any theory of  formation and evolution of these galaxies. Furthermore, their properties are linked with the distribution of luminous and dark matter,  that would be important when comparing with models of formation of elliptical galaxies.   Observational studies [e.g. Caon, Capaccioli \\& D'Onofrio 1993 (CCD93), Graham \\& Colles~1997, Binggeli \\& Jerjen 1998, D'Onofrio~2001 (D01), Trujillo~et~al.~2004] have found that the surface brightness density profiles of early-type galaxies are better described by  a S\\'ersic (1968) $R^{1/n}$--profile than  the classical de~Vaucouleurs (1948) $R^{1/4}$--profile. The index $n$ is  directly related with the cur\\-va\\-ture and ``concentration'' of the light profile (Trujillo, Graham \\& Caon~2001).   Several observational relationships have been found between the index $n$ and, for example, the total luminosity ($L$), effective radius ($R_{\\rm e}$) and central velocity dispersion (e.g. CCD93, Prugniel \\& Simien~1997, Graham \\& Guzm\\'an~2003). Also, it has been found a  linear relation among $\\log n$, $\\log R_{\\rm e}$ and $\\mu_0$ (central brightness) termed the Photometric Plane (PHP) for early-type galaxies [e.g. Khosroshahi~et~al.~2000 (K00), Graham~2002], analogous to the Fundamental Plane (Djorgovski \\& Davis 1987, Dressler~et~al.~1987). Recently, Rothberg~\\&~Joseph (2004, RJ04) have found  that nearby merger remnants have a peak in the $n$-distribution at $n\\approx 2$ with most values in the range of $1<n<6$, and in some cases it is found that $n > 8$.   On other hand, theoretical studies of  S\\'ersic properties of merger remnants have appeared recently.  For example, G\\'onzalez-Garc\\'{\\i}a~\\&~Balcells (2005, GGB) and  Naab~\\&~Trujillo~(2005, NT) find in collision-less simulations  that bulge-less progenitors lead to  ranges of $n\\! \\in\\! (2.4,3.2)$ and $(1.2,3.1)$, respectively; when a single S\\'ersic function is used to fit the entire remnant. For progenitors with a bulge component they  obtain about the same range of S\\'ersic index, $n\\! \\in \\!(3,8)$. Since ``bona fide'' ellipticals have values $n\\gta 4$, they reach the conclusion that collision-less merger remnants of pure disc galaxies do not lead to concentrations, indicated by $n$, similar to those found in intermediate or giant elliptical galaxies (e.g. Graham~et~al.~1996).   The above findings suggest that a primordial bulge in spirals is a necessary condition to form bona fide ellipticals in the hierarchical merging scenario. However, we show below,  collision-less mergers of pure discs can cover the range of observed values of the shape parameter $n$, and can reproduce adequately other observational correlations.   S\\'ersic model in a de-projected form has been recently used to represent the dark matter distribution in $\\Lambda$CDM haloes (Navarro~et~al.~2004, Merritt~et~al.~2005, Graham~et~al.~2005, Prada~et~al.~2005), in order to have a better estimation of the inner asymptotic logarithmic derivative. A mean value of a 3D S\\'ersic index $\\approx 6$, with a scatter of $\\approx 1$, has been found in these works. So it is of interest to determine the three-dimensional S\\'ersic parameters that characterise  our remnants.   In this work, we study the structural properties of remnants as provided by fitting  a  S\\'ersic profile to their luminous and dark mass distribution. The paper has been organised as follows: in $\\S$\\ref{sec:model} we present a summary of the properties of our progenitors, some details of the simulations performed, as well as some basic characteristics of  S\\'ersic profile; both projected and deprojected. In $\\S$\\ref{sec:results} we present  distributions and correlations, in two and three-dimensions, found among the  different S\\'ersic parameters for our remnants, and compare them with observations. S\\'ersic properties of the dark haloes of the remnants are determined, some correlations presented,  and compared with those obtained in cosmological simulations. Some final comments are given in $\\S$\\ref{sec:discussion} and a summary of our conclusions. ",
        "conclusions": "\\label{sec:discussion}  The results found here, as well as those of NT and GGB, show that the  merger scenario is capable of reproducing the S\\'ersic properties of observed elliptical galaxies. This work shows, however, that the presence of a ``primordial'' bulge in the progenitors is not necessary to satisfy, for example, the observed values of the shape parameter $n$; as was suggested by  NT and GGB.  It is likely that the different results with NT and GGB have their origin on  the initial properties of the progenitors; in particular,  their dark matter distribution. We have used cuspy (NFW-type) dark haloes in contrast to those used by NT (pseudo-isothermal) and GGB (Lowered Evans) that have  a constant density core. It is probable that the higher concentration of dark matter used here, affected the distribution of luminous matter in a way to increase the index $n$ which is correlated with the luminous concentration (Trujillo~et~al.~2001). Other initial conditions of the progenitors and of the encounters such as energy and angular momentum, both intrinsic and orbital, may have played a role in the final concentration of luminous matter of the remnants, as indicated by the index $n$. A systematic study of the way different dynamical elements determine the S\\'ersic index is, however, beyond the scope of the present work.   We have shown that haloes of remnants define a tight dark S\\'ersic plane (DSP) analogous to that of the luminous matter and with less dispersion. No indication of curvature is present, at difference to what is noted in the PHP at low values of $n$. We argue here that this curvature is real and is related to an intrinsic property of a S\\'ersic profile. Consider the expression for the total luminous matter associated with a S\\'ersic profile (\\ref{eq:Lmass}). This can be written in log-space as \\begin{equation} \\log L_{\\rm T} = \\log n - 0.4\\, \\mu_0 + 2 \\log R_{\\rm e} + \\log f_2(n) \\label{eq:plane2} \\end{equation} with the ``form factor'' $f_2(n)= 2 \\pi \\Gamma(2n)/b^{2n}$. A given set of galaxies with equal $L_{\\rm T}$ and different S\\'ersic parameters would define an exact log-plane, except for the presence of the $\\log f_2(n)$ term. In the 3D S\\'ersic function (\\ref{eq:3dsersic}) the analogous form factor is $f_3(n) = 4 \\pi \\Gamma(3n)/d^{3n}$. These non-constant terms introduce a systematic change in a PHP-like expression. The importance of the form factor is larger for values $n\\lta 1$ and smaller for $n \\gta 1$; both $f_2(n)$ and $f_3(n)$ are shown in Figure~\\ref{fig:Nfactor}. Thus, the form factor of the S\\'ersic model determines the curvature observed in the PHP. This explains why no curvature is found by  La~Barbera~et~al.~(2005) whose galaxies show $n \\gta 2$, but this can be seen in dwarf ellipticals with several values of $n\\lta 1$ (K04). Also, the DSP does not show such curvature since $n \\gta 3$ (see Figure~\\ref{fig:darkies}). Furthermore, the dispersion about these  ``planes'' is determined by the luminosity or dark mass range of the galaxy  sample.   \\begin{figure} \\centering \\includegraphics[width=7cm,height=7cm]{Nfactor.eps} \\vspace{-0.3cm} \\caption{Form factor $f_{2}(n)$ in equation (\\ref{eq:plane2}) contributing to the observed curvature of the photometric ``plane'' at low values of $n$; for an {\\sl ensemble} of equal total-luminous projected mass galaxies. An analogous $f_3(n)$ for a 3D distribution is also plotted. } \\label{fig:Nfactor} \\end{figure}    It remains to study the central phase-space densities of the remnants, to see if they   are consistent with the estimates for ellipticals (e.g. Carlberg~1986), and to analyse their kinematical properties with those observed in elliptical galaxies.  We plan to study these topics in a future work. In summary, our main conclusions are as follows:  \\begin{enumerate} \\item Collision-less mergers of \\emph{pure} disc galaxies yield values and distributions of  S\\'ersic parameters consistent with those observed for bona fide ellipticals. The existence of a bulge in merging spirals does not appear to be a necessary condition on grounds of S\\'ersic properties of the remnants.  \\item The suggested positive log-linear correlation between the size ($R_{\\rm e}$) and structure ($n$) in ellipticals is not supported. However, the strong $\\log~n$--$\\mu_0$ linear correlation found in observational studies is supported by our merger simulations. On other hand, the PHP is fairly well reproduced. For these results a  constant mass-to-light ratio is assumed.  \\item The final dark haloes of remnants show values of $n\\approx 4$ lower than those found in  cosmological simulations $n\\approx 6$. The difference may be attributable to the non-equivalence outer radius, where the dynamical virial radius was used in our case to carry out the fitting by a  S\\'ersic profile. Haloes define a tight Dark S\\'ersic Plane (DSP) in three dimensions, with no indication of curvature at the level of the smaller $n$ obtained.  \\item The curvature observed in the  PHP at low values of $n$ is an intrinsic manifestation of the properties of S\\'ersic model, due to the presence of a non-constant term dependent of $n$. \\end{enumerate}"
    },
    "0601/astro-ph0601624_arXiv.txt": {
        "abstract": "We are investigating the magnetic characteristics of pre-main sequence Herbig Ae/Be stars, with the aim of (1) understanding the origin and evolution of magnetism in intermediate-mass stars, and (2) exploring the influence of magnetic fields on accretion, rotation and mass-loss at the early stages of evolution of A, B and O stars. We have begun by conducting 2 large surveys of Herbig Ae/Be stars, searching for direct evidence of photospheric magnetic fields via the longitudinal Zeeman effect. From observations obtained using FORS1 at the ESO-VLT and ESPaDOnS at the Canada-France-Hawaii Telescope, we report the confirmed detection of magnetic fields in 4 pre-main sequence A- and B-type stars, and the apparent (but as yet unconfirmed) detection of fields in 2 other such stars. We do not confirm the detection of magnetic fields in several stars reported by other authors to be magnetic: HD 139614, HD 144432 or HD 31649. One of the most evolved stars in the detected sample, HD~72106A, shows clear evidence of strong photospheric chemical peculiarity, whereas many of the other (less evolved) stars do not. The magnetic fields that we detect appear to have surface intensities of order 1 kG, seem to be structured on global scales, and appear in about 10\\% of the stars studied. Based on these properties, these magnetic stars appear to be pre-main sequence progenitors of the magnetic Ap/Bp stars. ",
        "introduction": "Ap/Bp stars represent about 5\\% of all intermediate-mass main sequence stars and are characterized by strong, globally-ordered surface magnetic fields. The fundamental origin of these magnetic fields remains fundamentally a mystery. Because A and B type stars lack an outer convective envelope and exhibit none of the characteristics we have come to associate with dynamo-generated fields, it is generally believed that the magnetic fields of Ap stars are remnants, or fossils, of the interstellar magnetic field swept up during star formation.   A natural consequence of this ``fossil field'' hypothesis is the existence of magnetic field in some pre-main sequence (PMS) stars of intermediate-mass - stars which would become Ap/Bp stars when they reach the main sequence. Furthermore, we expect that the fields of those PMS stars will exhibit similar incidence, structure and intensity to those we observe in the Ap stars. The goals of our surveys were to search for such fields in the Herbig Ae/Be (HAeBe) stars, which are PMS stars of intermediate mass, of spectral types A and B, and which are characterised spectroscopically by the presence of strong, and often variable emission lines. ",
        "conclusions": "Magnetic fields can be detected and measured in Herbig Ae/Be stars using methods of measurement currently available. From our initial survey results we have detected (and confirmed) 4 new magnetic PMS A and B stars. Unconfirmed indications of magnetic field are also obtained for two additional PMS stars. On the other hand, we are unable to confirm claims of magnetic field in 3 other Herbig Ae/Be stars.  Analysis of the magnetic field measurements leads us to conclude that the magnetic field characteristics of PMS stars of intermediate-mass are qualitatively and quantitatively similar to those of their PMS descendants: they demonstrate similar incidence (about 10\\%), similar structure (primarily dipolar), and similar intensity (about 1~kG). Although chemical peculiarity cannot be confirmed in most of the magnetic Herbig Ae/Be stars, chemical spots are observed in the most evolved magnetic star of our sample. This strongly suggests that the magnetic Herbig Ae/Be stars and the magnetic Ap/Bp stars form an evolutionary sequence - that is, that the magnetic PMS stars that we have detected are the PMS progenitors of the Ap/Bp stars."
    },
    "0601/physics0601043_arXiv.txt": {
        "abstract": "A more restrictively general stability criterion of two-dimensional inviscid parallel flow is obtained analytically. First, a sufficient criterion for stability is found as either $-\\mu_1<\\frac{U''}{U-U_s}<0$ or $0<\\frac{U''}{U-U_s}$ in the flow, where $U_s$ is the velocity at inflection point, $\\mu_1$ is the eigenvalue of Poincar\\'{e}'s problem. Second, this criterion is generalized to barotropic geophysical flows in $\\beta$ plane. Based on the criteria, the flows are are divided into different categories of stable flows, which may simplify the further investigations. And the connections between present criteria and Arnol'd's nonlinear criteria are discussed. These results extend the former criteria obtained by Rayleigh, Tollmien and Fj\\o rtoft and would intrigue future research on the mechanism of hydrodynamic instability. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601630_arXiv.txt": {
        "abstract": "Type II quasars are luminous Active Galactic Nuclei whose centers are obscured by large amounts of gas and dust. In this contribution we present 3-band HST images of nine type II quasars with redshifts $0.25<z<0.4$ selected from the Sloan Digital Sky Survey based on their emission line properties. The intrinsic luminosities of these quasars are thought to be in the range $-24>M_B>-26$, but optical obscuration implies that host galaxies can be studied unencumbered by bright nuclei. Each object has been imaged in three filters (`red', `green' and `blue') placed between the strong emission lines. The spectacular, high quality images reveal a wealth of details about the structure of the host galaxies and their environments. Most galaxies in the sample are ellipticals, but strong deviations from de Vaucouleurs profiles are found, especially in the blue band. We argue that most of these deviations are due to the light from the nucleus scattered off interstellar material in the host galaxy. This scattered component can make a significant contribution to the broad-band flux and complicates the analysis of the colors of the stellar populations in the host galaxy. This extended component can be difficult to notice in unobscured luminous quasars and may bias the results of host galaxy studies. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601589_arXiv.txt": {
        "abstract": "We have considered  some results concerning   gas and stars kinematics of nearby galaxies recently obtained on the  SAO RAS 6m telescope using the panoramic spectroscopy methods. The circumnuclear regions of the galaxies were observed with integral-field spectrograph MPFS. The large-scale ionized gas kinematics was studied with the  scanning Fabry-Perot interferometer (FPI) in the multi-mode focal reducer SCORPIO. The main attention is given for kinematically decoupled regions in the galaxies: bars, spirals, polar disks and rings. ",
        "introduction": "The circumnuclear regions within the first kiloparsec in disk galaxies turn out to be decoupled on its dynamic characteristics. Using of the technique of panoramic spectroscopy makes it possible to study in detail the differences in kinematics of the stellar and gaseous components. I briefly review following types of kinematically decoupled regions in galactic disks:  \\begin{itemize} \\item Non-circular motions caused by the dynamical effects (spirals, bars, colliding rings).  \\item Circular rotation in different planes or directions (polar rings, counter-rotating disks).  \\item Non-circular motions  caused by a violent starformation (bubbles, high-velocity clouds). \\end{itemize} ",
        "conclusions": ""
    },
    "0601/astro-ph0601010_arXiv.txt": {
        "abstract": "First galaxies formed within halos of mass $M=10^{7.5}-10^9\\msun$ at $z=30-40$ in the standard cold dark matter (CDM) universe may each display an extended hydrogen 21-cm absorption halo against the cosmic microwave background with a brightness temperature decrement of $\\delta T=-(100-150)$~mK at a radius $0.3 \\le r\\le 3.0$~comoving Mpc, corresponding to an angular size of $10-100$ arcseconds. A 21-cm tomographic survey in the redshift shell $z=30-40$ (at $35-45$MHz), which could be carried out by the next generation of radio telescopes, is expected to be able to detect millions of first galaxies and may prove exceedingly profitable in enabling (at least) four fundamental applications for cosmology and galaxy formation. First, it may yield direct information on star formation physics in first galaxies. Second, it could provide a unique and sensitive probe of small-scale power in the standard cosmological model hence physics of dark matter and inflation. Third, it would allow for an independent, perhaps ``cleaner\" characterization of interesting features on large scales in the power spectrum such as the baryonic oscillations. Finally, possibly the most secure, each 21-cm absorption halo is expected to be highly spherical and faithfully follow the Hubble flow. By applying the Alcock-Paczy\\'nski test to a significant sample of first galaxies, one may be able to determine the dark energy equation of state with an accuracy likely only limited by the accuracy with which the matter density can be  determined independently. ",
        "introduction": "It is of wide interest to detect and understand the first generation of galaxies, expected to form in the redshift range $z=30-50$ in the standard CDM universe (Spergel \\etal 2003). Extensive literatures on 21-cm properties of neutral hydrogen in the dark ages and during cosmological reionization have long focused on large-scale fluctuations of the intergalactic neutral hydrogen and global spectral features (e.g., Hogan \\& Rees 1979; Scott \\& Rees 1990). In this {\\it Letter} we point out a unique feature possessed by the first {\\it individual} galaxies of mass $10^{7.5}-10^9\\msun$ formed at $z=30-40$ --- a large hydrogen 21-cm absorption halo against the cosmic microwave background (CMB). Each 21-cm absorption halo has a size $10^{''}-100^{''}$ with a brightness temperature decrement of $\\delta T=-(100-150)$~mK at $35-45$MHz, which could serve as a visible proxy for each galaxy that otherwise may be undetectable. The next generation of radio telescopes, such as LOFAR, may be able to detect such a signal. A range of fundamental applications is potentially possible with a redshift (i.e., 21-cm tomographic) survey of the first galaxies in the redshift shell $z=30-40$, which may hold the promise to revolutionize the field of cosmology and shed illuminating light on dark matter, dark energy and inflation physics. Throughout, a standard (Wilkinson Microwave Anisotropy Probe) WMAP-normalized CDM model is used (unless indicated otherwise): $\\Omega_M=0.31$, $\\Lambda=0.69$, $\\Omega_b=0.048$, $H_0=69$km/s/Mpc, $n_s=0.99$ and $\\sigma_8=0.90$. ",
        "conclusions": "It is shown that a first galaxy hosted by a halo of mass $M=10^{7.5}-10^9\\msun$ at $z=30-40$ possesses a large 21-cm absorption halo against the CMB with a brightness temperature decrement $\\delta T=-(100-150)$~mK and an angular size of $10^{''}-100^{''}$. A 21-cm tomographic survey of galaxies in the redshift shell at $z=30-40$ may detect millions of galaxies and may yield critical information on cosmology and galaxy formation. A successful observation may need an angular resolution of $\\le 1^{''}$, a spectral resolution of $\\le 4$kHz, and a sensitivity of $\\le 10$~mK at $35-45$~MHz. LOFAR appears poised to be able to execute this unprecedented task, at least for the high end of the distribution.  At least four fundamental applications may be launched with such a survey, which could potentially revolutionize cosmological study and perhaps the field of astro-particle physics. First, it may provide unprecedented constraint on star formation physics in first galaxies, for there is a proprietary sharp feature related to the threshold halo mass for efficient atomic cooling. Second, it may provide a unique and sensitive probe of the small-scale power in the cosmological model hence physics of dark matter and inflation, by being able to, for example, constrain $n_s$ to an accuracy of $\\Delta n_s=0.01$ at a high confidence level. Constraints on the nature of dark matter particles, i.e., mass or temperature, or running of index could be still tighter. Third, clustering of galaxies that may be computed with such a survey will provide an independent set of characterizations of potentially interesting features on large scales in the power spectrum including the baryonic oscillations, which may be compared to local measurements (Eisenstein \\etal 2005) to shed light on gravitational growth and other involved processes from $z=30$ to $z=0$. Finally, the 21-cm absorption halos are expected to be highly spherical and trace the Hubble flow faithfully, and thus are ideal systems for an application of the Alcock-Paczy\\'nski test. Exceedingly accurate determinations of key cosmological parameters, in particular, the equation of state of the dark energy, may be finally realized. As an example, it does not seem excessively difficult to determine $w$ to an accuracy of $\\Delta w\\sim 0.01$, if $\\Omega_M$ has been determined to a high accuracy by different means. If achieved, it may have profound ramifications pertaining dark energy and fundamental particle physics (e.g., Upadhye, Ishak, \\& Steinhardt 2005).  If a null detection of the proposed signal is found, as it might turn out, implications may be as profound. It might be indicative of some heating and/or reionizing sources in the early universe ($z=30-200$) that precede or are largely unrelated to structure formation, possibly due to yet unknown properties of dark matter particles or dark energy. Alternatively, star formation and/or BH accretion in first galaxies may be markedly different from our current expectations."
    },
    "0601/astro-ph0601683_arXiv.txt": {
        "abstract": "We show that as many as $4000$ SNeIa may be required to detect the effect of weak lensing on their flux distribution with a high level of significance. However, if the intrinsic SNeIa magnitude dispersion is unknown one needs an even higher number of SNeIa (an order of magnitude more) to reach a similar level of statistical significance. Moreover, the ability to separate the lensing contribution from the intrinsic scatter depends sensitively on the amplitude of the latter. Using a Kolmogorov - Smirnov (K-S) test we check how the required number of SNeIa changes with level of significance. Our model incorporates a completely analytical description of weak lensing which has been tested extensively against numerical simulations. Thus, future missions such as SNAP may be able to detect non-Gaussianity at a lower significance level of $10\\%$ (through the K-S test) only if the intrinsic scatter is known from external data (e.g. from low redshift observations) whereas ALPACA with $100,000$ SNe will definitely detect non-Gaussianity with a very high confidence even if the intrinsic magnitude dispersion is not known {\\it a priori}. ",
        "introduction": "Type Ia supernovae (SNeIa) are powerful probes of the recent expansion of the universe and they provided the main contribution to the discovery of the present acceleration of the universe (Riess et al. 1998; Perlmutter et al. 1999). Indeed, SNeIa are standard candles with a small luminosity dispersion so that by measuring the flux received on the earth one can derive the luminosity distance of the source. Then, by observing many SNeIa one can measure the redshift-distance relation which provides constraints on cosmological parameters (Goobar \\& Perlmutter 1995). However, even SNeIa are not perfect candles and are affected by various sources of noise such as the magnification produced by gravitational lensing, related to the fluctuations of the matter distribution along the line of sight (Kantowski et al. 1995; Frieman 1997). Flux conservation implies that the random magnification shift is zero (Weinberg 1976) but weak lensing distortions increase the observed SNeIa magnitude scatter and lead to an extended high-luminosity tail (Wambsganss et al. 1997; Valageas 2000). For a flux-limited survey weak lensing also leads to a slight bias towards larger luminosities close to the threshold (Valageas 2000) but this plays no significant role. Then, from the deviation of the magnitude distribution of 63 high redshift SNeIa (Riess et al. 2004) from a Gaussian, Wang (2005) claimed that weak lensing effects may have been detected. Although the distortion agrees at a qualitative level with the expectation from weak lensing magnification (i.e. there are three very bright SNeIa) the statistics was too small to draw firm conclusions. In this {\\it Letter} we revisit this issue by investigating how many SNeIa are needed to detect with high confidence weak lensing effects. In sect.~\\ref{Detecting_weak_lensing} we describe how weak gravitational lensing by large scale structures affects the apparent magnitude of SNeIa. Next, assuming that the intrinsic magnitude fluctuation (including all sources of noise except lensing) is Gaussian with a known variance we discuss a Kolmogorov-Smirnov test to assess whether a sample of $N$ supernovae may be drawn from such a Gaussian. Then, in sect.~\\ref{Known_intrinsic_variance} we compute for a survey such as the proposed SNAP\\footnote{http://snap.lbl.gov} experiment (Aldering et al. 2004) at which $N_*$ a deviation from this Gaussian is detected with a high confidence level. In sect.~\\ref{Marginalizing} we generalize this procedure to the case where the intrinsic magnitude variance is unknown. We discuss the dependence of our results on the amplitude of the intrinsic SNeIa magnitude dispersion in sect.~\\ref{Dependence_on_intrinsic_variance} and we conclude in sect.~\\ref{Conclusions}. ",
        "conclusions": "\\label{Conclusions}   In this {\\it Letter} we have addressed the issue of determining the number of observed SNeIa beyond which weak lensing effects can be detected with a high confidence. For the concordance $\\Lambda$CDM cosmology, using a model of the large-scale matter distribution which has been checked against numerical simulations, we found that $4000$ SNeIa are necessary to distinguish a weak lensing signature with a significance level of $2.5\\%$ through a Kolmogorov-Smirnov test. To reach a significance level of $10\\%$ we only need $2000$ SNeIa. This procedure compares the magnitude distribution of the observed SNeIa with a Gaussian of fixed variance, assuming that the latter describes all sources of noise except for weak lensing magnification. If we consider the variance to be a free parameter (e.g. the intrinsic SNeIa magnitude dispersion or the instrumental noise are not accurately known beforehand) we find that $45,000$ supernovae are required to detect with a high confidence (at a $2.5\\%$ level) non-Gaussian signatures. Therefore, future experiments such as those planned within the Joint Dark Energy Mission will exhibit clear weak lensing signatures if the intrinsic magnitude dispersion of SNeIa is well known. However to be more confident without any {\\it a priori} knowledge of $\\sigma_{\\rm int}$ we will have to wait for surveys such as ALPACA. Of course, the possibility of detecting weak lensing effects on SNeIa magnitude distributions also implies that such gravitational lensing effects should be taken into account or used as a complementary tool to constrain cosmology (e.g., Dodelson \\& Vallinotto 2005)."
    },
    "0601/astro-ph0601156_arXiv.txt": {
        "abstract": "We propose a simple mechanical model for relativistic explosions with both forward and reverse shocks, which allows one to do fast calculations of GRB afterglow. The blast wave in the model is governed by pressures $P_F$ and $P_R$ at the forward and reverse shocks. We show that the simplest assumption $P_F=P_R$ is in general inconsistent with energy conservation law. The model is applied to GRBs with non-uniform ejecta. Such \"stratified fireballs\" are likely to emerge with a monotonic velocity profile after an internal-shock stage. We calculate the early afterglow emission expected from stratified fireballs. ",
        "introduction": "Relativistic blast waves from GRBs are believed to produce the observed afterglow emission of the bursts. The blast-wave structure is schematically shown in Figure~1. To a reasonably good approximation, the whole blast [the region between FS and RS, including the contact discontinuity (CD)] may be assumed to have a common Lorentz factor $\\gamma_s$ \\cite{Piran}. The Lorentz factors of the forward shock (FS) and reverse shock (RS) are denoted by $\\gamma_F$ and $\\gamma_R$, respectively. Pressures $P_F$ and $P_R$ behind the forward and reverse shocks are calculated from the jump conditions.   \\begin{figure}[h] \\epsfig{file=schem.eps, height=14pc, width=20pc} \\caption{An illustrative diagram of 4 regions in a blast wave. Here $n_1$ and $n_4$ are the number densities of the ambient medium and the ejecta, respectively.} \\end{figure} ",
        "conclusions": ""
    },
    "0601/hep-th0601082_arXiv.txt": {
        "abstract": "In this paper we consider classical and quantum corrections to cosmological solutions of $11$D SUGRA coming from dynamics of membrane states. We first consider the supermembrane spectrum following the approach of Russo and Tseytlin for consistent quantization. We calculate the production rate of BPS membrane bound states in a cosmological background and find that such effects are generically suppressed by the Planck scale, as expected. However, for a modified brane spectrum possessing enhanced symmetry, production can be finite and significant. We stress that this effect could not be anticipated given only a knowledge of the low-energy effective theory. Once on-shell, inclusion of these states leads to an attractive force pulling the dilaton towards a fixed point of S-duality, namely $g_s=1$. Although the SUGRA description breaks down in this regime, inclusion of the enhanced states suggests that the center of M-theory moduli space is a dynamical attractor. Morever, our results seem to suggest that string dynamics does indeed favor a vacuum near fixed points of duality. ",
        "introduction": "Low-energy descriptions of string theory generically predict the existence of a large number of scalar fields, or moduli, which are associated with the size and shape of the extra dimensions, as well as the position and orientation of any branes present in the theory. These moduli are of interest for a number of reasons. From a theoretical point of view, the different vacuum expectation values (VEVs) of these fields correspond to different choices for the string vacuum, leading to an indeterminacy of the theory. From a more phenomenological viewpoint, light scalars can have a profound effect on both the early and late universe. If not fixed, these moduli can lead to a period of inflation, modifications of fundamental constants, and violations of the equivalence principle. If fixed, one must worry about the mass scales involved and the effects of the resulting relic density on cosmological observations, e.g. on the cosmic microwave background or Big Bang nucleosynthesis.  Recently it has been argued \\cite{Vafa:2005ui} (see also \\cite{Arkani-Hamed:2006dz}) that, by taking low-energy supergravity (SUGRA) as the effective description for string theory, we may miss certain crucial aspects of the underlying theory. As an example, when attempting to understand the vacuum structure of the theory one should remember that moduli space must be of finite size in order to have a realistic theory of gravity (finite $G_N$). Other crucial aspects that we wish to address in this paper are the role of dualities and the importance of dynamics. That is, as background fields evolve the effective mass of heavy states that are outside the realm of the low-energy theory can change. In fact, near points of enhanced symmetry (ESPs), frequently associated with dualities, these additional states can become massless and therefore play a vital role in the low-energy theory. It is important that these states lie beyond the naive low-energy SUGRA description and, without an underlying knowledge of the fundamental theory, effects associated with these additional degrees of freedom (such as particle creation and radiative corrections) would be missed. Indeed, recent work suggests that including these effects can have important and interesting effects both in string theory and cosmology \\cite{Silverstein:2003hf,Kofman:2004yc,Watson:2004aq,Alishahiha:2004eh}.  In this paper we will consider cosmological solutions of $11$D SUGRA and ask what corrections, if any, result from the dynamics of the moduli (i.e., dilaton and radii).  Focusing on the case of particle production, we show that production of BPS membrane bound states, whose mass depends on the evolution of the moduli, is generally suppressed by the Planck scale.  This is expected and assures us of the validity of the low-energy effective theory.  However, in the special case of membranes that become tensionless near fixed points of duality, we find that significant production can occur.  We are unaware of an explicit construction of such states, however we anticipate their existence given ubiquitous examples in the lower-dimensional string theory case.  We find that by including enhanced states in the low-energy theory, the center of the M-theory moduli space becomes an attractor, and the evolution of the moduli tends toward fixed points of duality.  In the case of the dilaton, which can be taken to be $R_{11}$ in the $11$D theory, we find that the evolution leads to the region near $g_s=1$, i.e. strong coupling. Naively trusting the equations of motion in this region, we find that all moduli will in fact become trapped and we recover a three dimensional radiation dominated universe at late times. Although we expect additional corrections to arise near the regime of strong coupling, it is intriguing that our results seem to provide an explicit example of the belief (\\cite{Brustein:2002xf} and \\cite{Dine:1998qr}) that a string vacuum consistent with gravity and particle phenomenology should lie near fixed points of duailty.   The paper is organized as follows.  In Section {\\ref{MthStates} we review a scheme developed by Russo and Tseytlin for quantizing the supermembrane in special limits.  This will allow us to obtain the BPS spectrum of membrane bound states, which we will utilize throughout the paper. In Section \\ref{ClTrapping} we briefly discuss cosmology with these states and show that they can lead to stabilization, but we also point out that such considerations seem inconsistent, since they lie outside the scope of the effective theory.  In Section \\ref{QmTrapping} we consider the possibility of producing membrane bound states within a cosmological background of $11$D SUGRA.  After reviewing the basic formalism for calculating production, we find an expression for the production rate and find that membrane bound states suffer exponential suppression, as expected.  However, we also find that, for a special class of states possessing enhanced symmetry, significant production can result. In Section \\ref{cosmo} we consider the modified effective theory in the presence of the produced states, and determine the new evolution of the moduli.  We find an attractor behavior towards the center of moduli space, and trusting the equations in the center region leads to trapping of the moduli.  We conclude with a discussion of the limitations of our approach and future work in progress. ",
        "conclusions": "We have considered both classical and quantum corrections to low-energy M-theory coming from dynamics of BPS membrane bound states whose effective mass depends on the radii of the extra dimensions. Including such states classically leads to an attractor mechanism that fixes moduli but is inconsistent with the use of the effective field theory approach.  Insisting on an effective field theory description, we then consider quantum mechanical production of these membranes in a time-dependent background, and find the expected result that the production suffers Planckian suppression. However, we do find the possibility of significant and finite production for states that exhibit enhanced symmetry. We believe that these should correspond to non-threshold bound states of membranes and gravitational waves, which become tensionless at the eleven-dimensional self-dual point, $R_{11}=l_p$. An exact construction of such states is challenging, due to the problems of quantizing the supermembrane and the lack of an effective description of the theory in this regime.  In \\cite{Russo:1996if} it was shown that one can obtain the BPS spectrum of the supermembrane in special limits and by utilizing properties of BPS configurations. However, the enhanced states we are interested in require non-vanishing vacuum energy, which is not compatible with the $\\mathcal{N}=2$ SUSY case considered in \\cite{Russo:1996if}. Therefore, guided by the analysis of \\cite{Russo:1996if}, we conjecture the existence of the enhanced states, awaiting a more concrete construction. One possibility for their existence is the case of $\\mathcal{N}=1$ heterotic M-theory, in analogy with the enhanced gauge symmetry of heterotic strings. Another intriguing possibility is the recent conjecture that the spectrum of string / M-theory contains states that lie below the BPS bound \\cite{Arkani-Hamed:2006dz} .  Assuming that such enhanced states exist, we find that they can have a critical impact on the evolution of moduli, which would have been missed in the naive low-energy effective theory neglecting dynamics. We have found that, by including the backreaction of the membrane states produced, the radii are dynamically attracted to values near fixed points of S- and T- duality. This effect would be missed if one did not have a knowledge of (enhanced) M-theory states having masses which depend on evolving moduli. Thus, the lesson we learn is that the presence of time-dependence introduces a dynamical mass scale that must be taken into careful consideration in the effective field theory. Furthermore, it is paramount not to forget the string theory origin of the low-energy effective action.  In addition to the corrections from on-shell particle production which we have explored here, there will be radiative corrections coming from the presence of the light membrane states. In fact, it was shown in \\cite{Silverstein:2003hf} that, for the case of colliding D-branes, open strings becoming light lead to corrections to the gauge theory propagator resulting in a speed limit for the moduli. In the case we have considered here a similar story should hold, but this time it is the closed string moduli which should be slowing down (the radii of the extra dimensions). This offers a challenge, since the gauge theory interpretation of such processes is unclear. We find the possibility of speed limits for radii intriguing, and we hope to report on it shortly."
    },
    "0601/astro-ph0601427_arXiv.txt": {
        "abstract": "The non--Gaussian cold spot in the 1--year WMAP data, described in Vielva et al. and Cruz et al., is analysed in detail in the present paper. First of all, we perform a more rigorous calculation of the significance of the non--zero kurtosis detected in WMAP maps by Vielva et al. in wavelet space, mainly generated by \\emph{the Spot}. We confirm the robustness of that detection, since the probability of obtaining this deviation by chance is 0.69\\%. Afterwards, the morphology of \\emph{the Spot} is studied by applying Spherical Mexican Hat Wavelets with different ellipticities. The shape of \\emph{the Spot} is found to be almost circular. Finally, we discuss if the observed non--Gaussianity in wavelet space can arise from bad subtracted foreground residues in the WMAP maps. We show that the flat frequency dependence of \\emph{the Spot} cannot be explained by a thermal Sunyaev--Zeldovich effect. Based on our present knowledge of Galactic foreground emissions, we conclude that the significance of our detection is not affected by Galactic residues in the region of \\emph{the Spot}. Considering different Galactic foreground estimates, the probability of finding such a big cold spot in Gaussian simulations is always below 1\\%. ",
        "introduction": "The Cosmic Microwave Background (CMB) is currently the most valuable tool to study the origin of the universe. This radiation was emitted $\\approx300000$ years after the Big--Bang and is the most ancient electromagnetic radiation that we can observe. It offers us information about the conditions in which it was emitted, soon after the recombination of the atoms that made the universe transparent to photons.  The most accepted theory of the early universe is the standard inflationary scenario. The CMB presents small temperature anisotropies. According to the standard model, these anisotropies represent a homogeneous and isotropic Gaussian random field on the celestial sphere. However, non-standard inflation (Bartolo et al. 2004) and topological defect models (Turok \\& Spergel 1990 and Durrer 1999) predict different non-Gaussian features in the CMB sky.  The most precise all--sky measurement of the CMB is up to date the Wilkinson Microwave Anisotropy Probe (WMAP, Bennett et al. 2003a). The analysis of the WMAP team did not find any evidence of deviations from Gaussianity of the cosmological signal in the 1--year data (Komatsu et al. 2003). Nevertheless in the last years many scientists studied the Gaussianity of these data and found non--Gaussian signatures or asymmetries using different methods: namely low multipole alignment statistics (de Oliveira--Costa et al. 2004a, Copi et al. 2004, 2005, Schwarz et al. 2004, Land \\& Magueijo 2005a,b,c, Bielewicz et al. 2005, Slosar \\& Seljak 2004); phase correlations (Chiang et al. 2003, Coles et al. 2004); hot and cold spot analysis (Larson \\& Wandelt 2004, 2005); local curvature methods (Hansen et al. 2004, Cabella et al. 2005); correlation functions (Eriksen et al. 2004a, 2005, Tojeiro et al. 2005); structure alignment statistics (Wiaux et al. 2006); multivariate analysis (Dineen \\& Coles 2005); Minkowski functionals (Park 2004, Eriksen et al. 2004b); gradient and dispersion analyses (Chyzy et al. 2005) or wavelets (Vielva et al. 2004, Mukherjee \\& Wang 2004, Cruz et al. 2005, McEwen et al. 2005). Also upcoming analyses with new tools, such as scalar statistics on the sphere, (Monteser\\'{\\i}n et al. 2005), or steerable filters, (Wiaux et al. 2005) may find some more non--Gaussian features.  In this paper we focus on the Gaussianity studies provided by Vielva et al. (2004) and Cruz et al. (2005), hereafter V04 and C05. These are based on the Spherical Mexican Hat Wavelet (SMHW) technique which allows us to perform a multiscale analysis and to localise the non--Gaussian signatures. In V04 an excess of kurtosis was detected convolving the data with the SMHW at scales between $3\\degr$ and $5\\degr$. The deviation from Gaussianity presented an upper tail probability of 0.38\\% considering all the sky and 0.11\\% in the southern hemisphere, analysing both Galactic hemispheres independently. Performing an analysis of the spots, C05 found that this deviation was due exclusively to a cold spot at galactic coordinates ($b = -57^\\circ, l = 209^\\circ$), called \\emph{the Spot} (see Fig.~\\ref{fig:SPOT}). Assuming the Gaussian hypothesis, the upper tail probability of having such a spot at a wavelet scale of $5\\degr$ is 0.18\\%. \\begin{figure} \\begin{center} \\includegraphics[width=5cm,angle=-90]{Fig_1.ps} \\end{center} \\caption{The combined and foreground cleaned Q--V--W WMAP map after convolution with the SMHW at scale $R_{9}$. The position of \\emph{the Spot} is marked.} \\label{fig:SPOT} \\end{figure} These detections were confirmed by other papers, also based on wavelets. Mukherjee \\& Wang (2004) and McEwen et al. (2005) recalculated the kurtosis analysis of V04 obtaining similar results. Cay\\'on, Jin \\& Treaster (2005) applied higher criticism statistics to maps in wavelet space. They detect a deviation from Gaussianity, finding that some pixels of \\emph{the Spot} are responsible for it.  One of the aims of this paper is to re--analyse carefully the robustness of V04 results. In this paper 15 different wavelet scales and two estimators, kurtosis and skewness were considered. Could it be that looking at that many estimators, the deviation from Gaussianity just happened by chance?  Anyway, the origin of non-Gaussian signatures is in general still not clear and controversial. Many scientists just believe that their findings are due to a deficient subtraction of the Galactic foreground emissions, (e.g. Chiang \\& Naselsky 2004 and Tojeiro et al. 2005). Some others (e.g. Eriksen et al. 2005) argue that foregrounds are unlikely to explain their results.  Considering the non--Gaussianity in wavelet space and related to the very big and cold spot, several possibilities were discussed in V04 and C05, namely systematics, Galactic foregrounds, thermal Sunyaev--Zeldovich effect (Sunyaev \\& Zeldovich 1970, hereafter SZ effect), topological defects or gravitational effects (Rees \\& Sciama 1968, Mart{\\'\\i}nez--Gonz{\\'a}lez et al. 1990a, b).  A better knowledge of \\emph{the Spot}'s shape can help us to find out its origin. For example, \\emph{the Spot} could be explained by considering topological defect models, as suggested in C05. In this context, if \\emph{the Spot} presents circular symmetry, it could have been generated by a texture. Finding out the shape of \\emph{the Spot}, can also inform us about a hypothetical gravitational potential that could be generating the non-Gaussian emission. Tomita (2005) suggests a local second--order gravitational effect as a possible origin of \\emph{the Spot}. These possibilities will be discussed in future papers.  The easiest  explanation assuming the widely accepted Gaussian hypothesis would be a bad foreground subtraction of the data. Hence it should be the first one to be tested. V04 and C05 performed several tests to discard the foregrounds as a cause of the detection. The most important test is the frequency dependence of \\emph{the Spot}, since the Galactic foregrounds have a strong frequency dependence. The amplitude and the number of pixels of \\emph{the Spot} have been shown to have a constant frequency dependence, as expected from a cosmological origin. Nevertheless, some recent works seem to contradict these conclusions. Recently, Liu \\& Zang (2005) analysed the non--Gaussianity results obtained with spherical wavelets, concluding that the most possible source of the non--Gaussian departures, are residual foregrounds. These conclusions do not agree with the results presented in V04 and C05. Also Coles (2005) suggest that Galactic foregrounds could be behind our and other non-Gaussian features.  Other different explanations are suggested for the large scale asymmetries and the non-Gaussian behaviour found in wavelet space. Jaffe et al. (2005a) assume an anisotropic cosmological model, namely the Bianchi type VII$_h$. The best fit Bianchi model introduces an additional CMB component, which presents a swirl pattern. Looking at this pattern, a hot and a cold spot can be seen close to the center of the swirl. The cold spot of the best fit model matches \\emph{the Spot}. Subtracting the swirl pattern, the significance of the kurtosis deviation and other large scale anomalies are largely reduced. Nevertheless a later work, Jaffe et al. (2005b), seems to indicate the incompatibility of the WMAP data with extended Bianchi models including a dark energy term. Freeman et al. (2005) points out a possible incorrect dipole correction as a plausible explanation for some of the observed asymmetries. This hypothesis is under examination but since \\emph{the Spot} has a much lower scale than the dipole, we do not expect to be able to explain it in this way.  Considering the relevance of the issue and the present controversial debate, the main part of this work is dedicated to investigate the origin of non--Gaussianity found in wavelet space and, in particular a possible foreground contribution to \\emph{the Spot}. Different techniques of foreground subtraction are taken into account.  Summarising, first we will describe shortly the data and simulations in section $\\S$2; the study of the robustness of the non--Gaussian detection in wavelet space will be discussed in section $\\S$3; the morphology of \\emph{the Spot} in section $\\S$4; the foreground tests are presented in $\\S$5, and our conclusions in section $\\S$6. ",
        "conclusions": "In this paper we address the issue of the origin of non--Gaussian behaviours observed in the WMAP data. In particular, a non--zero kurtosis in the distribution of wavelet coefficients was detected by V04 at angular scales ranging from $3\\degr$ and $5\\degr$. This non--Gaussian signal is mainly generated by the presence of a very cold spot in the southern hemisphere, at Galactic coordinates $b=-57\\degr$ and $l=209\\degr$ (V04, C05). Its dimension ($\\approx8\\degr$) and temperature in wavelet space makes this spot quite exceptional compared to Gaussian CMB simulations: less than $1\\%$ of the simulations have spots with similar characteristics.  As a first step, we have verified the robustness of the deviation from Gaussianity in the kurtosis. We have performed a test taking into account the estimators used in that detection (skewness and kurtosis) and the number of consecutive wavelet scales presenting a significant deviation from Gaussianity, namely 4 considering only the southern hemisphere. The significance for the non--Gaussian detection remains still high: we obtained that only $0.69\\%$ of the simulations have equal or higher deviation of skewness or kurtosis at any four consecutive scales and in any of the Galactic hemispheres.  Afterwards we studied the morphology of \\emph{the Spot} using Elliptical Mexican Hat Wavelets on the sphere. We observed that the maximum amplification of \\emph{the Spot} temperature is obtained for almost isotropic Mexican Hat Wavelets, meaning that the shape of the underlying signal is essentially circular. This result does not discard for instance topological defects like textures or a gravitational potential with a circular shape as possible explanations.  Finally, we focus on the possible foreground contamination of \\emph{the Spot} region in the clean WMAP maps, considering the SZ effect or bad--subtracted Galactic foregrounds. The SZ effect is clearly discarded by the flat frequency dependence of \\emph{the Spot} temperature. The Galactic foregrounds case requires a more detailed analysis, since a hypothetical foreground mixing could provide a flat foreground contribution.  In wavelet space, the contribution of Galactic foregrounds in the region of \\emph{the Spot} is extremely low, at least one order of magnitude less than CMB at the Q, V and W bands. The dominating foreground is the free--free emission at all WMAP frequencies except at the W band where free--free is at the same level as dust emission and below the noise level.  If the non--Gaussian analysis is affected by unsubtracted Galactic foregrounds, we would expect that the non--Gaussian detection is more significant at frequencies where the foreground emission is more relevant. But neither the kurtosis nor the area and amplitude of \\emph{the Spot} are more significant at the Q band, i.e. the band where the Galactic contribution is higher respect to the V and W bands. In addition we obtain very similar results from CMB maps produced by completely independent foregrounds subtraction techniques (e.g. WCM and TCM).  Because of the large uncertainties in Galactic emission at microwave frequencies, we have even considered the possibility that our templates provide an important underestimate of foregrounds. Nevertheless, the possibility of having a strong and frequency independent foreground residue, which could explain the non-Gaussian nature of \\emph{the Spot}, is very unlikely.  According to our knowledge on Galactic emissions, we can conclude that there is no evidence for a relevant contribution of unsubtracted foregrounds in the region of the sky which is responsible for the non--Gaussian detection in wavelet space."
    },
    "0601/astro-ph0601611_arXiv.txt": {
        "abstract": "{ The gas-phase chemistry of dark clouds has been studied with a treatment of uncertainties caused both by errors in individual rate coefficients and uncertainties in physical conditions.  Moreover, a sensitivity analysis has been employed to attempt to determine which reactions are most important in the chemistry of individual species.  The degree of overlap between calculated errors in abundances and estimated observational errors has been used as an initial criterion for the goodness of the model and the determination of a best `chemical' age of the source.  For the well-studied sources L134N and TMC-1CP, best agreement is achieved at so-called ``early times'' of $\\approx$ 10$^{5}$ yr , in agreement with previous calculations but here put on a firmer statistical foundation. A more detailed criterion for agreement, which takes into account the degree of disagreement, is also proposed.  Poorly understood but critical classes of reactions are delineated, especially reactions between ions and polar neutrals.  Such reactions will have to be understood better before the chemistry can be made more secure. Nevertheless, the level of agreement is low enough to indicate that a static picture of physical conditions without consideration of interactions with grain surfaces  is inappropriate for a complete understanding of the chemistry. ",
        "introduction": "The chemistry of dark clouds, now known as quiescent cores,  has been studied for over thirty years \\citep{1973ApJ...185..505H,1973ApJ...183L..17W}.  Indeed, a significant fraction of the more than 130 gas-phase species detected via high-resolution spectroscopy in the interstellar and circumstellar media has been seen in such sources, which include the well-studied clouds TMC1-CP and L134N \\citep{2004MNRAS.350..323S,1998ApJ...501..207T}.  In an attempt to reproduce the abundances of these species, a large number of gas-phase chemical models have been developed that compute the variation of the gas-phase concentrations as a function of time.  For many years, most of these models were based on the pseudo-time-dependent approximation \\citep{1980ApJS...43....1P,1984ApJS...56..231L,1984ApJ...277..581L,1985MNRAS.217..507M}, in which the chemistry evolves under fixed and homogeneous conditions from some initial abundances,  consisting of  totally atomic material except for a high abundance of H$_{2}$. In the last decade, more complex models, including surface chemistry \\citep{1992ApJS...82..167H,2000MNRAS.319..837R}, heterogeneity of the source \\citep{2001A&A...369..605H}, and assorted effects of dynamics and turbulence \\citep{2000ApJ...535..256M}, have been reported.  In all of these models, the connections among gas-phase species are described by up to thousands of chemical reactions, with rates quantified by rate coefficients, most of which are poorly determined by experimental or theoretical means. Nevertheless, even the older pseudo-time-dependent models were at least partially successful in reproducing the exotic nature of the chemistry (molecular ions, radicals, and metastable isomers), the unsaturated nature of the more complex products (e.g., cyanopolyynes), and the strong deuterium fractionation.  They have been less successful, however, in detailed comparisons with observations of fractional abundances and column densities for the large number of species observed in individual sources.  In the last published comparison between a pseudo-time-dependent theory and observation for the well-studied source TMC1-CP \\citep{2004MNRAS.350..323S}, it was found that at the time of best agreement ($1 \\times 10^{5}$ yr), only about 2/3 of the observed species have calculated fractional abundances within an order-of-magnitude of the observed values.  This result, obtained with the relatively new osu.2003 network and so-called ``low-metal'' elemental abundances, is actually worse than obtained with a previous network from the Ohio group, known as the ``new standard model'' (nsm), where nearly 4/5 of the molecular abundances were reproduced to within one order of magnitude \\citep{1998ApJ...501..207T}.   The order-of-magnitude criterion used in both studies is a subjective one, however,  and no comparison involving TMC1-CP has ever been done by taking into account the rate coefficient uncertainties in the computation of the theoretical abundances.  Without any quantification of the random model error, it is difficult to conclude if observed abundances can be reproduced or not by the model. Indeed, including the estimated uncertainties in rate coefficients is the only way to define those species with abundances not reproduced by chemical networks.  With a rigorous quantitative approach to uncertainty, moreover, one can begin to comprehend the deficiencies of gas-phase pseudo-time-dependent calculations, and determine where improvement is most necessary.  If, for example,  hydrogen-rich species, which may be synthesized on grain surfaces and desorbed into the gas via non-thermal means, are the only species with abundances to be poorly reproduced, the most important correction would involve the inclusion of surface processes.  If, on the other hand, smaller species such as O$_{2}$ and H$_{2}$O, with simple chemistry, cannot be reproduced within the estimated errors, then it is at least conceivable that a far more complex physical scenario is needed in which the formation and dynamics of the quiescent core play roles in the chemistry (Bergin \\& Melnick 2005).  Whether or not to give up on the pseudo-time-dependent approximation then depends critically on our estimates of uncertainties, because there is no point in violating Occam's Razor if a more simple theory is as adequate as a more complex one.  This paper is the second application of a method developed in \\citet{Wakelam2005} to quantify the random errors of chemical models. In \\citet{Wakelam2005}, we presented the method in great detail  and showed some consequences for hot core chemistry (T$> 100$~K).  Hot cores occur during an evolutionary stage of star formation, and can have complex structures not well understood  \\citep{2004A&A...414..409N}. Dark clouds, which have not yet started to collapse in an appreciable manner,  may, on the contrary,  have relatively simple temperature and density structures. These objects are then better laboratories to test the gas-phase chemical models usually used.  In this paper, we apply the method to estimate the statistical errors as a function of time for dark cloud chemistry, which is occurring at  an H$_2$ density of $\\sim 10^4$~cm$^{-3}$ and a temperature $< 20$~K.  A previous analysis of uncertainties for mainly steady-state conditions was performed for dark clouds by \\citet{2004AstL...30..566V}, although their error was defined in a different manner.  The  remainder of our paper is organized as follows.  In Section 2, we briefly describe the chemical network and approach to uncertainties used for this study. Some general results are presented in Section 3, where we also (i) consider differences between the two major networks, and (ii) make a comparison between our uncertainty calculations and those of \\citet{2004AstL...30..566V}.  Section 4 contains a comparison between our results and observed abundances towards the two dark clouds L134N and TMC1-CP, while  Section 5 discusses the specific cases of O$_2$ and H$_2$O, and the cloud ages. We conclude the paper in Section 7. ",
        "conclusions": "In this paper, we have reported the use of our Monte Carlo approach to uncertainties in calculated abundances to study the gas-phase chemistry of dark clouds, in particular the well-studied sources TMC-1CP and L134N.  We have utilized models in which the uncertainties in both rate coefficients and physical conditions are included; the latter can be thought of as either the actual observational uncertainties or physical heterogeneities in the sources.  With a specific criterion for agreement between observation and theory in which the errors in both techniques must overlap, we find that most but not species detected in the sources are reasonably well accounted for at so-called ``early time,'' a far from novel result although one determined more rigorously than in previous approaches.  A major goal of the study has been to gain some insight as to which failing of the simple picture utilized is the more important: the lack of grain-surface chemistry or the lack of a dynamical model of the sources.  For saturated (H-rich) species, with reasonably understood syntheses on grain surfaces,  our results support earlier indications that surface chemistry is an important, probably a critical omission, but for unsaturated species, the picture remains less clear because it is difficult to determine from our sensitivity analysis whether the discrepancies can be accounted for by poorly determined rates of critical chemical reactions.  This difficulty stems from the  fact that the chemistry of unsaturated species often seems to involve contributions from large numbers of poorly determined reaction rates.  The fact that discrepancies for the calculated abundances of individual species  between two chemical networks used, our osu.2003 network and the RATE99 network, can be larger than their calculated random uncertainties indicates that more attention must be paid in the laboratory to specific classes of reactions, such as ion-polar neutral systems,  before more definitive conclusions can be reached.  Further complicating the issue are other concerns such as the proper initial make-up of the gas, what elemental abundances are reasonable, and the range over which the cosmic ray ionization rate $\\zeta$ can be varied.  Nevertheless, our study indicates strongly that gas-phase models of static sources cannot represent the complete picture.  More complex treatments, involving cycles of surface adsorption, reaction, and desorption, possibly coupled with dynamical histories of the sources, would appear to be required."
    },
    "0601/astro-ph0601082_arXiv.txt": {
        "abstract": "{We present Temperature Programmed Desorption (TPD) experiments of CO and N$_2$ ices in pure, layered and mixed morphologies at various ice ``thicknesses'' and abundance ratios as well as simultaneously taken Reflection Absorption Infrared Spectra (RAIRS) of CO. A kinetic model has been developed to constrain the binding energies of CO and N$_2$ in both pure and mixed environments and to derive the kinetics for desorption, mixing and segregation. For mixed ices N$_2$ desorption occurs in a single step whereas for layered ices it proceeds in two steps, one corresponding to N$_2$ desorption from a pure N$_2$ ice environment and one corresponding to desorption from a mixed ice environment. The latter is dominant for astrophysically relevant ice ``thicknesses''. The ratio of the binding energies, $R_{\\rm BE}$, for pure N$_2$ and CO is found to be 0.936 $\\pm$ 0.03, and to be close to 1 for mixed ice fractions.  The model is applied to astrophysically relevant conditions for cold pre-stellar cores and for protostars which start to heat their surroundings.  The importance of treating CO desorption with zeroth rather than first order kinetics is shown. The experiments also provide lower limits of 0.87 $\\pm$ 0.05 for the sticking probabilities of CO-CO, N$_2$-CO and N$_2$-N$_2$ ices at 14 K. The combined results from the desorption experiments, the kinetic model, and the sticking probability data lead to the conclusion that these solid-state processes of CO and N$_2$ are very similar under astrophysically relevant conditions. This conclusion affects the explanations for the observed anti-correlations of gaseous CO and N$_2$H$^+$ in pre-stellar and protostellar cores. ",
        "introduction": "\\label{intro}  CO and N$_2$ are two of the most abundant species in molecular clouds and therefore control the abundances of many other molecules. CO is the second most abundant molecule after H$_2$, both in the gas phase and in the solid state. Gaseous CO abundances up to 2.7$\\times$ 10$^{-4}$ with respect to H$_2$ are found in warm regions \\citep{lacy1994}, indicating that CO contains most of the carbon not locked up in refractory material. In cold clouds, CO ice absorption features are seen superposed on the spectra of background sources or embedded protostars \\citep[e.g.,][]{chiar1994,pontoppidan2003}.  The solid CO abundance varies strongly from source to source, but can be as high as $10^{-4}$ with respect to H$_2$ in the coldest cores \\citep{pontoppidan2005}. Such high abundances are consistent with indirect determinations of the amount of CO frozen out in the densest parts of pre-stellar cores based on submillimeter line and continuum data, which suggest that more than 90\\% of the CO is removed from the gas \\citep[e.g.,][]{caselli1999,tafalla2004,jorgensen2005}.  The amount of N$_2$ present in the gas and solid state is more uncertain, since N$_2$ cannot be detected directly as it lacks a permanent dipole moment. The abundance of gas phase N$_2$ is usually inferred from the presence of the daughter species N$_2$H$^+$. Early work by \\citet{womack1992} inferred gas phase N$_2$ abundances of 2--6 $\\times 10^{-6}$ with respect to H$_2$ in star-forming regions, indicating that N$_2$ contains at most 10\\% of the nitrogen abundance.  Up to an order of magnitude higher abundances were found \\citet{dishoeck1992}, suggesting that at least in some sources the transformation to molecular form is complete. More recent determinations of the N$_2$ abundance have focused on dark cores for which the physical structure is well determined from complementary data. For example, \\citet{bergin1995} and \\citet{bergin2002} find typical gas-phase N$_2$ abundances of $1-2\\times 10^{-5}$.  Indirect indications for N$_2$ freeze-out onto grains can be obtained from analysis of the millimeter N$_2$H$^+$ data, which suggest a decline in the gas-phase abundance by a at least a factor of two in the centers of dense cores \\citep{bergin2002,belloche2004}. Constraints on the amount of solid N$_2$ that might be present come from analysis of the solid CO band profile \\citep{elsila1997}. The most stringent limits indicate that the N$_2$:CO ratio must be less than 1:1, derived for sources for which both $^{12}$CO and $^{13}$CO ices have been detected \\citep{boogert2002,pontoppidan2003}.  This limit only holds for mixed ices of CO and N$_2$, not when N$_2$ ice has formed a separate layer.  The chemistries of CO, N$_2$ and their daughter products are intimately linked, even though the two molecules belong to different elemental families. This is due to the fact that CO is one of the main destroyers of N$_2$H$^+$ in the gas phase. When CO is frozen out onto the grains, N$_2$H$^+$ is enhanced, as confirmed observationally by the anti-correlation of the abundances of N$_2$H$^+$ with CO and HCO$^+$ in pre- and protostellar regions \\citep{bergin2001,tafalla2002,difranco2004,pagani2005,jorgensen2004}. This anti-correlation is often quantitatively explained by a factor of 0.65 difference in the binding energies for CO and N$_2$, allowing N$_2$ to stay in the gas phase while CO is frozen out. These models do not contain an active grain-surface chemistry, but only include freeze-out and desorption. The relative freeze-out behavior of CO and N$_2$ also affects the abundance of H$_3^+$ and its level of deuterium fractionation \\citep{roberts2002}. Indeed, observations of H$_2$D$^+$ in cold cores and in protoplanetary disks often invoke large (relative) depletions of CO and N$_2$ \\citep{ceccarelli2005}.   The above discussion clearly indicates the need for a good understanding of the processes by which CO and N$_2$ freeze-out and desorb from the grains under astrophysically relevant conditions. To describe desorption, accurate values for the binding energies and the kinetics of the process are needed. For freeze-out, the sticking probability is the main uncertainty entering the equations.  In an earlier paper \\citep[][hereafter paper I]{oberg2005}, we presented a limited set of experiments using our new ultra-high vacuum (UHV) set-up to show that the ratio of the binding energies $R_{\\rm BE}$ for CO and N$_2$ in mixed and layered ices is at least 0.923 $\\pm$ 0.003 and in many circumstances close to unity.  This result can be understood chemically by the fact that the two molecules are iso-electronic.  Indeed, the sublimation enthalpies calculated from the IUPAC accredited data for pure ices were found to be 756 $\\pm$ 5 K and 826 $\\pm$ 5 K for pure N$_2$ and CO ices respectively, giving a ratio of 0.915 \\citep{lide2002}.  This experimental ratio is much larger than the value $R_{\\rm BE}=0.65$ adopted in chemical models to explain the observational data \\citep{bergin1997,ceccarelli2005}. In an alternative approach, \\citet{flower2005} used the results from paper I and instead varied the sticking probabilities of CO and N$_2$, which were assumed to be 1 below 15 K in all previous models. They could only reproduce the observed anti-correlation of N$_2$H$^+$ and HCO$^+$ if the sticking probability for N$_2$ was lowered to 0.1 compared with 1 for all other molecules.  In this paper, we present new experiments on CO--N$_2$ ices, both in pure, layered and mixed ice morphologies with varying ice ``thicknesses'' and relative abundances. In addition to TPD, RAIRS is used to probe the mixing, segregation and desorption processes in the ices. The aim of these experiments is to understand the CO--N$_2$ ice system to an extent that the experimental desorption kinetics can be modeled and reproduced, and to subsequently use these model parameters to predict the behavior of CO and N$_2$ under astrophysically relevant conditions. The key parameters to be derived for the CO--N$_2$ ice are: i) the CO-CO, CO-N$_2$, and N$_2$-N$_2$ binding energies, ii) the desorption kinetics (i.e., the desorption rates), iii) the diffusion kinetics (i.e., the mixing and segregation rates), and iv) lower limits to the sticking probabilities.   This paper is organized as follows: Sect. \\ref{expt} focuses on the experimental procedure and choice of ice layers and mixtures, Sect. \\ref{results} presents the experimental results on desorption, Sect. \\ref{model} a kinetic model of the experimental data, and Sect. \\ref{sticking} experiments on the sticking probabilities.  Sect. \\ref{astro} discusses how the kinetic model can be applied to astrophysically relevant situations and predicts the desorption behavior of CO and N$_2$ for astrophysically relevant heating rates. In Sect. \\ref{conclusion} all important conclusions are summarized. ",
        "conclusions": "\\label{conclusion}  New experimental data have been presented for the desorption and sticking of CO-N$_2$ ice systems.  Furthermore a kinetic model has been constructed that allows for accurate simulations of the TPD experiments as well as predictions for the behavior of CO and N$_2$ ices under astrophysical conditions. The key results are: \\begin{itemize} \\item The ratio for the binding energies for N$_2$ and CO in pure ices is 0.936 $\\pm$ 0.03. For mixed ices, the ratio for the binding energies is 1.0 (see Sect. \\ref{results_model}). \\item Desorption of N$_2$ from layered ices occurs in two steps, due to mixing of N$_2$ with CO. This indicates that for astrophysically relevant ice abundances, desorption from the mixed layer dominates, with less than 50\\% of N$_2$ able to desorb prior to CO. \\item In mixed ices, segregation causes the peak temperature for N$_2$ desorption to shift to lower temperatures for higher ice thicknesses, even though most of the ice desorbs from a mixed ice environment. Since the onset of segregation is concurrent with desorption, a single broad desorption step is observed for N$_2$. For astrophysically relevant ice thicknesses, N$_2$ desorption occurs close to the CO desorption temperature. \\item The desorption kinetics for CO ice are zeroth order instead of the commonly adopted first order process, resulting in an error in the desorption timescale of 12.5\\%, with a shift to lower temperatures for the first order process. Since this corresponds to 50\\% of the difference between N$_2$ and CO desorption, it results in a comparatively large effect on the relative desorption behavior of layered ices of N$_2$ and CO. \\item The lower limits on the sticking probabilities for N$_2$ and CO are found to be the same within experimental error, 0.87 $\\pm$ 0.05 at 14~K (see Sect. \\ref{sticking}). In reality the sticking probabilities will be even closer to 1.0 for lower temperatures. \\end{itemize}  The main conclusion from this work is that the solid-state processes of CO and N$_2$ are very similar under astrophysically relevant conditions."
    },
    "0601/astro-ph0601561_arXiv.txt": {
        "abstract": "We present a study of the colors, structural properties, and star formation histories for a sample of $\\sim1600$ dwarfs over look-back times of $\\sim3$ Gyr ($z=0.002-0.25$). The sample consists of 401 distant dwarfs drawn from the Galaxy Evolution from Morphologies and SEDs (GEMS) survey, which provides high resolution {\\it Hubble Space Telescope (HST)} Advanced Camera for Surveys (ACS) images and accurate redshifts, and of 1291 dwarfs at 10--90 Mpc compiled from the Sloan Digitized Sky Survey (SDSS). The sample is complete down to an effective surface brightness of 22 mag arcsec$^{-2}$ in $z$ and includes dwarfs with $M_g=-18.5$ to $-14$ mag. Rest-frame luminosities in Johnson $UBV$ and SDSS $ugr$ filters are provided by the COMBO-17 survey and structural parameters have been determined by S\\'ersic fits. We find that the GEMS dwarfs are bluer than the SDSS dwarfs by $\\sim0.13$ mag in $g-r$, which is consistent with the color evolution over $\\sim2$ Gyr of star formation histories involving moderate starbursts and long periods of continuous star formation. The full color range of the samples cannot be reproduced by single starbursts of different masses or long periods of continuous star formation alone. Furthermore, an estimate of the mechanical luminosities needed for the gas in the GEMS dwarfs to be completely removed from the galaxies shows that a significant number of low luminosity dwarfs are susceptible to such a complete gas loss, {\\it if} they would experience a starburst. On the other hand, a large fraction of more luminous dwarfs is likely to retain their gas. We also estimate the star formation rates per unit area for the GEMS dwarfs and find good agreement with the values for local dwarfs. ",
        "introduction": "Until recently the study of dwarf galaxies has been primarily concentrated either on clusters, where a large number of dwarfs can be observed within a relatively small area of sky \\citep[e.g.,][]{bin91,tre98,dri01}, or on the Local Group, where dwarfs can be studied in great detail due to their proximity \\citep[for recent reviews see][]{mat98,ber00,gre00}. Dwarfs outside of these regimes are numerous, but widespread. Many of them are associated with giant galaxies, forming galaxy groups \\citep{kar05}, which to some extent have been studied photometrically \\citep{bre98,bre99,bre00,jer00} and kinematically \\citep{bot90,kar99}. Studies of dwarfs between clusters and groups are rather rare and are mainly composed of small samples, which have been selected more or less randomly \\citep{mak99,bar01,par02} or examine specific types of dwarfs, e.g., blue compact dwarfs \\citep[BCDs;][]{noe03,gil03}.  The evolution of dwarfs is a complex problem, where evolutionary paths may depend on a variety of external and internal factors. Our knowledge of the local volume ($<8$ Mpc) has deepened, in particular due to strong efforts in determining distances to many nearby galaxies \\citep[and references therein]{kar03}. However, it is still unclear what governs the evolution of dwarfs in low density regions and how the different morphological types form. Progress has in part been hampered by the fact that many early studies are based on small samples, suffer from small number statistics and had systematic selection biases. Headway toward constraining the evolutionary history of dwarf galaxies and identifying the fundamental physical processes involved requires, as a first step, the study of large, statistically significant samples of dwarfs, selected over a large region of the sky without systematic biases.  In a cosmological context, empirical constraints on dwarf galaxies at different look-back times are essential for testing hierarchical models of galaxy formation \\citep{whi91,ste02}. One such aspect is the possibility that star formation in low mass dark matter halos gets suppressed due to a number of reasons, including the ejection of gas out of a shallow potential well by early, possibly primordial, episodes of star formation \\citep{sil03,bur04}.  Many questions remain unresolved in simulations. Over what epochs does the suppression happen: is this a primordial process only, or do dwarfs present a few Gyrs ago also exhibit signs of blowout? What types of masses and star formation histories do dwarf galaxies exhibit {\\it empirically}, and with such empirically established star formation histories, are they expected to retain their gas? What type of descendants would dwarfs at these epochs yield, if star formation was turned off?  In this paper, we present a study of the properties of dwarf galaxies over the last 3 Gyr ($z=0.002-0.25$)\\footnote{We assume in this  paper a flat cosmology with $\\Omega_M = 1 - \\Omega_{\\Lambda} = 0.3$ and $H_{\\rm 0}$=70~km~s$^{-1}$~Mpc$^{-1}$.} drawn from GEMS \\citep[Galaxy Evolution from Morphologies and SEDs,][]{rix04} and the Sloan Digitized Sky Survey \\citep[SDSS,][]{aba04}. Our initial sample consists of 988 dwarfs from the GEMS survey in the redshift range $z\\sim0.09-0.25$ (corresponding to look-back times $T_{back}$ of 1 to 3 Gyr), and a comparison local sample of 2847 dwarfs with $z<0.02$ ($T_{back}\\approx0.3$ Gyr) from SDSS. In this paper, we concentrate on the colors, star formation rates (SFR), star formation histories, and structural properties of dwarf galaxies. In a follow-up paper (Barazza et al. in prep.) we will present a detailed morphological analysis of the dwarf sample from GEMS and also take into account their environment. Throughout the paper we will refer to the dwarf samples from GEMS and SDSS just as $dwarfs$ without specifying classes (e.g., dE, Im etc.). However, we can assume that most of them are Sm, Im, or BCD in the classification scheme of \\citet{san84}, which was also confirmed by a rough visual inspection. This assumptions is based on the fact that both surveys particularly cover low density environments, where late type dwarfs dominate \\citep{kar04}. The rest of the paper is organized as follows: in $\\S$ \\ref{sur} we give a brief description of the GEMS survey and explain the sample selections; in $\\S$ \\ref{res} follows the presentation of the comparisons performed and the resultant implications as well as a discussion of the findings; finally the summary and conclusions are given in $\\S$ \\ref{sum}. ",
        "conclusions": "\\label{sum} We present a study of the colors, structural properties, and star formation histories in a sample of $\\sim1600$ dwarfs present over the last 3 Gyr ($z=0.002-0.25$). Our sample consists of 401 dwarfs over $z=0.09-0.25$ (corresponding to $T_{back}$ of 1 to 3 Gyr) from the GEMS survey and a comparison sample of 1291 dwarfs with $z<0.02$ ($T_{back}<0.3$ Gyr) from the SDSS. Our final sample is complete down to an effective surface brightness of 22 mag arcsec$^{-2}$ in $z$ and includes dwarfs with $M_g=-18.5$ to $-14$ mag. The main results are:  The mean global color of GEMS dwarfs ($g-r=0.57$ mag) is bluer than that of the SDSS dwarfs ($g-r=0.70$ mag) by 0.13 mag. Using {\\it Starburst 99}, we find that the full range of global colors shown by SDSS and GEMS dwarfs, over look-back times of 3 Gyr, is consistent with a star formation history involving bursts of star formation, combined with periods of relatively constant SFRs or exponentially decreasing SFRs. In particular, we note that a single burst model or a history of constant star formation alone, cannot reproduce the full color range. We also find that the bluer colors typical of GEMS dwarfs can evolve by $\\sim0.1$ mag into the redder colors of SDSS dwarfs over $\\sim2$ Gyr.  We identify a population of low luminosity ($M_B=-16.1$ to $-14$ mag), blue ($\\bv<0.26$ mag) dwarfs (LLBDs) in the GEMS sample, with hardly any counterparts among local SDSS dwarfs. The very blue colors of the LLBDs are comparable to those of systems, such as I Zw 18 and SBS 1415+437. Their magnitudes and colors are consistent with a recent intermediate-to-low mass starburst. However, we have to stress the caveat that a large fraction of these objects might in fact be star forming galaxies at much higher redshifts. This is indicated by an abnormally high IR emission of some of these objects and the fact, that 10 LLBDs (out of 48) exhibit a second peak around $z\\sim1$ in their redshift probability distributions.  We performed S\\'ersic fits to the GEMS and SDSS dwarfs using GALFIT making sure that both samples are fitted using the exact same procedure and weighting schemes. Our experience shows that this approach is essential for avoiding large spurious differences (see Appendix). We find that $\\sim76\\%$ of GEMS dwarfs and $\\sim81\\%$ of SDSS dwarfs have S\\'ersic $n<2$. The majority of GEMS dwarfs have half light radii comparable to those of SDSS dwarfs.  We estimate the SFR per unit area of GEMS dwarfs using the rest-frame UV (2800\\AA) luminosity and find values in the range $5\\times10^{-3}$ to $5\\times10^{-1}$ M$_{\\sun}$ yr$^{-1}$ kpc$^{-2}$. We note that the LLBDs have magnitudes and colors consistent with a recent low-to-intermediate mass starburst, small half light radii, and resultant high SFRs per unit area.  Finally, we estimate the mechanical luminosities needed for the gas in the GEMS dwarfs to become unbound and lost (blowaway). We then compare these to maximum mechanical power that would be injected by starbursts that are consistent with the derived star formation histories of these dwarfs. We find that the luminous ($M_B=-18$ to $-16$ mag) dwarfs are likely to retain their gas and avoid blowaway. However, there are a fair number of low luminosity dwarfs ($M_B=-16$ to $-14$ mag) that are susceptible to a complete blowaway of gas, {\\it if they were} to experience a starburst."
    },
    "0601/hep-ph0601054_arXiv.txt": {
        "abstract": "I derive an upper bound on the electron neutrino component of the diffuse supernova neutrino flux from the constraint on the antineutrino component at SuperKamiokande.  The connection between antineutrino and neutrino channels is due to the similarity of the muon and tau neutrino and antineutrino fluxes produced in a supernova, and to the conversion of these species into electron neutrinos and antineutrinos inside the star.  The limit on the electron neutrino flux is $5.5 ~{\\rm cm^{-2} s^{-1}}$ above 19.3 MeV of neutrino energy, and is stronger than the direct limit from Mont Blanc by three orders of magnitude.  It represents the minimal sensitivity required at future direct searches, and is intriguingly close to the reach of the Sudbury Neutrino Observatory (SNO) and of the ICARUS experiment.  The electron neutrino flux will have a lower bound if the electron antineutrino flux is measured.  Indicatively, the first can be smaller than the second at most by a factor of 2-3 depending on the details of the neutrino spectra at production. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601031_arXiv.txt": {
        "abstract": "On 2004 August 15, we observed a fast (shorter than 10 hours) state transition in the bright black-hole transient GX 339--4 simultaneously with RossiXTE and INTEGRAL. This transition was evident both in timing and spectral properties. Combining the data from PCA, HEXTE and IBIS, we obtained good quality broad-band (3-200 keV) energy spectra before and after the transition. These spectra indicate that the hard component steepened.  Also, the  high-energy cutoff that was present at $\\sim$70 keV before the transition was not detected after the transition. This is the first time that an accurate determination of the broad-band spectrum across such a transition has been measured  on a short time scale. It shows that, although some spectral parameters do not change abruptly through the transition, the high-energy cutoff increases/disappears rather fast. These results constitute a benchmark on which to test theoretical models for the production of the hard component in these systems. ",
        "introduction": "Black-hole candidate X-ray binaries (BHCs) are known to show transitions between different spectral states since the first observations of Cyg X-1 (Tananbaum et al. 1971). When X-ray instrumentation became sufficiently sophisticated to allow detailed variability studies on short time scales, the definitions of states were refined to include fast timing properties (van der Klis 1995, 2005; McClintock \\& Remillard 2005). The number and defining properties of these states have changed with time (see e.g. Homan et al. 2001), but it is now clear that fast variability is key ingredient which needs to be considered in order to have a complete view of the states and state transitions. The spectral evolution of BHCs has recently been described in terms of the pattern in an X-ray hardness-intensity diagram (HID) (see Homan \\& Belloni 2005; Belloni et al. 2005; Belloni 2005). Original states are found to correspond to different branches/areas of a $q$-like HID pattern. Four main states are identified within this framework. Two of them correspond to the original states discovered in the seventies. The Low/Hard State (LS), observed usually at the beginning and at the end of an outburst, and the most common state for the persistent system Cyg X-1 (see e.g. Dove et al. 1998; Pottschmidt et al. 2003; Cadolle Bel et al. 2005), is identified by the presence of strong ($> 30$\\% fractional rms) band-limited noise in the power spectrum and by a hard energy spectrum. The High/Soft State (HS) shows weak variability, in the form of a few \\% fractional rms power-law component in the power spectrum, and an energy spectrum dominated by a thermal disc component, with the presence of an additional weak power-law component. The HS is usually observed, if at all, in the central intervals of an outburst; it is the most common state in the persistent systems LMC X-1 and LMC X-3 (Nowak et al. 2001; Wilms et al. 2001; Haardt et al 2001). For a comparative example of these two states, see Belloni et al. (1999). In between these two well-established states, the situation is rather complex and has led to a number of different classifications. Homan \\& Belloni (2005) identify two additional states, clearly defined by spectral/timing transitions. In the evolution of a transient, after the LS comes a transition to the Hard Intermediate State (HIMS): the energy spectrum softens as the combined result of a  steepening of the power-law component and the appearance of a thermal disc component. At the same time, the characteristic frequencies in the power spectrum increase and the total fractional rms decreases. A type-C QPO (see Casella, Belloni \\& Stella 2005) appears, also with centroid frequency increasing as the source softens (see Belloni et al. 2005). The transition to the Soft-Intermediate State is very sharp (sometimes over a few seconds, see Nespoli et al. 2003) and is marked by the disappearance of the type-C QPO and by the appearance of a type-B QPO. In the 3-20 keV range, the corresponding variations in the energy spectrum are rather minor (Belloni et al. 2005). Together with the association of this transition with the ejection of fast relativistic jets, this has led to the identification of a ``jet line'' in the HID, separating these states (Fender, Belloni \\& Gallo 2004). The jet line can be crossed more than once during an outburst (as in the case of XTE J1859+226:  Casella et al. 2004; Brocksopp et al. 2002).  While the physical nature of the soft component is commonly associated with an optically thick accretion disc, there is no consensus as to the origin of the hard power-law component, where a power law is probably a simplification of a much more complex reality. The energy spectra include also additional components, important for the physics of accretion around black holes, as emission line features  (see e.g. Reynolds \\& Nowak 2003, Miller et. al 2002, Miller et al., 2004a) and Compton reflection bumps (see e.g. George \\& Fabian 1991, Zdziarski et al. 2001, Frontera et al. 2001). The spectral slope of the hard component is around $\\sim$1.6 for the LS, 1.6-2.5 in the SIMS/HIMS and 2.5-4 for the HS. One important parameter to measure is the presence/absence of a high-energy cutoff in the spectrum, for which observations at energies $>$20 keV are necessary. It is known since a long time (Sunyaev \\& Tr\\\" umper 1979) that the LS spectrum does show such a cutoff around $\\sim$100 keV. A comparative measurement of a number of systems with CGRO/OSSE has been presented by Grove et al. (1998). Here the energy spectra can clearly be divided into two classes: the ones with a strong soft thermal component and {\\it no} evidence of a high-energy cutoff until $\\sim$1 MeV, and those with no soft component and a $\\sim$100 keV cutoff. While the second class can clearly be identified with the LS, an identification of the first class is not clear. Since both the SIMS/HIMS and the HS show a strong soft component (see e.g. Nespoli et al. 2003) and OSSE integration times are very long, different states could be mixed, making a precise identification of the source state uncertain. More recently, Zdziarski et al. (2001) and Rodr\\' \\i guez et al. (2004) measured the high-energy spectrum of GRS 1915+105 and found no direct evidence for a high-energy cutoff, but spectra which appeared to contain two components. A hybrid model was found to be necessary to interpret the spectra of Cyg X-1 (Malzac et al., 2005; Cadolle Bel et al. 2005).  The nature of the hard component in BHCs is not known, nor is it known whether the ones observed in the HS and the LS have the same physical origin. For the LS, different models include thermal comptonization (e.g. Dove et al. 1997;1998), bulk motion comptonization (e.g. Laurent \\& Titarchuk 1999), jet syncrotron plus comptonization (e.g. Markoff, Falcke \\& Fender 2001; Markoff, Nowak \\& Wilms 2005). For the HS, the situation is less clear, partly because of the limited statistics available in the soft spectra and partly because of the absence of an observed characteristic energy such as a high-energy cutoff. Possible models involve non-thermal/hybrid  comptonization (e.g. Gierli\\'nski et al. 1999; Zdziarski et al. 2001; Malzac et al. 2005; Cadolle Bel et al. 2005) and bulk-motion comptonization  (e.g. Turolla, Zane \\& Titarchuk 2002).  GX 339--4 was one of the first two BHCs for which a complete set of transitions has been observed and studied (see Miyamoto et al. 1991; Belloni et al. 1997; M\\'endez \\& van der Klis 1997). This transient black-hole candidate is known to spend long periods in outburst. Historically, it was found prevalently in a hard state, although several transitions were reported (Maejima et al. 1984; Ilovaisky et al. 1986; Miyamoto et al. 1991). From the launch of the Rossi X-Ray Timing Explorer (RXTE) until 1999 it remained detectable with the RXTE All-Sky Monitor, mostly in the Low/Hard state but with a transition to a softer state in 1998 (see Nowak, Wilms \\& Dove 1999; Wilms et al. 1999; Belloni et al. 1999; Nowak, Wilms \\& Dove 2002; Corongiu et al. 2003). In 1999, the source went into quiescence, where it was detected with Chandra BeppoSAX at low flux levels (Kong et al. 2000; Corbel et al. 2003).  After roughly one year in quiescence, a new outburst started in 2002 (Smith et al. 2002a,b; Nespoli et al. 2003; Belloni 2004; Homan et al. 2005) and ended in 2003 (Buxton \\& Bailyn 2004a). This new outburst was followed in detail through timing and color analysis (Belloni et al. 2005), which made it the prototype for the HID evolution and the definitions of the HIMS/SIMS. In particular, a very clear HIMS-SIMS transition was observed on 2003 May 17, with fast changes in the variability properties, but almost no variations in the 3-20 keV energy spectrum (Homan et al. 2005). A relativistically broadened iron emission line has been detected in the X-ray spectrum of GX 339--4, indicating the presence of a non-zero angular momentum in the black hole (Miller et al. 2004a,b). A comparative spectral analysis of all existing RXTE data from GX 339--4 is presented by Zdziarski et al. (2004).  GX 339--4 was also the first BHC showing radio/X-ray and near-IR/X-ray correlations in the LS (Hannikainen et al. 1998; Corbel et al. 2003; Markoff et al. 2003;  Homan et al. 2005). Radio observations during the 1999 HS showed clear evidence of a strong decrease of core radio emission during this state (Fender et al. 1999). During the 2002/2003, near to the transition to the VHS (see Smith et al. 2002c), a bright radio flare was observed (Fender et al. 2002), which led to the formation of a large-scale relativistic jet (Gallo et al. 2004). Fender, Belloni \\& Gallo (2004) associated this flare and subsequent ejections with the crossing of the jet line, i.e. the HIMS-SIMS transition reported by Nespoli et al. (2003).  A high mass function (5.8$\\pm$0.5 $M_\\odot$) has been measured for the system (Hynes et al. 2003), indicating strong dynamical evidence for the black-hole nature of the compact object. The distance to GX 339--4 is not well known, with a lower limit of $\\sim$6 kpc (Hynes et al. 2004; see also Zdziarski et al. 1998).  It is clear that GX 339--4 is an extremely important source for our understanding of the accretion and ejection properties of stellar-mass black holes. In February 2004, after having returned to quiescence, GX 339--4 started a new outburst (Buxton et al. 2004, Smith et al. 2004, Belloni et al. 2004,  Kuulkers  et al. 2004, Israel et al. 2004). ",
        "conclusions": "The results presented above constitute one of the first determination of the changes of the broad-band X-ray spectrum of a BHC across the HIMS-SIMS transition, which is crucial for the development and testing of theoretical models. Although the transition from the LS to the HS is a process that takes days to weeks (see e.g. Belloni et al. 2005), a sharp transition in the properties of fast time variability takes place on a much shorter time scales. We showed that this transition corresponds also to a change in the high-energy properties."
    },
    "0601/astro-ph0601207_arXiv.txt": {
        "abstract": "This paper uses the First {\\it XMM-Newton} Serendipitous Source Catalog compiled by the {\\it XMM-Newton} Science Center to identify low-$z$ X-ray selected normal galaxy candidates.  Our sample covers a total area of $\\approx \\rm 6\\, deg^2$ to the 0.5-2\\,keV limit $\\rm \\approx 10^{-15} \\, erg \\, s^{-1} \\, cm^{-2}$. A total of 23 sources are selected on the basis of low X-ray--to--optical flux ratio $\\log f_X/f_{opt}  < -2$, soft X-ray spectral properties and optical spectra, when available, consistent with stellar than AGN processes. This sample is combined with similarly selected systems from the Needles in the Haystack Survey (Georgantopoulos et al. 2005) to provide a total of 46 unique $z \\la 0.2$ X-ray detected normal galaxies, the largest low-$z$ sample yet available. This is first used to constrain the  normal galaxy $\\log N - \\log S$  at bright fluxes ($\\rm 10^{-15} -10^{-13} \\, erg \\, s^{-1} \\,  cm^{-2}$). We estimate a slope of $-1.46\\pm0.13$ for the cumulative number counts consistent with the euclidean prediction. We further combine our sample with 23 local ($z \\la 0.2$) galaxies from the {\\it Chandra} Deep Field North and South surveys to construct  the local X-ray luminosity function of normal galaxies. A Schechter form provides a good fit to the data with a break at $\\rm \\log L_\\star = 41.02_{-0.12}^{+0.14} \\rm \\, erg \\, s^{-1}$ and a slope of $\\alpha=-1.76\\pm 0.10$. Finally, for the sample of 46 systems we explore the association between X-ray luminosity and host galaxy properties, such as star-formation rate and stellar mass. We find that the $L_X$ of the emission-line systems correlates with $\\rm H\\alpha$ luminosity and 1.4\\,GHz radio power, both providing an estimate of the current star-formation rate. In the case of early type galaxies with absorption line optical spectra we use the $K$-band as proxy to stellar mass and find a correlation of the form $L_X \\propto L_K^{1.5}$. This is flatter than the $L_X - L_B$ relation for local ellipticals. This may be due to either $L_K$ providing a better proxy to galaxy mass or selection effects biasing our sample against very luminous early-type galaxies, $L_X > 10^{42} \\rm erg \\, s^{-1}$. ",
        "introduction": "The launch of the {\\it XMM-Newton} and the {\\it Chandra} missions has opened new opportunities in the study of the X-ray properties of normal galaxies. Apart from targeted observations of local systems, surveys performed by these telescopes have allowed, for the first time, study of normal galaxies at X-ray wavelengths outside the nearby Universe, thus opening the way for evolutionary investigations. Stacking analysis studies using these missions for example, have provided  constraints on the mean X-ray properties ($L_X$, hardness ratios, $L_X/L_B$) of star-forming systems over the redshift range $z=0.1-3$ (Hornschemeier et al. 2002; Nandra et al. 2002; Georgakakis et al. 2003a, b; Laird et al. 2004). The stacking analysis results from the above independent studies when combined together are consistent with luminosity evolution of the form $\\approx (1+z)^3$ at least to $z\\approx1.5$ (e.g. Georgakakis et al. 2003a, b).  The X-ray stacking analysis provides information on the mean properties of systems that are too X-ray faint  to be individually detected. A major breakthrough however, of the {\\it XMM-Newton} and the {\\it Chandra} missions has been the  {\\it detection} of normal galaxies at low and moderate redshifts. Bauer et al. (2002) and Alexander et al. (2002) studied the radio and mid-infrared properties of X-ray sources in the  Chandra Deep Field (CDF) North and argue that systems with X-ray emission dominated by  starburst activity are detected in this field out to $z\\approx1$. Hornschemeier et al. (2003) demonstrated that more quiescent normal galaxies with X-ray--to--optical flux ratio $\\log f_X / f_{opt}< -2$ and a median redshift $z \\approx 0.3$ not only constitute a non-negligible component of the X-ray source population in the CDF-North but  are also likely to outnumber AGNs at fluxes below $f_X(\\rm 0.5 - 2 \\, keV) \\approx 10^{-17} \\rm \\, erg \\, s^{-1} \\, cm^{-2}$. Norman et al. (2004) extended the Hornschemeier et al. (2003) study and identified over 100 starburst/quiescent galaxy candidates to $z\\approx1$ in the combined CDF-North and South albeit with optical spectroscopy limited to a fraction of them. These authors estimate the X-ray luminosity function of normal galaxy candidates and find evidence for X-ray evolution of the form $\\propto (1+z)^3$ to $z\\approx1$, i.e. similar to that inferred from other wavelengths (e.g. Hopkins 2004 and references therein).  Parallel to the studies above focusing on distant sources ($z \\ga 0.3$), there have also been efforts to compile X-ray selected normal galaxy samples at low redshifts $z<0.2$. This is essential to complement the deep CDF studies and to provide a comparison sample in the nearby Universe. Georgakakis et al. (2004) and Georgantopoulos, Georgakakis \\& Koulouridis (2005) used public {\\it XMM-Newton} data overlapping with the Sloan Digital Sky Survey (SDSS; Stoughton et al. 2002; York et al. 2000) to identify normal galaxies at a median redshift of about 0.07 over an $\\rm 11 \\, deg^2$ area. This dataset, the Needles in the Haystack Survey, has been used to constrain the galaxy $\\log N - \\log S$ at bright fluxes, estimate the X-ray luminosity function of absorption and emission line galaxies separately and to assess their  contribution to the diffuse X-ray background under different evolution scenarios.  A similar study using the HRI detector onboard {\\it ROSAT} has been performed by Tajer et al. (2005) aiming to determine the number counts of low X-ray--to--optical flux ratio systems ($\\log f_X/f_{opt} < -1$) at very bright fluxes ($\\rm 10^{-14} - 10^{-12} \\, erg \\, s^{-1} \\, cm^{-2}$). More  recently, Hornschemeier et al. (2005) explored the X-ray properties of $z\\approx0.1$ galaxies identified in public {\\it Chandra} pointings overlapping with the SDSS. They use this sample to investigate the link between X-ray luminosity,  host galaxy star-formation rate (SFR) and stellar mass. These relations are consistent with a picture where the X-ray emission  arises in a combination of low- and high-mass X-ray binaries as well as hot gas interstellar medium.  Despite significance progress, the number of X-ray selected normal galaxies at low-$z$ (e.g. $z \\la 0.2$) remains small hampering statistical studies. In this paper we further expand existing samples using as a resource the 1st {\\it XMM-Newton} Serendipitous Source Catalogue produced by the  {\\it XMM-Newton} Science Center. Our aim is to demonstrate the power of this unique and continuously expanding database and how it can be exploited to advance our understanding of the properties of X-ray selected normal galaxies. Among others, such a sample can provide a tight anchor point for comparison with deeper surveys probing on average higher-$z$, calibrate the relation between SFR diagnostics and X-ray luminosity and improve our understanding of the  association between host galaxy properties (e.g. stellar mass) and X-ray emission. Throughout this paper we adopt $\\rm H_{o} = 70 \\, km \\, s^{-1} \\, Mpc^{-1}$, $\\rm \\Omega_{M} =   0.3$ and $\\rm \\Omega_{\\Lambda} = 0.7$. ",
        "conclusions": "In this paper we demonstrate the power of the First {\\it XMM-Newton} Serendipitous Source Catalog for studies of X-ray selected normal galaxies. Our sample is compiled from X-ray sources detected on {\\it XMM-Newton} pointings that have  (i) EPIC-PN detector as prime instrument operated in full-frame mode, (ii) exposure time $>7$\\,ks, (iii) declinations $\\rm DEC(J2000)  > - 10 \\, deg$, (iv) galactic latitude $\\rm  |b_{II}|>20 \\, deg$ and (v) right ascension $\\rm RA(J2000)>4$\\,hours. A total of 51 fields fulfil the above criteria, covering a total area of $\\approx \\rm 6\\, deg^2$ to the 0.5-2\\,keV band limit  $f_X(\\rm 0.5 - 2 \\,  keV) \\approx 10^{-15} \\, erg \\, s^{-1}   \\, cm^{-2}$.  The USNO A2.0 catalogue provides a homogeneous platform that allows identification of X-ray selected normal galaxy candidates on the basis of the low X-ray--to--optical flux ratios of these sources, $\\log f_X /f_{opt} < -2$. Reliable star/galaxy separation is available from the APM, allowing us to exclude Galactic stars that also have $\\log f_X /f_{opt} < -2$. Optical spectroscopy is then used to identify systems that show evidence for AGN activity resulting in a sample of 23 normal galaxy candidates: 9 with narrow emission lines, 12 with absorption lines only and 2 without optical spectroscopic information. Future releases of the  {\\it XMM-Newton} Serendipitous Source Catalog will provide much wider areal coverage resulting in significantly larger low-$z$ normal galaxy samples.  At present we increase our sample size by combining it with X-ray selected normal galaxy candidates from the Needles in the Haystack Survey  (Georgantopoulos et al. 2005). This provides a total of 46 $z \\la 0.2$ X-ray detected normal galaxies, the largest low-$z$ sample yet available. Such a large number of sources provides a unique opportunity to constrain the normal galaxy $\\log N - \\log S$  at bright fluxes ($\\rm 10^{-15} -10^{-13} \\, erg \\, s^{-1} \\,   cm^{-2}$). We estimate a slope of $-1.46\\pm0.13$ consistent with the euclidean prediction and in agreement with previous determinations (Hornschemeier et al. 2003; Tajer et al. 2005).   Our enlarged sample is further combined with 23 local ($z \\la 0.2$) galaxies from the {\\it Chandra} Deep Field North and South surveys to construct  the local X-ray luminosity function of  normal galaxies. This is fit with a  Schechter function with a break at $\\rm \\log L_\\star = 41.02_{-0.12}^{+0.14} \\rm \\, erg \\, s^{-1}$ and a slope of $\\alpha=-1.76\\pm 0.10$. The large sample size allows us to also estimate the luminosity function of emission and absorption line systems separately and to provide the most accurate estimates of the 0.5-2\\,keV X-ray emissivity of these systems in the nearby Universe. We find $j_X= (0.24 \\pm 0.02) \\times 10^{38} $ and $(0.22 \\pm 0.03) \\times 10^{38} \\rm \\, erg \\, s^{-1}\\, Mpc^{-3}$ for emission and absorption line systems respectively. We note that these are lower than previous (less direct) estimates based on stacking analysis (e.g. Georgakakis et al. 2003b) or optically selected local galaxies (Georgantopoulos et al. 1999).  Finally, for the combined sample of 46 systems we explore the association between X-ray luminosity and host galaxy properties, such as SFR and stellar mass. We use the 1.4\\,GHz radio and $\\rm H\\alpha$ luminosities as star-formation indicators.  We find that the $L_X - L_{1.4}$ relation for our sources is consistent with that determined by Ranalli et al. (2003) for the local Universe. For the $\\log L_X - \\log L_{H\\alpha}$ correlation we estimate a slope of $\\approx 0.7$ in fair agreement with previous results (Zezas 2000). The above relations suggest that the X-ray luminosity in the emission-line galaxies in our sample is directly related to the current star-formation activity as measured by the 1.4\\,GHz radio and $\\rm H\\alpha$ luminosities. For the early type galaxies in the sample showing  absorption optical lines only we use the $K$-band as proxy to their stellar mass and find a linear correlation between $\\log L_X$ and $\\log L_K$ with a slope of $\\approx 1.5$. This is flatter than the slope of $\\approx 1.8$ for the $L_X - L_B$ relation for local ellipticals of Fabbiano et al. (1992) or the $L_X$/stellar-mass correlation for absorption line galaxies in the SDSS determined by Hornschemeier et al. (2005). This may be due to a possible bias in our sample against very luminous galaxies, $L_X > 10^{42} \\rm erg \\, s^{-1} \\, cm^{-2}$, likely introduced by the low X-ray--to--optical flux ratio selection. Alternatively, this may indicate that the $K$-band provides a better representation of the galaxy mass."
    },
    "0601/astro-ph0601494_arXiv.txt": {
        "abstract": "On January 12 of 2006, the Hubble Space Telescope (HST) observed the peculiar double extragalactic object CSL-1, suspected to be the result of gravitational lensing by a cosmic string. The high resolution image shows that the object is actually a pair of interacting giant elliptical galaxies. In spite of the weird similarities of the energy and light distributions and of the radial velocities of the two components, CSL-1 is not the lensing of an elliptical galaxy by a cosmic string. ",
        "introduction": "The Capodimonte-Sternberg-Lens candidate n. 1 (CSL-1; $\\alpha_{2000}=12^h$ $23^m$ $30\\fs5$, $\\delta_{2000}=-12\\degr$ $38\\arcmin$ $57\\farcs0$), a peculiar extragalactic object observed in the OACDF by \\citet{csl1}, hereafter Paper I, appears as a double source projected onto a low density field. The two components of the source are 1.9 arcsec apart, well resolved (i.e. extended) and, when observed from the ground, show roundish and identical shapes. Low resolution spectra look identical with a high confidential level (c.l.) and both give the same redshift of $0.46\\pm 0.008$. Photometry matches the properties of giant elliptical galaxies. Additional medium-high resolution observations carried on with FORS1 at VLT in March 2005 confirmed the similarity of the spectra of the two objects at a 98\\% c.l. \\citet{saz05} and showed that the differential radial velocity is compatible with zero $\\pm 20$ km s$^{-1}$.  In view of all that, already in Paper I we argued that CSL-1 could be either {\\it i)} the chance alignment of two very similar giant E galaxies at the same redshift, or {\\it ii)} a gravitational lensing phenomenon. In this second case detailed modelling shows that the properties of CSL-1 are compatible only with the lensing of an E galaxy by a cosmic string.  Cosmic strings were first introduced by \\citet{kib76} and extensively discussed in the literature (cf. \\citet{zel} and \\citet{vil}). Recent work has also shown the relevance of cosmic strings for both fundamental physics and cosmology (cf. \\citet{dav05}) and has shown that cosmic strings could be observable, in principle, via several effects, the most important being gravitational lensing, as we extensively discussed in the above quoted papers.  Is CSL-1 a cosmic string? In Paper I, we suggested an {\\it experimentum crucis} \\citep{csl1}: the detection of the sharp edges at faint light levels which are expected to be generated by a cosmic string (Fig.\\ref{string}). This test requires high angular resolution and deep observations to be performed with the Hubble Space Telescope (HST). In what follows we shortly summarise the most immediate outcomes of these observations.  \\begin{figure*} \\begin{center}\\label{HST1} \\includegraphics[height=5.5cm]{HST_vs_model.eps} \\end{center} \\caption{Left panel: numerical simulation of the image of an E galaxy with a de Vaucouleurs [$r^{1/4}$] light profile, as it would appear if lensed by a cosmic string (the situation roughly corresponds to that of CSL-1 and the noise level is that expected in our HST observations.). Right panel: the image region surrounding CSL-1 as observed by HST. A logarithmic colorscale has been adopted to enhance the  morphological details.} \\label{string} \\end{figure*} ",
        "conclusions": ""
    },
    "0601/astro-ph0601177_arXiv.txt": {
        "abstract": "{} {We report on a deep X-ray survey of the young ( $\\sim$ 140 Myr), rich open cluster NGC 2516 obtained with the EPIC camera on board the XMM-Newton satellite.} {By combining data from six observations, a high sensitivity, greater than a factor 5 with respect to recent Chandra observations, has been achieved. Kaplan-Meier estimators of the cumulative X-ray luminosity distribution are built, statistically corrected for non members contaminants and compared to those of the nearly coeval Pleiades. The EPIC spectra of the X-ray brightest stars are fitted using optically thin model plasma with one or two thermal components.} {We detected 431 X-ray sources and 234 of them have as optical counterparts cluster stars spanning the entire NGC 2516 Main Sequence. On the basis of X-ray emission and optical photometry, we indicate 20 new candidate members of the cluster; at the same time we find 49 X-ray sources without known optical or infrared counterpart.   The X-ray luminosities of cluster stars span the range $\\log \\mathrm L_\\mathrm X$ (erg s$^{-1}$) $=$ 28.4 -- 30.8. The representative temperatures span the 0.3 -- 0.6 keV (3.5 -- 8 MK) range for the cool component and 1.0 - 2.0 keV (12 -- 23 MK) for the hot one; similar values are found in other young open clusters like the Pleiades, IC 2391, and Blanco 1. While no significant differences are found in X-ray spectra, NGC 2516 solar type stars are definitely less luminous in X-rays than the nearly coeval Pleiades. The comparison with a previous ROSAT survey evidence the lack of variability amplitudes larger than a factor 2 in solar type stars in a $\\sim 11$ yr time scale of the cluster and thus activity cycles like in the Sun are probably absent or different by period and amplitude in young stars.} {} ",
        "introduction": "X-ray observations of stars have shown that X-ray emission is a common feature along the main sequence \\citep{Vaiana1981}. The emission mechanisms depends on masses: in massive O and early B type stars the X-rays are likely generated by plasma shocked and heated in stellar winds, although a role of magnetic field is not ruled out \\citep{Lucy1980,Feldmeier95,Babel97,Cassinelli2000,Waldron01}, whereas solar mass stars present coronae composed by hot,  magnetically confined emitting plasma, as directly observed in the Sun \\citep{Vaiana1981}. Only B-late -- A-early stars are not expected to emit significant X-rays because neither strong stellar winds nor convective regions, necessary to have a solar magnetic dynamo, are present in these stars. Focusing on solar mass stars, the X-ray activity is strictly linked to the internal structure and the rotation. The coupling of convective motions and rotation produces the solar-type dynamo. The losses of angular momentum is the main factor of decrease of X-ray activity during stellar evolution, especially in the first Gyr, although the role of other factors like, e.g., stellar chemical composition or environmental forming conditions cannot be ruled out.  In this context the studies based on stellar open clusters with different properties are of great importance to better define the relevance of age differences, chemical composition and environment influence. An open cluster like the Pleiades constitutes a prototype for the study of activity and evolution at an age around 100 Myr, given its richness and nearby distance. The evolution of X-ray emission with stellar age has been studied in the last five years with XMM-Newton and Chandra X-ray observatories, which provided high quality surveys of open clusters like Pleiades, Blanco 1, IC 2391, NGC 6383, NGC 188, M67 (cf. \\citealp{Briggs2003}, \\citealp{Pilli2004}, \\citealp{Marino04}, \\citealp{Rauw2003}, \\citealp{vanBerg04} and \\citealp{Gondoin05}). In this paper we study the X-ray emission of the young open cluster NGC 2516.  NGC 2516 is a rich open cluster with an age of 140 Myr \\citep{Meynet93}, slightly older than the Pleiades, located in the southern hemisphere at a height of $\\sim$100 pc on the Galactic Plane. The richness of star and its relatively nearby distance (about 387 pc) have motivated the growing interest for the study of this cluster both in the optical and in the X-ray band. In the optical band the  photometric study  of \\citet{Dachs70} and \\citet{DachsKab89} have shown a higher number of stars and a larger fraction of peculiar Ap stars with respect to the Pleiades. Episodes of subsequent star formation, as due to passages across the Galactic Plane or interaction with the Gum Nebula, have been also speculated by \\citet{DachsKab89} in order to explain younger blue stragglers in this cluster. More recently, the deep photometric survey of \\citet{Jeffries01} allowed to trace the cluster sequence down to V $\\simeq$ 20 and to build an optical catalog of 1254 photometric members. The {\\em Mass Function} of the cluster  is consistent with the empirical Salpeter's power law $dN/dM \\propto M^{-1.35}$ but with index 1.47--1.67, depending on the assumed metallicity. The chemical composition remains to date quite uncertain for this cluster: while photometry suggests a significant underabundance of metals ([Fe/H] = -0.3, \\citealp{Jeffries01}), more recently \\citet{Terndrup02}, on the basis of high resolution spectroscopy of two stars, strongly suggested a rather solar-like metallicity for NGC 2516. Furthermore, these authors conclude that an age of 150 Myr and a distance modulus of 7.93 mag (i.e. a distance of $\\sim$ 385.5 pc) are more appropriate.  In the X-ray band \\citet{Jef97} and \\citet{Micela00} analyzed ROSAT PSPC and HRI observations leading to 47 detections among cluster stars with spectral type from B2 down to K. The G- and K-type stars of the cluster appeared under-luminous when compared to the Pleiades thus inducing to consider the low metallicity of the cluster (estimated at that time by photometry) to have a role in determining the X-ray emission. With the advent of {\\em Chandra} and {\\em XMM-Newton} missions, the cluster has been studied by \\citet{Harnden01}, \\citet{Sciortino01} and \\citet{Damiani2003}, confirming the under-luminosity of solar type stars with respect to the Pleiades and reporting X-ray emission from several Ap stars. In M-type stars the X-ray emission was found similar to the Pleiades. The time X-ray variability of stars in NGC2516 has been studied by \\citet{Wolk04} with Chandra observations over a 2 year time span.  Unfortunately, cluster membership is determined only through photometry, since cluster proper motion is close to the solar value, making very difficult to obtain reliable astrometric membership determination. Therefore the low X-ray emission level in solar type stars could be attributed to both a significant contamination of the cluster sample or to a real effect, related to slightly older age of NGC~2516 with respect to the Pleiades.  In the present paper we report on a deep X-ray survey of NGC 2516 obtained through a series of observations with XMM-Newton satellite. Our aim is to determine with the highest sensitivity the X-ray emission of the cluster. In particular we will focus on the X-ray properties of solar type stars of this cluster and their comparison with those of the very similar and well studied Pleiades. The structure of the paper is as follows: in Sect. 2 we present the data and the analysis method. In Sect. 3 we discuss the spectral properties and the X-ray emission level of solar type stars. In Sect. 4 we compare the X-ray luminosities from XMM-Newton with those of Chandra and ROSAT. Sect. 5 reports a discussion on other stars suggested as possible new members of the cluster, based on X-ray emission and optical photometry. In Sect. 6 we give a brief summary of relevant results. ",
        "conclusions": "\\subsection{X-ray emission along the Main Sequence} Fig. \\ref{vbv} shows the Color-Magnitude Diagrams (CMD) in B, V, I bands for the photometric cluster members. Big dots are X-ray detections, small dots are the undetected cluster stars. Hotter stars are shown only in the left panel (CMD with B--V) while cooler stars with unreliable B--V are shown only in the CMD with V--I color index. The X-ray detections span the Main Sequence of the cluster from B type stars down to cool, low mass M-type stars.   \\begin{figure*} \\centering \\includegraphics[height=\\textwidth,angle=-90]{grafici.ps} \\caption{\\label{vbv} Color-Magnitude diagrams in B, V, I bands of NGC 2516 stars (small dots) from \\protect{\\citet{Jeffries01}} and \\protect{\\citet{DachsKab89}}. Big dots are X-ray detected stars.} \\end{figure*}  \\begin{figure} \\includegraphics[height=\\columnwidth,angle=-90]{bv_t.ps} \\caption{\\label{tem} Coronal temperatures as a function of B--V color. Filled dots refer to the cool component whereas empty diamonds refer to the hot component.} \\end{figure}  \\subsection{X-ray spectral properties} The sample of stars of which we have analyzed the spectra are constituted by solar G- and K- type stars and by hot, massive B- and early A- type stars. We were unable to analyze the spectra of stars with B--V$\\sim0.4$ (late-A spectral type stars) because of their lower X-ray fluxes  and hence poor statistic spectra. Although the models we used are too simple and unrealistic, a distribution of plasma temperatures being more realistic, we may determine the dominant temperatures of the coronae as usually done with low resolution X-ray spectra.  In Fig. \\ref{tem} we plot the temperatures derived from global fitting as a function of B--V color. We were able to fit most of the spectra of solar type stars with 2 temperature models. High energy tails of spectra are very scarce of photons, consequently hot components have been estimated with large errors. The temperatures are in 0.3--0.7 keV and 1.0--2.0 keV range, respectively, and are similar to those detected in the coronae of Pleiades and Blanco 1 open clusters, which are nearly coeval to NGC 2516 \\citep{Briggs2003,Pilli2004}. The ratio of hot to cool emission measure in 2-T models suggests that the two thermal components are comparable, the ratio being between 0.5 and 3.0 with a median of 1.3.  Spectra of all but 3 early type stars (B--V $\\la 0.5$) have low count statistics: we may  obtain a good fit of them with 1-T models. In this case the temperatures range around 0.6 keV, suggesting that the spectra are softer than those of solar mass stars.  The star DK 56 (a binary system with a B2V star and an unresolved late type companion) was bright enough in each observation to fit its spectrum with a 2-T model. We found a hot spectrum with temperatures of 0.8 and 3.0 keV and a spectrum shape similar to those of solar type stars. In the case of a wind driven X-ray emission, typical of massive stars, the spectrum is expected to be softer than in the case of coronae of solar mass stars. Therefore, based on the spectrum shape and temperatures, we suggest that the X-ray emission of DK 56 is likely due or strongly affected by the late type companion. This is a conclusion opposite to that of \\citet{Jef97} who, on the basis of its high X-ray luminosity, attributed all the X-ray emission of DK56 to the massive primary. This system is a blue straggler of the cluster, younger than $\\sim 25$Myr \\citep{DachsKab89}. Based on this age estimate, these authors discussed the hypothesis of subsequent episodes of star formation in NGC 2516 determined by the crossing of the cluster through the Galactic Plane or the Gum Nebula. Given this age estimate, the high X-ray luminosity of order of 10$^{30}$ erg s$^{-1}$ is not an atypical value for solar-type stars of age of 25 Myr or less. \\begin{figure*} \\includegraphics[height=\\textwidth]{xlf.ps} \\caption{\\label{xlf} Cumulative distribution functions of X-ray luminosities of NGC 2516 stars (solid black line) with different spectral types, evaluated with Kaplan Meyer estimators. Dotted lines are the distributions of Pleiades \\protect{\\citep{Giusi999}}. Solid gray line in G-, K- and M-type stars are the distribution of detections alone. Gray area in G-, K- and M-type panels marks the envelop of simulated Kaplan-Meier estimators after statistical correction for contaminants; the black square inside each area is the average of their medians. } \\end{figure*} \\subsection{\\label{xlfs} The X-ray luminosities of NGC 2516 late type stars} We have calculated the Kaplan-Meier estimators \\citep{Feigelson85} of the distribution of X-ray luminosities for F, G, K and M spectral types. The distributions take into account the presence of the upper limits in $\\log \\mathrm L_\\mathrm X$ that censorize the sample. When the lowest values are upper limits the distribution does not reach unity. Fig. \\ref{xlf} shows the X-ray distribution functions for F-, G-, K- and M-type stars of NGC 2516 (solid lines); for comparison we have added the ROSAT data of Pleiades (\\citealp{Giusi999}, dotted line) and the distribution considering only detections of NGC 2516 (solid gray line). This is an upper bound to the true distribution which must include undetected members. The Kaplan-Meier estimators of Pleiades and NGC 2516 are different \\footnote{By applying several two sample tests, the difference between the two curves is significant at a level $\\ga$ 99.9\\%.}. In particular, the luminosities of NGC 2516 of G-, K- and M- type stars are systematically lower than those of the Pleiades, the difference being larger for G- and K-type stars where the contamination of the optical catalog is higher than in the other spectral ranges, as stated by \\citet{Jeffries01}. Especially among K-type stars the Kaplan-Meier estimator and cumulative distribution of X-ray detections alone show the largest difference.  The contamination by less active field stars among G- and K-type stars in NGC 2516, hampers the evaluation of L$_\\mathrm X$ distributions of NGC 2516. However, also the X-ray luminosities of detections alone of G-, K- and M-type stars are systematically lower than the luminosities of the Pleiades at a level $\\geq 99\\%$ as it results by applying several two sample tests. The distribution of detections alone shown in Fig. \\ref{xlf} may be interpreted as an upper limit for the {\\em true} cumulative distribution of X-ray luminosity. If we were able to take into account the upper limits of only {\\em true} cluster members, the Kaplan-Meier estimators should be at lower luminosities with respect to the distribution of detections alone.  We have evaluated how many field stars are expected to fall among our X-ray detections. We focused on the main sequence field stars expected to be in a volume determined by the photometric selection of cluster Main Sequence and the sky area of our survey. By using the stellar density given in \\citet{Allen2000}, Table 19.14 (valid for stellar density on the galactic plane, NGC 2516 is $\\sim$100 pc above the Galactic Plane) we estimated 11 and 19 G and K-type main sequence field stars comprised in that volume and not related to the cluster.  From the X-ray luminosity distributions reported in \\citet{Schmitt95} and \\citet{Schmitt97} G-type field stars have $\\log \\mathrm L_\\mathrm X/(\\mathrm{erg}\\ \\mathrm{s}^{-1})$ in the range 26.5--29.5 with median 27.3; K-type stars occupy a narrower range (27.1--28.4) with a median of $\\log \\mathrm L_\\mathrm X/(\\mathrm{erg}\\ \\mathrm{s}^{-1})=$ 27.6. M-type stars in the solar neighborhood have luminosities $25.5 \\leq \\log \\mathrm L_\\mathrm X/(\\mathrm{erg}\\ \\mathrm{s}^{-1}) \\leq 29.1$ and a median of $\\log \\mathrm L_\\mathrm X/(\\mathrm{erg}\\ \\mathrm{s}^{-1}) \\sim 27$. Hence we estimated that less than 3 field stars in each spectral type should have been detected among the photometric sample (at the same time a sample of 10 to 18 stars should be undetected in the surveyed volume). More distant giants should not be detected and thus they should not affect the distribution of detections alone. The expected field star detections are thus very few compared to the number of detected cluster members in each sample (52 G-, 74 K- and 59 M-type detections), thus the distribution of detections alone should not be affected by field stars.  We have taken into account the expected fraction of contaminants given by \\citet{Jeffries01} in order to estimate {\\em unbiased} X-ray luminosity distribution functions of NGC 2516 G-, K- and M-type stars. By following the method used in \\citet{Damiani2003}, we excluded a fraction of upper limits from the X-ray luminosity sample in each spectral type according to the membership probability and assuming that the contaminants yield only upper limits; then we calculated the Kaplan-Meier estimator. This simulation procedure was iterated 500 times producing thus the bunch of curves plotted in light gray in Fig. \\ref{xlf} and which define the region where more likely the {\\em true} distribution function should lie. Only in the case of G-type stars we observe a larger difference between the curve with all upper limits and the family of contaminant corrected curves. In the K- and M-type stars the large number of upper limits is less affected by the statistical correction for contaminants. We point out that in all cases the gray region is at lower X-ray luminosities with respect to the Pleiades cumulative distributions.  The averages of the 0.5 quantile of each simulated curve are marked in each panel with a black square; the $\\log$ L$_ \\mathrm X /(\\mathrm{erg}\\ \\mathrm{s}^{-1})$ values are 29.04, 28.82 and 28.68 for G-, K- and M-type stars, respectively; the width of the gray region at the 0.5 quantile is $\\sim$0.05 dex which we assume as the uncertainty on the previous averages. \\citet{Guedel04} and \\citet{Favatarev03} discuss X-ray emission as a function of age for stars from young stellar objects to aged open clusters and the Sun. We have plotted in Fig. \\ref{lxage} the medians of $\\log \\mathrm L_\\mathrm X$ for F-, G-, K- and M-type stars reported by \\citet{Guedel04} and our $\\log \\mathrm L_\\mathrm X$ medians (after contaminant correction) for the same spectral types as a function of age. We have assumed here an age of 140 Myr for NGC 2516. It is evident the decrease of the mean level of X-ray luminosities with age, after a plateau at $\\sim10^{30}$ erg s$^{-1}$ which extends from Orion and Taurus Star forming Region ages (few Myr) to Alpha Persei cluster age (50 Myr). The X-ray under-luminosity of NGC 2516 with respect to the nearly coeval Pleiades remains evident, also taking into account the statistical correction for non member contamination and any systematic uncertainty discussed in Sect. 2.2.  \\begin{figure} \\includegraphics[height=\\columnwidth,angle=-90]{lx_age.ps} \\caption{\\label{lxage} $\\log \\mathrm{L}_\\mathrm X$ vs. age for different open clusters, star forming regions and the Sun. Open symbols: data of all but NGC 2516 clusters from \\protect{\\citet{Guedel04}}; solid symbols: medians of NGC 2516 F-, G-, K- and M-type stars (contaminants corrected, see. Sect. 3.3) from the present work.} \\end{figure} The causes of the different level of luminosity of solar type stars in NGC 2516 and the Pleiades could be attributed to several factors. The rotation is a key parameter for X-ray emission \\citep{Pallavicini81} up to rotation as fast as v sin$i\\sim 20$ km/s, when the L$_\\mathrm X$ -- rotation relation flattens with respect to rotation. Furthermore, angular momentum losses due to magnetic braking link stellar rotation to age and X-ray activity. A different distribution of rotation rates in the Pleiades and NGC 2516 could explain the low X-ray luminosity observed in the latter cluster. \\citet{Jef98} studied the stellar rotation of solar type stars of NGC 2516, finding a lower rotational rate in F- early and G- type stars with respect to the Pleiades whereas no clear difference between these two clusters was found among K-type stars. However \\citet{Terndrup02} concluded that no significant difference in rotation distribution is present between the Pleiades and NGC 2516, if this latter is assumed slight older than Pleiades. The under-luminosity of NGC 2516 in X-rays seems arising essentialy from the slight difference of age between the two clusters (80--100 Myr for Pleiades, 140 Myr for NGC 2516), although the past history of NGC 2516, its location and kinematics could have played a role in determining peculiar forming conditions, perhaps different from those of the Pleiades with an influence on the X-ray emission evolution.  \\subsection{L$_\\mathrm X$ to L$_\\mathrm{bol}$ ratio} An indicator of activity is the ratio of X-ray to stellar bolometric luminosity, \\lxlbol. Because of the dependence of X-ray luminosity on rotation, the \\lxlbol\\ ratio follows a power law relation with index -2 with rotational period and the Rossby number, i.e. the ratio between rotational period and convective turnover time in {\\em non saturated} stars (see \\citealp{Patten96}, \\citealp{Randich00}, \\citealp{Nicola2003}). ROSAT observations of young open clusters in the band 0.1--2.4 keV have also shown that low mass stars exhibit a saturation of this ratio at the level of $\\sim$ -3, i.e. the X-ray luminosity reaches at most 1/1000 of bolometric luminosity of the star. Furthermore, saturation seems to occur at earlier types in younger stars (B--V$ = 0.7$ or late G-type stars at the age of the Pleiades, B--V $ = 1$ or mid K-type stars at 220 Myr, \\citealp{Prosser95}).  We calculated the bolometric luminosities of NGC 2516 by interpolating L$_\\mathrm{bol}$ as a function of (B--V) based on a isochrone model with Z=0.02 and age of 140 Myr \\citep{Siess2000}; for stars with (B--V)$_0\\geq 0.3$ we used (V--I) color instead of (B--V). Fig. \\ref{lxlb} shows \\lxlbol\\ of NGC 2516 stars versus their B--V$_0$ and V--I$_0$ colors, with open and filled circles for single and binaries, respectively. We did not correct L$_\\mathrm{bol}$ for unresolved binary systems. These unresolved binaries trace an upper sequence in the plots. Arrows represent upper limits and the bottom curve traces approximately the detection limit at the cluster distance ($\\log$L$_\\mathrm{X}\\sim 28.4$). A spread of $\\sim$1.4 dex is observed among detections in $0\\leq B-V \\leq1.5$ where \\lxlbol\\ may reach values up to -2.4. The highest point at $\\sim$-1.4 is due to a very cool star (V--I= 3.06, id. 6160 in \\citealp{Jeffries01} catalog): however this star is very close to the edge of field of  view and was observed only during the first observation, hence its X-ray luminosity could be quite uncertain.  The \\lxlbol\\ ratio flattens at (V--I)$_0 \\sim 1.5$ or (B--V)$_0 \\sim 0.7$, which corresponds to the spectral type K5, according to the color--temperature calibration by \\citet{kenhart95} and we observe a saturation value of \\lxlbol\\ $\\sim -2.5$ instead of -3. A comparison with L$_\\mathrm{bol}$ used by \\citet{Jef97} shows that our L$_\\mathrm{bol}$ are on average lower by 0.1 dex for stars detected in both surveys. This difference would produce essentially a shift toward high \\lxlbol\\ ratio, while the trend and the spectral type at which the saturation of \\lxlbol\\ occurs should not be changed.  \\begin{figure*} \\includegraphics[height=\\textwidth,angle=-90]{lxlbol.ps} \\caption{\\label{lxlb} Logarithmic X-ray to bolometric luminosity ratio of NGC 2516 stars. Open circles: detected single stars, filled circles: detected binaries, vertical arrows: upper limits}. Lines below data points roughly trace the detection limit at the cluster distance ($\\log L_\\mathrm X \\sim 28.4$). \\end{figure*}"
    },
    "0601/astro-ph0601388_arXiv.txt": {
        "abstract": "We have modelled diffuse far-ultraviolet spectrum observed by {\\it FUSE} near M42 as scattering of starlight from the Trapezium stars by dust in front of the nebula. The dust grains are known to be anomalous in Orion with R$_{V} = 5.5$ and these are the first measurements of the FUV optical properties of the grains outside of ``normal'' Milky Way dust. We find an albedo varying from 0.3 $\\pm$ 0.1 at 912 \\AA\\ to 0.5 $\\pm$ 0.2 at 1020 \\AA\\, which is consistent with theoretical predictions. ",
        "introduction": "The Orion Nebula (M42) was first observed as one of the brightest diffuse sources in the ultraviolet (UV) sky by \\citet{Car77}. They identified this light to be due to the radiation from the bright Orion stars scattered by dust in the Orion Molecular Cloud. Further observations allowed \\citet{Wen} to suggest a model for the M42 region in which M42, itself, is a thin blister of ionized gas in front of the Orion Molecular Cloud. The scattering arises in a neutral sheet in front of the Nebula known as the Veil \\citep{Odell92} where there is a deficiency of small grains compared to the diffuse ISM \\citep{Baade,Costero,Cardelli}.  The first observations of far ultraviolet (FUV: 905 -- 1187 \\AA ) emission in this region were made by \\citet{Mu05} who used serendipitous pointings of the {\\it Far Ultraviolet Spectroscopic Explorer} ({\\it FUSE}) to find intensities as high as 3 $\\times$ 10$^{5}$ ph cm$^{-2}$ s$^{-1}$ sr$^{-1}$ \\AA$^{-1}$. In this work, we have modelled these observations in order to determine the optical properties of the dust grains in the FUV. We find that the albedo of these grains varies from 0.3 $\\pm$ 0.1 at 912~\\AA\\ to 0.5 $\\pm$ 0.2 at 1020 \\AA\\ which is consistent with the theoretical predictions of \\citet{Draine}. ",
        "conclusions": "We have modelled the intense diffuse light observed near M42 by \\citet{Mu05} as starlight from the Trapezium stars scattered by interstellar dust in Orion's Veil, a sheet of neutral hydrogen in front of the Orion Nebula. Most of the absorption seen in the spectrum is due to material between the Trapezium and the scattering location and is much less than the absorption along the direct line of sight to the Trapezium.  If we fix $g$ at 0.55, the albedo of the interstellar grains increases from a value of 0.3 $\\pm$ 0.02 at 912 \\AA\\ to 0.5 $\\pm$ 0.03 at 1020~\\AA\\, close to previously observed values but with a different R$_{V}$. On the other hand, if we assume a $g$ of 0.85, as observed by \\citet{Burgh02}, the albedo increases from 0.40 $\\pm$ 0.02 to 0.68 $\\pm$ 0.03. We have restricted our analysis up to 1020~\\AA\\ since molecular hydrogen fluorescence contaminates the spectrum longward of 1020~\\AA. These observations are the first of dust grains with an R$_{V}$ significantly different from the Galactic norm in the FUV; however we do not find a significant difference in the optical properties."
    },
    "0601/astro-ph0601341_arXiv.txt": {
        "abstract": "{We solve the integral equation describing the propagation of light in an isothermal plane-parallel atmosphere of optical thickness $\\tau^*$, adopting a uniform thermalization parameter $\\epsilon$. The solution given by the ALI method, widely used in the field of stellar atmospheres modelling, is compared to the exact solution. Graphs are given that illustrate the accuracy of the ALI solution as a function of the parameters $\\epsilon$, $\\tau^*$ and optical depth variable $\\tau$. ",
        "introduction": "The solution of the radiative transfer equation (RTE) is at the heart of the stellar atmospheres modelling, since this equation has to be solved typically thousands of times in order to construct a realistic model.  It is thus crucial to get a clear idea of the accuracy with which the RTE is solved, and the effect it has on the determination of the main physical quantities of the model: populations, electron density, temperature, etc.  In this article, we focus on the first point checking the ALI method for solving the integral form of the RTE, since this method is nowadays at the basis of most numerical schemes used to determine the radiation field in stellar atmospheres. We recall that ``ALI'' means Accelerated (or Approximate) Lambda Iteration, the Lambda operator being defined by Eqs.~(\\ref{eq_lambda})-(\\ref{eq_e1}) below for the scattering law we adopt here. The ALI code used in this paper is a combination of an accelerated iterative method (with a diagonal $\\Lambda$-operator) and a formal solver based on parabolic short characteristics. Recent reviews on this approach are Paletou (2001), Hubeny (2003) and section 3 of Trujillo Bueno (2003).  The accuracy of our ALI code is tested while applied to a well-known problem consisting of a homogeneous, isothermal slab with isotropic and monochromatic light scattering (Sec.~\\ref{sec_2}). Indeed, this idealized problem can be solved exactly, which allows for a direct comparison with the solution given by the ALI method. This problem is very simple on physical grounds but implies analytical and numerical calculations that are far from trivial. It contains the seeds of most of the difficulties met when solving the RTE in a thick, highly scattering medium.  It thus provides an excellent test for numerical codes since very accurate analytical solutions are available (Sec.~\\ref{sec_3}).  After a brief description of our ALI code, we move to the numerical tests in Sec.~\\ref{sec_4}, which is the main part of this paper.  The link with previous studies on the subject (Trujillo Bueno \\& Fabiani Bendicho 1995, Trujillo Bueno \\& Manso Sainz 1999) is finally commented in Sec.~\\ref{sec_5}. ",
        "conclusions": "Our ALI code has been subjected to a wide range of tests, revealing at the same time its capabilities and its limits.  Before developing these two points, we note that our conclusions are relative to the particular code we have used (based on Jacobi's method), specifically solving the standard problem (\\ref{eq_s1})-(\\ref{eq_e1}). The accuracy of the code is ultimately determined by the accuracy of the formal solver we have used (parabolic short characteristics).  We have checked the great robustness of our code, which is certainly its most remarkable feature.  It is able to solve the standard problem (\\ref{eq_s1})-(\\ref{eq_e1}) for a wide range of input parameters $\\epsilon$ and $\\tau^*$, with no important lack of performance when $\\epsilon \\to 0$ and/or $\\tau^* \\to +\\infty$.  However, the lowest accuracy of the ALI numerical solutions happens in the outermost layers of a star, corresponding to $\\tau$ lower than the thermalization depth $1/k(\\epsilon)$, these layers forming, by definition, the atmosphere of the star. The accuracy of our code is not better than, say $10^{-2}$, when we choose $n_\\tau = 9$, $n_\\mu = 5$ and limit the number of iterations to $N<100$, as it is currently done in stellar atmospheres modelling. To improve the accuracy of the calculations up to $\\sim 10^{-3}$, the parameters $n_\\tau$, $n_\\mu$ and $N$ should be set to larger (but today impractical) values when solving the radiative transfer equation on a large frequency spectrum, i.e. at thousands of frequencies. Indeed we pointed out a truly noticeable improvement of the accuracy when using finer grids in $\\tau$ or $\\mu$. Such an observation was made easier by the use of a very accurate reference solution. Of course, increasing the level of refinement of both spatial and angular quadratures has a strong impact upon the number of iterations needed for convergence. However, to overcome this difficulty while keeping the same accuracy on the numerical solutions, methods based on Gauss-Seidel and successive over-relaxation iterations were already proposed by Trujillo Bueno \\& Fabiani Bendicho (1995).  Another important question is relative to the propagation of errors in a stellar atmosphere model: to what extent are the main quantities provided by the model (populations of heavy particles, electron density, pressure, etc.) sensitive to the accuracy on the RTE solution? We intend to tackle this subject in a future work by constructing an accurate -- but still very idealized -- stellar atmosphere model, in which the main quantities are first derived from an exact solution to the RTE, and then from the solution given by a ALI-based numerical method."
    },
    "0601/astro-ph0601355_arXiv.txt": {
        "abstract": "{} {We study the evolution of the luminosity function (LF) of type-1 and type-2 Active Galactic Nuclei (AGN) in the mid-infrared, and then derive the contribution of the AGN to the Cosmic InfraRed Background (CIRB) and the expected source counts to be observed by $S\\!pitzer$ at 24\\,$\\mu$m.} {We used a sample of type-1 and type-2 AGN selected at 15$\\,\\mu$m ({\\it ISO}) and 12$\\,\\mu$m ({\\it IRAS}), and classified on the basis of their optical spectra. Local spectral templates of type-1 and type-2 AGN have been used to derive the intrinsic 15$\\,\\mu$m luminosities.  We adopted an evolving smooth two-power law shape of the LF, whose parameters have been derived using an un-binned maximum likelihood method.} {We find that the LF of type-1 AGN is compatible with a pure luminosity evolution ($L(z)=L(0)(1+z)^{k_L}$) model where $k_L\\sim$2.9.  A small flattening of the faint ($L_{15}<L^{*}_{15}$) slope of the LF with increasing redshift is favoured by the data.  A similar evolutionary scenario is found for the type-2 population with a rate $k_L$ ranging from $\\sim$1.8 to 2.6, depending significantly on the adopted mid-infrared spectral energy distribution.  Also for type-2 AGN a flattening of the LF with increasing redshift is suggested by the data, possibly caused by the loss of a fraction of type-2 AGN hidden within the optically classified starburst and normal galaxies. The type-1 AGN contribution to the CIRB at 15$\\,\\mu$m is (4.2--12.1)$\\times\\,10^{-11}\\;\\mathrm{W m^{-2} sr^{-1}}$, while the type-2 AGN contribution is (5.5--11.0)$\\times\\,10^{-11}\\;\\mathrm{W\\,m^{-2}\\,sr^{-1}}$.  We expect that $S\\!pitzer$ will observe, down to a flux limit of $S_{24\\,\\mu \\mathrm{m}}$=0.01\\,mJy,  a density of $\\sim$1200 deg$^{-2}$ type-1 and $\\sim$1000 deg$^{-2}$ type-2 optically classified AGN.} {AGN evolve in the mid-infrared  with a rate similar to the ones found in the optical and X--rays bands. The derived total contribution of the AGN to the CIRB (4-10\\%) and $S\\!pitzer$ counts should be considered as lower limits, because of a possible loss of type-2 sources caused by the optical classification.} ",
        "introduction": "Active Galactic Nucleus (AGN hereafter) are expected to have played an important role in the formation and evolution of the galaxies in the Universe. An example is the observation of the correlation between the mass of the central black holes ($M_{BH}$) and the mass of the bulges (Magorrian et al. 1998) or the velocity dispersion of gas and stars in the bulges of galaxies ($M_{BH}$--$\\sigma$ relation; Ferrarese \\& Merritt 2000; Tremaine et al. 2002). Thus, the measure of the history of the density of AGN as a function of the luminosity (the luminosity function, LF hereafter) can provide fundamental clues to explain the present day universe (e.g. Balland et al.  2003; Menci et al. 2004; Granato et al. 2004; Di Matteo et al. 2005).  The LF will not only provide informations on the demographics of AGN, but will also place constraints on the physical model of AGN, the origin and accretion history into supermassive black hole and the formation of structures in the early universe. Moreover, the measure of the AGN LF over the whole spectral range will provide information on the relative importance of accretion power in the overall energy budget of the universe.  AGN have been historically classified into two groups (type-1 and type-2), depending on the presence or absence in their optical spectra of broad emission lines (FWHM$>$1200 $km\\,s^{-1}$). The unified model for AGN assumes that the two types of sources are intrinsically the same (Antonucci et al. 1993), and that the observed differences in the optical spectra are explained by an orientation effect. The presence of obscuring material ({\\it torus} or {\\it warped disc}) around the central energy source and its orientation respect to the observer could prevent a direct view of the central region around the nuclei responsible for the broad line emission observed in the optical band. AGN are then classified in two groups: the un-obscured (type-1) and obscured (type-2) sources depending  whether the line-of-sight intersects or not the obscuring material. However, nowdays it seems clear that the orientation-dependent model is probably just a 0th-order approximation to the true nature of AGN (e.g. Barger et al. 2005; La Franca et al. 2005).  The major efforts to understand the evolution of AGN have been historically concentrated at optical wavelengths and in the last years have produced the largest compilations of AGN from the Two Degree Field (2dF; Boyle et al. 2000; Croom et al. 2004) and the Sloan Digital Sky Survey (SDSS; York et al. 2000; Hao et al. 2005). However the optical bands have proved to be very inefficient in the selection of obscured sources and the assembly of unbiased samples of type-2 objects is difficult.  In the last years many new windows have been opened at various wavelength bands for the observation of high redshift galaxies and AGN: sub-mm ({\\it SCUBA}), Infrared ({\\it IRAS}, {\\it ISO} and {\\it Spitzer}) and the X--ray ({\\it ROSAT}, {\\it XMM} and {\\it Chandra}). Thanks to their smaller dependence on dust obscuration if compared to the optical surveys, the Infrared and the hard X--ray wavelengths have proved to be much more efficient in the detection of type-2 sources.  X--ray observations have been so far consistent with population synthesis models based on the 0th-order unified AGN scheme (e.g. Comastri et al. 1995) or its modifications (Pompilio et al. 2000; Gilli et al. 2001; Ueda et al. 2003; La Franca et al. 2005).  The hard spectrum of the X--ray background is explained by a mixture of absorbed and unabsorbed AGN, evolving with cosmic time. According to these models, most AGN spectra should be heavily obscured as the light produced by accretion is absorbed by gas and dust.  For this reason an important mid-infrared (4-40\\,$\\mu$m) thermal emission is expected from reprocessed radiation by gas and dust grains directly heated by the central black hole (e.g. Granato et al. 1997; Oliva et al. 1999; Nenkova et al. 2002). AGN could be in this case important contributors to the Cosmic Infrared Background (CIRB). Indeed, {\\it IRAS} has observed strong mid-IR emission in all local ($z<0.1$) AGN (Miley et al. 1985; de Grijp et al. 1985; Neugebauer et al. 1986; Sanders et al. 1989). Complete and largely unbiased samples of AGN have been produced at 12$\\,\\mu$m (Rush, Malkan \\& Spinoglio 1993, RMS hereafter) and 25$\\,\\mu$m (Shupe et al. 1998) based on {\\it IRAS} observations.  These samples provide an important understanding of the infrared emission from local galaxies and a firm base to compare the properties of galaxies in the local universe with the high redshift populations uncovered by {\\it ISO} ($z\\leq 1$) and most recently by {\\it Spitzer} ($z\\leq 2$).  Statistical studies of AGN in the mid-IR have been based on large but local samples from the {\\it IRAS} observations (e.g. RMS, Shupe et al. 1998) and deep but small {\\it ISO} fields in the {\\it Hubble Deep Field North} (HDFN; Aussel et al. 1999a, 1999b), {\\it  South} (HDFS; Oliver et al. 2002) and the {\\it Canada-France Redshift Survey} (CFRS; Flores et al. 1999) that provided only a handful of objects.  Therefore, due to the small number of high redshift objects available, the shape and evolution of the mid-IR LF of these objects at high redshift was largely unknown. Only recently, large and highly complete 15$\\,\\mu$m spectroscopic samples, based on the {\\it European Large Area ISO Survey} (ELAIS hereafter; Oliver et al. 2000), have been released (Rowan-Robinson et al.  2004; La Franca et al. 2004; FLF04 hereafter). Based on the analysis of the {\\it ISO} observations in the ELAIS-S1 field, Matute et al. (2002, M02 hereafter) derived the first estimate of the type-1 LF in the mid-IR (15\\,$\\mu$m)  finding an evolution similar to the ones observed in the optical (Croom et al. 2004) and X--rays (e.g. Hasinger et al. 2005; La Franca et al. 2005). The evolution of the LF for normal and starburst galaxies detected at 15$\\mu$m in the southern ELAIS fields (S1 \\& S2) has been derived and discussed by Pozzi et al. (2004).  In this paper we investigate the separate evolution of type-1 and type-2 AGN in the mid-IR and their contribution to the CIRB using: $a$) the local IR population uncovered by {\\it IRAS}, $b$) deep observations from {\\it ISO} fields and, $c$) the most recent spectroscopic identifications of 15$\\mu$m sources in the southern fields of the ELAIS Survey (Pozzi et al. 2003; FLF04)\\footnote{Data and related papers about the ELAIS southern survey are available at: {\\sf http://www.fis.uniroma3.it/$\\sim$ELAIS\\_S}}.  The AGN sample used in our work is described in section 2. Section 3 introduces the general method used to derive the parameters of the luminosity function and presents the results. In section 4 we discuss our results and compare them with the LFs derived in the mid-IR by previous analysis of mid-IR observations. The contribution to the CIRB and the predicted counts in the mid-IR {\\it Spitzer} band at 24\\,$\\mu$m are computed and discussed in sections 5 and 6. Section 7 summarises the main results.  A cosmology with H$_\\mathrm{o}=75$ km s$^{-1}$ Mpc$^{-1}$ and $\\Omega_{\\mathrm \\lambda}=0.7$ , $\\Omega_{\\mathrm m}=0.3$ was adopted for the work presented in this paper. ",
        "conclusions": "Combining {\\it ISO}/15$\\,\\mu$m observations in ELAIS fields, HDF-N, HDF-S and the CFRS with the local IR population detected by {\\it IRAS} at 12$\\,\\mu$m, we have derived the evolution for the mid-IR selected  and optically classified AGN. While a similar study has already been done for QSOs+Seyfert-1 sources in the past (Matute et al. 2002), we have presented here for the first time the rate of evolution shown by the obscured, type-2 AGN. Our results are briefly summarised as: \\begin{itemize} \\item[-]{Type-1 AGN evolve following a double power law LF and a rate of evolution of the form (1+$z$)$^{k_L}$ with $k_L\\sim$2.9. The data are consistent, within the errors, with a PLE model. However,  a flattening at high redshift of the faint luminosity end of the luminosity function is marginally favoured by the data.} \\item[-]{Type-2 sources evolve with a slightly lower rate than type-1, with $k_L$ ranging from $\\sim$1.8 to $\\sim$2.6 depending strongly on the assumed SED used to compute the K-correction. The best fit solution favours a luminosity dependent luminosity evolution (LDLE) model, while the PLE model is statistically acceptable. However, integral counts reported for the deep 15$\\,\\mu$m observations in the HDF-N, -S and CFRS suggest an evolutionary scenario closer to the PLE model. We expect a significant number of type-2 AGN to be hidden within the optically classified normal and starburst galaxies.} \\item[-]{In the volume of the universe where both types of sources (type-1 \\& type-2) are observed ($z$=[0,0.7]) the ratio of type-2 to type-1 is $\\sim$2--6.  This value is in agreement, within the errors, with the one measured from the local optically selected samples, where a ratio close to 4 is found (Maiolino \\& Rieke 1995).} \\item[-]{The total contribution of AGN to the CIRB at 15$\\,\\mu$m is of the order of $\\sim$4--10$\\%$ of the total background light measured by Metcalfe et al. (2003), divided approximately equally between the two classes (type-1 \\& type-2). The contribution of type-1 sources is in good agreement with the results presented in previous studies (Matute et al. 2002). Comparison with results from X--ray selected samples (e.g. Fadda et al. 2002, Silva et al. 2004) shows that a significant fraction of obscured type-2 sources are missed in our sample. We argued that the optical spectral classification of the mid-IR sample, and not the mid-IR selection, is the principal responsible for the missed fraction.} \\item[-]{Estimates for the number of mid-IR selected and optically classified AGN, expected from {\\it Spitzer} at 24$\\,\\mu$m, are given according to the best model results. We expect $\\sim$1200 type-1 and $\\sim$1000 type-2 optically classified AGN per sq. degree at a flux limit of $S_{24\\mu m}=0.01$ mJy. Our 24\\,$\\mu$m counts derived for type-1 sources agree very well with previous works. At a flux limit of $S_{24\\mu m}=0.01$ mJy the expected counts for the obscured population from our best fit models are a factor $\\sim$5 lower than the expectations from models based on X--ray LFs (e.g Silva et al. 2004).} \\end{itemize}  The results presented here for type-1 sources are quite robust since a high completeness is expected for these sources. Unlike type-1, type-2 AGN show a larger spread in their mid-IR and optical  properties and a significant fraction can be misclassified as starburst or normal galaxies. As a consequence, the results for type-2 AGN derived here can only be considered as a lower limit to their true density and a first approximation to the evolution of these sources in the mid-IR. The true mid-IR space density of obscured sources has to be determined combining their mid-IR properties and the optical classification with the information available at other wavelengths, especially in the X--rays."
    },
    "0601/astro-ph0601449_arXiv.txt": {
        "abstract": "We present the interpretation of the muon and scintillation signals of ultra-high-energy air showers observed by AGASA and Yakutsk extensive air shower array experiments. We consider case-by-case ten highest energy events with known muon content and conclude that at the 95\\% confidence level (C.L.) none of them was induced by a primary photon. Taking into account statistical fluctuations and differences in the energy estimation of proton and photon primaries, we derive an upper limit of 36\\%  at 95\\% C.L.\\  on the fraction of primary photons in the cosmic-ray flux above $10^{20}$~eV. This result disfavors the $Z$-burst and superheavy dark-matter solutions to the GZK-cutoff problem. ",
        "introduction": "\\label{sec:intro} One of the most intriguing puzzles in astroparticle physics is the observation of air showers initiated by particles with energies beyond the cutoff predicted by Greisen and by Zatsepin and Kuzmin~\\cite{gzk}. Compared to lower energies, the energy losses of protons increase sharply at $\\approx 5\\times 10^{19}$~eV since pion production on cosmic microwave background photons reduces the proton mean free path by more than two orders of magnitude. This effect is even stronger for heavier nuclei, while photons are absorbed due to pair production on the radio background with the mean free path of a few Mpc. Thus, the cosmic-ray (CR) energy spectrum should dramatically steepen at $\\approx 7\\times 10^{19}$~eV for any homogeneous distribution of CR sources. Despite the contradictions in the shape of the spectrum, the existence of air showers with energies in excess of $10^{20}$~eV is firmly established by several independent experiments using different techniques (Volcano Ranch~\\cite{exp}, Fly's Eye~\\cite{Bird}, Yakutsk~\\cite{YakutskExperiment}, AGASA~\\cite{agasares}, HiRes~\\cite{HiRes} and Pierre Auger~\\cite{Auger} experiments). Some explanations for these showers, like the $Z$-burst or top-down models, predict a significant fraction of photons above typically $8\\times 10^{19}$~eV (for reviews see, e.g., Refs.~\\cite{reviews}). Indications for the presence of neutral particles at lower energies were found in Refs.~\\cite{neutral}. Thus, the determination of the photon fraction in the CR flux is of crucial importance, and the aim of this work is to derive a stringent limit on this fraction in the integral CR flux above $10^{20}$~eV. To this end, we compare the reported information on signals measured by scintillation and by muon detectors for observed showers with those expected by air shower simulations. We focus on the surface detector signal density at 600 meters $S(600)$ (known as charged particle density) and the muon density at 1000 m, $\\rho_{\\mu}(1000)$, which are used in experiments as primary energy and primary mass estimators, respectively.  We study individual events of AGASA~\\cite{AGASA_Eest} and of the Yakutsk extensive air shower array (Yakutsk in what follows)~\\cite{YakutskExperiment} with \\textit{reconstructed} energies above $8\\times 10^{19}$~eV and measured muon content. We reject the hypothesis that any of showers considered was initiated by a photon primary at the 95\\% confidence level (C.L.). We then derive as our main result an upper limit of 36\\% (at 95\\%~C.L.)  on the fraction $\\epsilon_\\gamma$ of primary photons with \\textit{original} energies above $10^{20}$~eV (the difference between original and reconstructed energies is discussed in Sec.~\\ref{sec:data}).  The rest of the paper is organized as follows. In Sec.~\\ref{sec:data} we discuss the experimental data set which we use for our study. In Sec.~\\ref{sec:simulations}, the details of the simulation of the artificial shower libraries and comparison of the simulated and real data are given. This section contains the description of our method and the main results. We discuss how robust these results are with respect to changes in assumptions, to analysis procedure, and to variations in the experimental data, in Sec.~\\ref{sec:robustness}. In Sec.~\\ref{sec:comparison}, we discuss the differences between our approach and previous studies, which allowed us to put a significantly more stringent limit on the gamma-ray fraction. Our conclusions are briefly summarized in Sec.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions} To summarize, we have studied the possibility that the highest-energy events observed by the AGASA and Yakutsk experiments were initiated by primary photons. Comparing the observed and simulated muon content of these showers, we reject this possibility for each of the ten events at $E>8\\cdot 10^{19}$~eV at least at the 95\\% C.L.  An important ingredient in our study is the careful tracking of differences between the original and reconstructed energies.  This allows us to put an upper bound of 36\\% at 95\\% C.L.\\ on the fraction $\\epsilon_\\gamma $ of primary photons with original energies $E_0>10^{20}$~eV, assuming an isotropic photon flux and $E_0^{-2}$ spectrum. This limit is the strongest one up to date. It strongly disfavors the $Z$-burst and constrains severely superheavy dark-matter models. The method that we have used is quite general and may be applied at other energies and to other observables.  We are indebted to M.~Kachelrie\\ss, K.~Shinozaki and M.~Teshima for numerous helpful discussions and collaboration at initial stages of this work. We thank L.~Bezroukov, V.~Bugaev, R.~Engel, D.~Heck, A.~Ringwald, M.~Risse, V.~Rubakov, D.~Semikoz and P.~Tinyakov for helpful discussions and comments on the manuscript. This study was performed within the INTAS project 03-51-5112. We acknowledge also support by fellowships of the Russian Science Support Foundation and of the Dynasty foundation (D.G.\\ and S.T.), by the grants NS-2184.2003.2 (D.G., G.R.\\ and S.T.), NS-1782.2003.2, RFFI 03-02-16290 (L.D., G.F., E.F.\\ and T.R.), NS 748.2003.2, RFFI 03-02-17160, RFFI 05-02-17857, FASI 02.452.12.7045 (A.G., I.M., M.P.\\ and I.S.). G.R.\\ and S.T.\\ thank the Max-Plank-Institut f\\\"ur Physik (Munchen), where a significant part of this work was done, for warm hospitality. Computing facilities of the Department of Theoretical Physics, Institute for Nuclear Research (Moscow), were used to perform the simulations of air showers."
    },
    "0601/astro-ph0601163_arXiv.txt": {
        "abstract": "Baryon acoustic oscillations (BAO) in the galaxy power spectrum allows us to extract the scale of the comoving sound horizon at recombination, a cosmological standard ruler accurately determined by the cosmic microwave background anisotropy data. We examine various issues important in the use of BAO to probe dark energy. We find that assuming a flat universe, and priors on $\\Omega_m$, $\\Omega_m h^2$, and $\\Omega_b h^2$ as expected from the Planck mission, the constraints on dark energy parameters $(w_0,w')$ scale much less steeply with survey area than (area)$^{-1/2}$ for a given redshift range. The constraints on the dark energy density $\\rho_X(z)$, however, do scale roughly with (area)$^{-1/2}$ due to the strong correlation between $H(z)$ and $\\Omega_m$ (which reduces the effect of priors on $\\Omega_m$). Dark energy constraints from BAO are very sensitive to the assumed linear scale of matter clustering and the redshift accuracy of the survey. For a BAO survey with $0.5\\leq z \\leq 2$, $\\sigma(R)=0.4$ (corresponding to $k_{max}(z=0)=0.086\\,h\\,$Mpc$^{-1}$), and $\\sigma_z/(1+z)=0.001$, $({\\sigma}_{w_0},{\\sigma}_{w'})=$(0.115, 0.183) and (0.069, 0.104) for survey areas of 1000 (deg)$^2$ and 10000 (deg)$^2$ respectively. We find that it is critical to minimize the bias in the scale estimates in order to derive reliable dark energy constraints. For a 1000 (10000) square degree BAO survey, a 1$\\sigma$ bias in $\\ln H(z)$ leads to a 2$\\sigma$ (3$\\sigma$) bias in $w'$. The bias in $w'$ due to the same scale bias from $\\ln D_A(z)$ is slightly smaller and opposite in sign. The results from this paper will be useful in assessing different proposed BAO surveys and guiding the design of optimal dark energy detection strategies. ",
        "introduction": "The most intriguing mystery in cosmology today is the nature of the dark energy that is driving the accelerated expansion of the universe as evidenced by supernova \\citep{Riess98,Perl99} and other data. Although current observational data are consistent with a cosmological constant (see for example, \\cite{WangTegmark04,WangTegmark05,Daly05,Jassal05}), viable dark energy models abound. See for example, \\cite{Freese87,Linde87,Peebles88,Wett88,Frieman95,Caldwell98,SH98,Parker99,DGP00, Mersini01,Freese02,Carroll04,OW04,Alam05,Cai05,Cardone05,Kolb05,MB05}. \\cite{Pad} and \\cite{Peebles03} contain reviews of many models.  The powerful complementarity of the cosmic microwave anisotropy (CMB) and galaxy survey data in precision cosmology has long been noted (see for example, \\cite{Bahcall99}, \\cite{Eisen99}, and \\cite{Wang99}). An important development in this complementarity is to use the baryonic acoustic oscillations (BAO) in the galaxy power spectrum as a cosmological standard ruler to probe dark energy \\citep{BG03,SE03}, made possible by accurate determination of the standard ruler scale by CMB data from WMAP \\citep{Bennett03,Spergel03} and Planck \\footnote{See Planck mission blue book at http://www.rssd.esa.int/SA/PLANCK/docs/Bluebook-ESA-SCI(2005)1.pdf}.  The ``wavelength'' of the baryonic acoustic oscillations in $k$-space is determined by $k_A=2\\pi/s$, where $s$ is the comoving sound horizon at the drag epoch (when baryons are released from the Compton drag of the photons\\footnote{For $\\Omega_bh^2 \\la 0.03$, the drag epoch follows the last scattering of the photons \\citep{EisenHu98}.}), \\be s=\\int_0^{t_{d}} c_s dt/a =H_0^{-1} \\int_{z_{d}}^\\infty dz\\, \\frac{c_s}{ H(z)} \\ee The sound speed is \\ba c_s &= &\\frac{1}{\\sqrt{3(1+R^b)}}, \\nonumber\\\\ R^b &=& \\frac{3\\rho_b}{4\\rho_{\\gamma}}=31.5\\,\\Omega_bh^2 (T_{CMB}/2.7\\mbox{K})^{-4} (z/1000)^{-1}, \\ea and the Hubble parameter is \\be \\label{eq:H(z)} H(z)=\\sqrt{\\Omega_m(1+z)^3+ \\Omega_{rad}(1+z)^4+ \\Omega_k(1+z)^2+ \\Omega_X X(z)}, \\ee with $X(z)\\equiv \\rho_X(z)/\\rho_X(0)$ denoting the dark energy density function. For a cosmological constant, $X(z)=1$ and $\\Omega_X =\\Omega_\\Lambda$. Since $z_{d}\\gg 1$, we have \\ba s &\\simeq &\\frac{H_0^{-1}}{\\sqrt{\\Omega_m}} \\int_0^{a_{d}} da \\,\\frac{c_s }{\\sqrt{a+a_{eq}}}\\nonumber\\\\ &=&\\frac{1}{\\sqrt{\\Omega_m H_0^2}} \\frac{2c}{\\sqrt{3z_{eq} R^b_{eq}}} \\, \\ln \\frac{ \\sqrt{1+R^b_{d}}+ \\sqrt{R^b_{d}+ R^b_{eq}} }{1+ \\sqrt{R^b_{eq}}}, \\ea where ``d'' and ``eq'' refer to the drag epoch and matter-radiation equality respectively, with $z_{d}\\sim 1089$, and $z_{eq}= 2.5\\times 10^4 \\Omega_m h^2 (T_{CMB}/2.7\\mbox{K})^{-4}$. Hence the ``wavelength'' of the baryonic oscillations depends very strongly on $\\Omega_m$, but has negligible dependence on dark energy (which only became important recently, at redshifts a few and less). It can therefore be used as a standard ruler to probe dark energy.  The systematic effects of BAO as a standard ruler are: bias between luminous matter and matter distributions, nonlinear effects, and redshift distortions \\citep{BG03,SE03}. Cosmological N-body simulations are required to quantify these effects \\citep{Angulo05,SE05,Springel05,White05}. Fisher matrix formalism is useful in investigating the impact of various effects and assumptions that can be parametrized in the observed galaxy power spectrum \\citep{SE03}.  In this paper, we examine various issues important in the use of baryon acoustic oscillations (BAO) to probe dark energy using the Fisher matrix formalism. In particular, we extend the work of Seo \\& Eisenstein (2003) by deriving the biases in dark energy parameters ($w_0$, $w'$) due to systematic errors in the estimated scales. Throughout this paper, we assume spatial flatness as motivated by inflation and consistent with current CMB data.   Sec.2 describes the method used in our calculations. We present results in Sec.3 and summarize in Sec.4. ",
        "conclusions": "Baryon acoustic oscillations (BAO) provide an important new method for probing dark energy. To extract reliable dark energy constraints from BAO surveys, we need to understand how the expected errors depend quantitatively on the assumptions made.  For simplicity and transparency, we use the Fisher information matrix formalism in this paper. Our results are consistent with those of \\cite{SE03}, and agree to about 10\\% with those found using Monte Carlo methods by \\cite{BG03} and \\cite{GB05}, for the same $k_{max}(z=0)$ and priors on cosmological parameters. Although Fisher matrix yields the smallest possible error bars, we have made the most conservative assumptions. In deriving the dark energy constraints, we marginalize over many additional physical parameters \\citep{SE03}: the linear redshift distortion $\\beta$, linear growth function $G(z)$, and an unknown shot noise $P_{shot}$ in each redshift slice; the matter density ratio $\\Omega_m$, the matter density $\\Omega_m h^2$, and the baryon density $\\Omega_b h^2$. This approach seems to work rather well, as \\cite{SE05} found that their results from cosmological N-body simulations are close to what they found using the Fisher matrix formalism \\citep{SE03}.  We have examined various issues important in the use of BAO to probe dark energy. In particular, we extend the work of Seo \\& Eisenstein (2003) by deriving the biases in dark energy parameters ($w_0$, $w'$) due to systematic errors in the estimated scales.  We find that assuming priors on $\\Omega_m$, $\\Omega_m h^2$, and $\\Omega_b h^2$ as expected from the Planck mission, the constraints on dark energy parameters $(w_0,w')$ scale much less steeply with survey area (for a given redshift range) than (area)$^{-1/2}$ (which holds in the absence of the priors). For a BAO survey with $0.5\\leq z \\leq 2$, $\\sigma(R)=0.4$ (corresponding to $k_{max}(z=0)=0.086\\,h\\,$Mpc$^{-1}$), and $\\sigma_z/(1+z)=0.001$, $({\\sigma}_{w_0},{\\sigma}_{w'})=(0.115,0.183)$ and $(0.069,0.104)$ for survey areas of 1000 (deg)$^2$ and 10000 (deg)$^2$ respectively (see Fig.1). This provides a useful guide in designing optimal survey strategies that provide robust and accurate dark energy constraints.  The constraints on the dark energy density $\\rho_X(z)$, however, do scale roughly with (area)$^{-1/2}$ due to the strong correlation between $H(z)$ and $\\Omega_m$ (which reduces the effect of priors on $\\Omega_m$). This is interesting since $\\rho_X(z)$ provides direct and model-independent constraints on dark energy \\citep{WangGarnavich01,Tegmark02,WangFreese04}.  The BAO constraints on the dark energy equation of state parameters ($w_0$,$w'$) are very sensitive to the assumed linear scale of matter clustering $k_{max}(z)$\\footnote{Once realistic N-body simulations are available for calibrating the analysis of real data, we will be able to extract additional information on dark energy parameters by using the data from the quasi-linear regime.} and the redshift accuracy of the survey $\\sigma_z/(1+z)$ (see Figs.2-3). This is important to note since different proposed BAO surveys are often compared by their expected errors in ($w_0$,$w'$). Such comparisons are not appropriate unless the same assumptions are made about $k_{max}(z)$, and the claimed $\\sigma_z/(1+z)$ can be demonstrated to be feasible. The latter is a key issue for surveys using photometric redshifts \\citep{BG03,SE03,Zhan05}.   It should be noted that the radial BAO are {\\it not} observable for photometric redshift accuracies. The $H(z)$ measurement from using photometric redshifts is derived from the degree of damping of broad-band power in the radial direction on very large scales ($k < 0.05 $h$\\,$Mpc$^{-1}$). Hence this measurement is much less robust (more susceptible to systematic error) than the spectroscopic case, where the radial BAO are detectable and give $H(z)$ measurements directly.  We find that it is critical to minimize the bias in the scale estimates ($\\ln H(z)$ and $\\ln D_A(z)$) in order to derive reliable dark energy constraints. For a 1000 (10000) sq deg BAO survey, a 1$\\sigma$ bias in $\\ln H(z)$ leads to a 2$\\sigma$ (3$\\sigma$) bias in $w'$. The bias in $w'$ due to the same scale bias from $\\ln D_A(z)$ is slightly smaller and opposite in sign. In addition to the scale bias due to the inaccurate modeling of nonlinear effects (which impacts the results in the linear regime, see \\cite{SE05}), systematic biases in the $H(z)$ measurement could arise when photometric redshifts are used. It is important to quantify the expected bias in $\\ln H(z)$ and $\\ln D_A(z)$ using cosmological N-body simulations. Note that $H(z)$ and $D_A(z)$ are measured as independent parameters. Since $D_A(z)$ is related to $H(z)$ through an integral for a given cosmological model in the absence of systematic biases, comparing the directly measured $D_A(z)$ to the $D_A(z)$ derived from the $H(z)$ measurement provides us with a cross-check to help model unknown systematic effects.    The planned BAO surveys from both ground and space telescopes will play an important role in unraveling the nature of dark energy. The results from in this paper will be useful in assessing different proposed BAO surveys and guiding the design of optimal dark energy detection strategies.  \\bigskip  {\\bf Acknowledgements} I am grateful to Chris Blake for making his Monte Carlo results available to me for comparisons, and for useful discussions. I thank Max Tegmark, Dan Eisenstein, and Hee-Jong Seo for helpful discussions. This work was supported in part by NSF CAREER grants AST-0094335."
    },
    "0601/astro-ph0601480_arXiv.txt": {
        "abstract": "Isothermal and adiabatic shocks, which are produced from fast expansion of the gas, is simulated with smoothed particle hydrodynamics~(SPH). The results are compared with the analytic solutions. The algorithm of the program is explained and the package, which is written in Fortran, is presented in the appendix of this paper. It is possible to change (to complete) the program for a wide variety of applications ranging from astrophysics to fluid mechanics. ",
        "introduction": "Gas dynamical processes are believed to play an important role in the evolution of astrophysical systems on all length scales. Smoothed particle hydrodynamics~(SPH) is a powerful gridless particle method to solve complex fluid-dynamical problems. SPH has a number of attractive features such as its low numerical diffusion in comparison to grid based methods. An adequate scenario for SPH application is the unbound astrophysical problems, especially on the shock propagation~(see, e.g., Liu \\& Liu 2003). In this way, the basic principles of the SPH is written in this paper and the simulation of isothermal and adiabatic shocks are applied to test the ability of this numerical simulation to produce known analytic solutions.  The program is written in Fortran and is highly portable. This package supports only calculations for isothermal and adiabatic shock waves. It is possible to change (to complete) the program for a wide variety of applications ranging from astrophysics to fluid mechanics. The program is written in modular form, in the hope that it will provide a useful tool. I ask only that: \\begin{itemize} \\item If you publish results obtained using some parts of this program, please consider acknowledging the source of the package. \\item If you discover any errors in the program or documentation, please promptly communicate the to the author. \\end{itemize} ",
        "conclusions": ""
    },
    "0601/astro-ph0601025_arXiv.txt": {
        "abstract": "We present a spectroscopic survey of 2046 red giant stars, distributed over the central 4~kpc$\\times$2~kpc of the Small Magellanic Cloud (SMC). After fitting and removing a small velocity gradient across the SMC ($7.9$~km~s$^{-1}$~deg$^{-1}$ oriented at 10$^\\circ$ E of N), we measure an rms velocity scatter of $27.5\\pm0.5$~km~s$^{-1}$. The line of sight velocity distribution is well-characterized by a Gaussian and the velocity dispersion profile is nearly constant as a function of radius. We find no kinematic evidence of tidal disturbances.  Without a high-precision measurement of the SMC's proper motion, it is not possible to constrain the SMC's true rotation speed from our measured radial-velocity gradient.  However, even with conservative assumptions, we find that $v < \\sigma$ and hence that the SMC is primarily supported by its velocity dispersion. We find that the shape of the SMC, as measured from the analysis of the spatial distribution of its red giant stars, is consistent with the degree of rotational flattening expected for the range of allowed $v/\\sigma$ values. As such, the properties of the SMC are consistent with similar low luminosity spheroidal systems.  We conclude that the SMC is primarily a low luminosity spheroid whose irregular visual appearance is dominated by recent star formation. A simple virial analysis using the measured kinematics implies an enclosed mass within 1.6~kpc of between 1.4 and $1.9\\times10^9 M_\\odot$, and a less well constrained mass within 3~kpc of between 2.7 and $5.1\\times10^9 M_\\odot$. ",
        "introduction": "\\label{sec:intro}  There is an abundant and growing body of evidence that gravitational interactions have been critical in the evolution of the Magellanic Clouds. One prominent piece of evidence is the existence of the Magellanic Stream \\citep{mat74}, a roughly contiguous filament of neutral hydrogen that stretches more than 150~kpc behind the Clouds and which numerical simulations suggest was stripped by a recent interaction among the Clouds and the Milky Way \\citep{gn96}.  The Clouds are the nearest  galaxies that are globally affected by interactions yet are far from being disrupted. As such, they are detailed examples of galaxies expected to be common in a hierarchical model of structure formation and may provide general insights on how such weak interactions affect both the observed characteristics and evolution of galaxies.  For the Small Magellanic Cloud (SMC), we have presented results on the effects of the interactions on both its global properties and its evolution. First, we demonstrated that the irregular appearance of the SMC is due solely to recent star formation, because its older stellar populations are distributed in a regular, smooth ellipsoid \\citep{zar00}. This result cautions us that galaxies that appear to be irregular, particularly low mass galaxies like the SMC, may indeed have quite a different underlying structure --- a conclusion that is obviously critical to such considerations as whether dIrr's can evolve into dE's.  Second, we found that the SMC has experienced periodic episodes of global enhanced star formation, suggesting triggering by perigalactic passes with the Large Magellanic Cloud and with the Milky Way \\citep{hz04,zh04}.  These results demonstrate for one galaxy, and suggest for others, that basic observational parameters such as morphology, luminosity, and color are significantly affected by star formation induced by weak interactions.  Here, we investigate (1) whether such interactions also induce detectable kinematic disturbances, and so whether common kinematical analyses, such as the classification of galaxies by the ratio of rotational velocity to velocity dispersion and the application of the virial theorem to estimate the mass, are valid for galaxies affected by weak interactions and (2) whether the classification of the SMC as a dE based on the distribution of the older stars \\citep{zar00} is substantiated by the kinematics.  The conventional picture of the SMC as a disturbed galaxy is derived from the kinematics of stars in its outer regions \\citep{hat93,hat97}, and from a large observed line-of-sight depth \\citep{hh89, gh91, cro01}.  It is important to note that these line-of-sight depth studies also generally sample the outer regions of the SMC.  The depth of the central regions of the SMC has been probed using Cepheid variables, but there seems to be no consensus on the resulting depth.  While some authors find a large line-of-sight depth of around 20~kpc \\citep{mat86, mat88, gro00}, others conclude the depth is only 6--8~kpc \\citep{wel87, lah05}. Much of the controversy can be resolved by simply recognizing that the term ``depth'' is used differently by the two groups of authors. Those who define the depth as the full front-to-back extent of the tracer population obtain a depth around 20~kpc, while those who define the depth as the full width at half-maximum of the distribution of distances obtain a much smaller depth around 6--8~kpc.  The conclusion to be drawn from this body of work is that while the outer regions of the SMC may be extended along the line of sight (up to 20~kpc), especially toward the Northeast, its central region seems to be only modestly extended, with a depth of 6--8~kpc.  Our knowledge of the internal kinematics of the main body of the SMC, the region that would typically be observed in more distant galaxies, is surprisingly limited, considering the SMC's proximity and suspected complex internal structure.  Previous studies have determined mean radial velocities and velocity dispersions, based on modest numbers of tracer populations:  44 central planetary nebulae \\citep{dop85}, 12 outer RGB stars \\citep{sun86}, 131 central Carbon stars \\citep{har89}, 30 outer red clump stars \\citep{hat93}, and 70 outer Carbon stars \\citep{hat97}. In each case, the observed sample was too small to support analysis beyond the mean velocity and velocity dispersion of the sample.  In this paper, we present a spectroscopic survey of over 2000 red giant branch stars in the main body of the SMC, using the Inamori Magellan Areal Camera and Spectrograph \\citep[IMACS;\\ ][]{big98} at the Magellan Baade Telescope at Las Campanas Observatory in Chile. In Section~\\ref{sec:obsdata}, we describe the observations and data reduction.  In Section~\\ref{sec:kinematics}, we discuss whether the SMC kinematics follow the expectations for a dE and whether we find any kinematic signature of the weak interactions the SMC has recently experienced.  We summarize the results in Section~\\ref{sec:summary}.  In addition to obtaining radial-velocity kinematics of a large sample of SMC field stars, the observations were also designed to yield metallicities based on the equivalent width of the Ca~{\\small II} triplet feature.  The metallicities of our sample and their constraint on the chemical enrichment history of the SMC are presented in a subsequent paper \\citep{har05}. ",
        "conclusions": "\\label{sec:summary}  We present the first spectroscopic survey of field red giant stars over the main body of the Small Magellanic Cloud.  Our spectroscopic sample includes over 3000 red giants in 14 28$^\\prime$ fields spanning the central $2^\\circ\\times1^\\circ$ of the SMC.  After rejecting two fields due to unresolved pointing errors and losing additional spectra to overlap issues, our final kinematic sample contains 2079 stars throughout the central 3.9~kpc$\\times$1.8~kpc, of which 2046 are determined to be SMC members.  The global velocity distribution is well-fit by a Gaussian centered at 146~km~s$^{-1}$ with a width or $\\sigma$ of 28~km~s$^{-1}$, but we also find evidence for a velocity gradient across the SMC.  Removing the gradient using a planar fit to the velocities results in an rms scatter ($27.5\\pm0.5$~km~s$^{-1}$) that is effectively unchanged from the global velocity dispersion measured prior to any correction because the amplitude of the fitted rotation gradient (7.9~km~s$^{-1}$~deg$^{-1}$) is small. The origin of the velocity gradient could be either internal rotation or an illusion caused by the differential projection of the SMC's tangential velocity along non-parallel lines of sight. Without a high precision measurement of the SMC's proper motion, a sufficiently precise correction required to recover the SMC's intrinsic rotation speed is impossible.  However, even adopting the maximum correction we find that $v < \\sigma$ and hence that the SMC is primarily supported by its velocity dispersion.  We find that the shape of the SMC, as measured from the analysis of the spatial distribution of red giant stars, is consistent with the degree of rotational flattening expected for the observed $v/\\sigma$ (Figure \\ref{fig:vsigma}). As such, the properties of the SMC are consistent with similar low luminosity spheroidal systems. We conclude that the SMC is a low luminosity spheroid whose visual appearance is dominated by star formation resulting from a recent accretion of gas \\citep[see also][]{zar00}. The underlying morphology and dynamics are consistent with those of low luminosity ellipticals.  A simple virial analysis using the measured kinematics implies an enclosed mass within 1.6 kpc of between 1.4 and 1.9$\\times 10^9 M_\\odot$, and a less well constrained mass of between 2.7 and $5.1\\times 10^9 M_\\odot$ within 3 kpc. These values are larger than those inferred from the H{\\small I} kinematics \\citep[$2.4\\times10^9 M_\\odot$,][] {ssj04}, but the simplifying assumptions used could easily account for differences at the factor of two level.   \\vskip 1in \\noindent Acknowledgments: The authors thank Gus Oemler for his generous assistance in getting the COSMOS software working with the misaligned 600~l~mm$^{-1}$ grating; John Moustakas for providing an IDL program for robust spectral line fitting; and J.~D. Smith for general IDL assistance. JH is supported by NASA through Hubble Fellowship grant HF-01160.01-A awarded by the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Inc., under NASA contract NAS 5-26555.  DZ acknowledges financial support from National Science Foundation CAREER grant AST-9733111 and a fellowship from the David and Lucile Packard Foundation."
    },
    "0601/gr-qc0601062_arXiv.txt": {
        "abstract": "We study the collapse towards the gravitational radius of a macroscopic spherical thick shell surrounding an inner massive core. This overall electrically neutral \\emph{macroshell\\/} is composed by many nested $\\delta$-like massive \\emph{microshells\\/} which can bear non-zero electric charge, and a possibly non-zero cosmological constant is also included. The dynamics of the shells is described by means of Israel's (Lanczos) junction conditions for singular hypersurfaces and, adopting a \\emph{Hartree\\/} (mean field) approach, an effective Hamiltonian for the motion of each microshell is derived which allows to check the stability of the matter composing the macroshell. We end by briefly commenting on the quantum effects which may arise from the extension of our classical treatment to the semiclassical level. ",
        "introduction": "\\setcounter{equation}{0} \\label{sec:introduction} The dynamics of collapsing bodies is a subject which attracted much attention amongst physicists since the very first formulation of General Relativity. However, the general problem of gravitational collapse is extremely complicated and simplified models have therefore been constructed which allow one to describe the dynamics at least for the simplest configurations (for a review of analytical solutions, see Ref.~\\cite{adler}). In this perspective, a lot of effort has been dedicated to the study of the collapse of thin spherical shells (see, e.g.~Refs.~\\cite{grav,Israel,Pim}), for they represent a quite general and simple model and can be regarded as the starting point for the analysis of more complicated systems. In recent years, attention has been focused on quantum effects that can arise from strong gravitational fields characterizing the final stages of the collapse and several approaches have been attempted (for a very limited list, see, e.g.~Refs.~\\cite{Farhi,Balbi,Hajicek}). \\par In particular, a formalism was introduced in Ref.~\\cite{shellCQG} which allows one to study the semiclassical behaviour of the shell by adopting a ``minisuperspace'' approach~\\cite{mini} to its dynamics very much in the same spirit as that previously employed for the Oppenheimer-Snyder model in Ref.~\\cite{OS}. Yet, many aspects of the collapse still need to be investigated at the classical level, which may alter the quantum (semiclassical) picture in a substantial way. For example, the issue of whether the radial degree of freedom of the shell can be described by a quantum wave function strictly depends on the fact that the matter composing the shell be confined within a very small thickness (i.e.~of the order of the \\emph{Compton\\/} wavelength of the elementary particles composing the shell) around the shell mean radius, so that the assumption of a coherent state is reliable. We refer to this issue as the ``stability'' problem for thick shells. In fact, a realistic description of the collapse should involve finite quantities (that is, a shell with finite thickness) whereas, already for the classical dynamics, the very use of standard junction equations~\\cite{Israel} requires the validity of the \\emph{thin shell limit\\/} (i.e.~the shell's thickness be much smaller than its gravitational radius). \\par In this paper we address the classical dynamics of a collapsing spherical shell (which we call \\emph{macroshell\\/}) of proper mass $M$ with an inner massive core of mass $M_c$ and study the conditions under which the matter composing the shell remains confined around its mean radius. In order to verify this property, we discretize the continuous distribution of matter and consider the macroshell as if composed by a large number $N$ of homogeneous nested spherical $\\delta$-like \\emph{microshells\\/}~\\cite{shellPRD}. Exploiting the spherical symmetry, we can adopt a Schwarzschild-like coordinate system and the motion of the microshells can be described by the usual ``areal'' radial coordinates $r_i$ ($i=1,\\ldots,N$), ordered so that $r_1<r_2<...r_N$. Therefore, the thickness of the macroshell is given by $\\delta=(r_N-r_1)$ and its mean radius is $R=(r_N+r_1)/2$. Each microshell will be assumed to have the same proper mass $m_i=m$ and a (possibly zero) electric charge $Q_i$. We will limit our study outside the gravitational radius of the system, $R_H=2\\,G\\,M_s/c^2$, where $M_s$ is the total ADM mass, so that the radius of the innermost microshell $r_1$ will always be greater than $R_H$. \\par To summarise, the model will be constructed assuming the following conditions: \\begin{enumerate} \\item $M_c\\gg M$. The mass of the core is much bigger than the macroshell's, so that the total mass-energy of the system is approximately $M_s\\simeq M_c$; \\label{1con} \\item $m\\ll Nm=M$. Each microshell is much lighter than the macroshell; \\item $(R>)\\,r_1\\gg \\delta$. The inner radius of the macroshell (and therefore the radius of each microshell) is much larger than its thickness. This implies that the macroshell thickness also remains much smaller than $R_H$ and this condition will limit the temporal range of our analysis; \\item $\\displaystyle{\\sum_{i=1}^{N}Q_i=0}$. The macroshell is overall electrically neutral; \\item $R(0)\\gg R_H$ and $\\displaystyle{\\left.\\frac{\\rmd R}{\\rmd t}\\right|_{t=0}\\simeq 0}$. The collapse starts in the weak field regime. One can therefore take $t$ as the Schwarzschild time, although it will then be more convenient to use the (micro)shell proper times (for a detailed comparison, see the Appendix~B in Ref.~\\cite{shellPRD}). \\label{lcon} \\end{enumerate} \\par As mentioned above, the classical equations of motion for a thin shell in General Relativity can be obtained from the standard Israel's junction equations. In the present paper, we shall adopt a minisuperspace approach in which the embedding metrics (interior and exterior with respect to the shell) are held fixed and chosen to be particular solutions of the (bulk) Einstein equations, and derive the shell's dynamical equations from a variational principle applied to an effective action. This (equivalent) treatment will allow us to take into account the matter internal degrees of freedom and can be extended to the semiclassical level along the lines of Refs.~\\cite{shellCQG,shellPRD}. In fact, the equations thus obtained are equivalent to the junction conditions, with the exception of one more equation which has no counterpart as a junction condition and represents a secondary constraint for the time variation of the gravitational mass of the shell. The relevant dynamical equation is a first order ordinary differential equation for the shell radius which contains ADM parameters~\\cite{ADM} (those defining the inner and outer metrics). In order to determine the shell trajectory, one can fix the initial shell radius, but the initial velocity is then uniquely determined {\\em if\\/} the ADM parameters were also fixed. Alternatively (and more sensibly), one can set the initial radius and velocity and adjust the ADM parameters correspondingly. \\par The task of extending the above analysis to a system of (micro)shells would be extremely involved in either cases. If we choose to set the initial radii and velocities, then the ADM parameters will have to be reset each time two shells cross. Alternatively, if we set the ADM parameters once and for all, it is possible that the system runs into some spurious (cusp or ``superficial'') singularity when two shells cross~\\cite{adler,singh}. In light of a future extension to the semiclassical level, we shall therefore find it more convenient to derive an effective Hamiltonian for the motion of a single microshell in the mean gravitational field generated by the core and all the other microshells treated as a whole. This will allow us to obtain a one-particle effective potential which describes the tidal forces the microshell is subject to, and, in particular, we will see that the binding potential varies with time and confines the microshells within an ever decreasing thickness (at least sufficiently away from the gravitational radius of the system). \\par We will work in units in which the speed of light $c=1$ but will explicitly display the gravitational (Newton) constant $G$. ",
        "conclusions": "\\setcounter{equation}{0} \\label{conc} We have examined the collapse of a macroscopic self-gravitating spherical shell of matter in the presence of a massive core in the framework of General Relativity. The (macro)shell is assumed to be formed by a large number of thin (micro)shells initially bundled within a very short thickness. The system is then shown to remain classically stable, since the effective potential which drives the microshells becomes more and more binding as the macroshell approaches the gravitational radius of the system (but remains sufficiently away from it so that our approximations $R-R_H\\gg\\delta$ and $|\\dot R|<1$ hold). Further, including a (reasonably valued) cosmological constant does not alter this behaviour, nor does the presence of electric charges carried by the microshells. \\par A semiclassical analysis may however change the above conclusion, since the bound states of microshells are just meta-stable and there may be significant changes in their lifetimes (or other, non-adiabatic, effects) as the macroshell collapses. In Refs.~\\cite{shellPRD,ac}, it was shown that the microshells can indeed be described by means of localized quantum mechanical wave functions, however in the linear approximation (first order in $|\\bar r|/R$) for the effective potential. We have now found that, for sufficiently large values of $|\\bar r|$ (but still much smaller than $R$), the complete effective potential reaches maxima and then decreases, thus suggesting a non-vanishing probability for the microshells to tunnel out. It was also found in Refs.~\\cite{shellPRD,ac} that, upon coupling the microshell wave functions to a radiation field, the shell spontaneously emit during the collapse as a non-adiabatic effect. A further analysis of such semiclassical effect in light of the present results is therefore in order, which could also affect the usual thermodynamical description~\\cite{thermo}. \\par Let us conclude by mentioning that inclusion of (or a possible relation with) the Hawking radiation~\\cite{hawk} as the shell approaches the gravitational radius still remains to be investigated."
    },
    "0601/hep-ph0601089_arXiv.txt": {
        "abstract": "We examine the effect of a prior that favours low values of fine-tuning on Bayesian multi-dimensional fits of the constrained minimal supersymmetric standard model (CMSSM or mSUGRA) to current data. The dark matter relic density, the anomalous magnetic moment of the muon and the branching ratio of $b \\rightarrow s \\gamma$ are all used to constrain the model via a Markov Chain Monte Carlo sampler. As a result of the naturalness prior, posterior probability distributions skew towards lighter higgs and sparticle masses, the effect being most pronounced in the gaugino sector. Interestingly, slepton masses are an exception and skew towards heavier masses. The lightest CP-even Higgs $h^0$-pole annihilation mechanism becomes allowed at the 2$\\sigma$ level for the latest combination of measurements of $m_t=172.7 \\pm 2.9$ GeV, provided we allow for a theoretical error in the prediction of its mass $m_{h^0}$.  $m_{h^0}$ is constrained to be less than 120 GeV at the 95$\\%$ C.L. Probing the branching ratio of $B_s \\rightarrow \\mu^+ \\mu^-$ to the level of 2$\\times 10^{-8}$, as might be achieved by the Tevatron experiments, would cover 32$\\%$ of the probability density, irrespective of which of the two priors is used. ",
        "introduction": "Although weak-scale supersymmetry (SUSY)~\\cite{haber} arguably offers the best candidate for physics beyond the Standard Model, it potentially suffers from a fine-tuning problem~\\cite{ftrefs}. Radiative electroweak symmetry breaking conditions imply that the mass of the $Z^0$-boson $M_Z$ is related to the superpotential $\\mu$-term, the two soft supersymmetry breaking Higgs mass parameters $m_{H_{1,2}}$ and the ratio of the Higgs vacuum expectation values $v_2/v_1=\\tan \\beta$ by \\begin{equation} \\frac{M_Z^2}{2} = \\frac{m_{H_1}^2 - m_{H_2}^2 \\tan^2 \\beta}{\\tan^2 \\beta - 1} -\\mu^2   \\label{rewsb} \\end{equation} at tree level~\\cite{haber}. In the following numerical calculations, the full one-loop MSSM corrections are added to this relation. Experiments at LEP2 and the Tevatron have placed increasingly severe bounds upon sparticle and Higgs masses~\\cite{ftrefs,PDBook} and in constrained models they push up the minimum value of $\\mu$ to several hundred GeV. In order to reproduce the observed value of $M_Z=91$ GeV, there then must be a large degree of cancellation between the {\\em a priori} unrelated terms on the right hand side of Eq.~\\ref{rewsb}. This is considered to be unnatural. In order to quantify the amount of fine-tuning, one can construct~\\cite{ftrefs} a fine-tuning measure $c_i$ for a parameter of the model $p_i$ by noting that $M_Z$ is unstable to small changes of $p_i$ as a result of the fine-tuning: \\begin{equation} c_i \\equiv \\left|\\frac{\\partial \\ln M_Z}{\\partial \\ln p_i}\\right|, \\qquad c\\equiv \\mbox{max} \\{c_i \\}. \\label{fineT} \\end{equation} $c_i$ then measures the fractional change in $M_Z$ (the partial derivative is calculated while {\\em not} using the constraint in Eq.~\\ref{rewsb}) produced by some small fractional  change in the parameter $p_i$. We shall here use the largest $c_i$ to quantify the amount of fine-tuning $c$ for a given point in parameter space.  It is our purpose in this paper to study the effect of a prior probability distribution that favours a low value of $c$ in combined fits to the CMSSM\\@. We wish to extract any collider implications, as well as implications for the dark matter annihilation mechanism. The MSSM contains several good candidates for cold dark matter, the most studied being the lightest neutralino $\\chi_1^0$. The assumption that thermally produced  neutralinos make up the dominant component of dark matter leads to strict constraints upon the parameter space of the MSSM\\@. We shall here take this assumption, as well as the assumption of R-parity (which provides stability for the lightest neutralinos). In order to keep the dimensionality of the parameter space down, we shall further specialise to the CMSSM. The CMSSM unifies all SUSY breaking scalar mass parameters to $m_0$, the Standard Model gaugino mass parameters to $M_{1/2}$ and the SUSY breaking trilinear scalar couplings to $A_0$ at a high energy scale, commonly taken to be the scale of electroweak gauge unification $M_{GUT}$.  Our choice of parameters from which to calculate the fine-tuning includes the SUSY breaking Higgs bilinear $B$: \\begin{equation} p_i \\in \\{ m_0, A_0, M_{1/2}, B, \\mu \\}. \\end{equation} $B$ is used in preference to $\\tan \\beta$ since it directly appears in the Lagrangian and is therefore considered to be more fundamental. Our definition of $c$ does not include the sensitivity to the GUT-scale top Yukawa coupling $h_t$. We can argue about whether to include it or not in the definition since its value is determined by aspects of flavour physics in the model, an orthogonal aspect to the SUSY breaking aspects studied in the present paper. Were we to include it in the definition of $c$, a strong suppression of the high $m_0$ regime would be the likely result of the naturalness prior, since $c_{h_t}$ is very high there~\\cite{Feng:1999mn}.  In ref.~\\cite{Giusti:1998gz}, exclusion limits from LEP1 and LEP2 were placed upon the CMSSM and a probability of $1/c$ was applied to the remaining parameter space. This resulted in probability distributions for sparticle masses and various observables. This work is similar in spirit to the ingredient of the naturalness prior that we will follow. A more common approach in the literature involves applying separate 95$\\%$ C.L. constraints upon a two-dimensional sub-space of the model, see recent examples in ref.~\\cite{classic}. The advantage of this simple approach is that it is easy to discern which constraints act on different parts of parameter space. On the other hand, it was argued in Ref.~\\cite{myold} that a more complete approach would perform a combined likelihood fit in the full dimensionality of the parameter space, including variations of important Standard Model input parameters. In fact, the feasibility of the approach had already been demonstrated in a four-dimensional subspace~\\cite{gondolo}. This was achieved by utilising the Metropolis algorithm in a Markov Chain Monte Carlo (MCMC) algorithm~\\cite{mackay}. The MCMC approach to parameter sampling has a calculation time that scales linearly with the number of dimensions, rendering it very useful for investigating parameter spaces with dimensionality greater than 3. We point the interested reader to refs.~\\cite{myold,mackay} for details on the algorithm, but the important point for the present paper is that the algorithm provides a list (or ``chain'') of points in parameter space. The density of points that it produces is proportional to some specified function of the parameters. In reference~\\cite{myold}, the function chosen was the likelihood, the probability distribution function (pdf) of reproducing data $d$ given a point in CMSSM parameter space $m$: ${\\mathcal L}=p(d|m)$. One ideally wishes to determine the posterior pdf that the model point $m$ is correct given the data $d$: $P(m)=p(m|d)$. The relative posterior between two different points $m_{1,2}$ in parameter space can be determined provided one assumes a prior pdf for the model parameters $p(m)$: \\begin{equation} \\frac{P(m_1)}{P(m_2)} = \\frac{p(d|m_1)}{p(d|m_2)} . \\frac{p(m_1)}{p(m_2)}. \\end{equation} Thus, the likelihood distribution is proportional to the posterior distribution for the case of flat prior distributions, i.e. when $p(m_1)/p(m_2)$ is a constant across the relevant parameter space. For a brief discussion on a Bayesian interpretation of the prior probability distributions, see the Appendix.   Here, we will go beyond the MCMC analysis of Ref.\\cite{myold} by using a plausible {\\em naturalness}-favouring prior proportional to $1/c$. This will then disfavour regions of parameter space with high fine-tuning, arguably a commendable effect. We will then compare and contrast the resulting probability distributions of observables, parameters and sparticle masses with the flat-prior case.  Having obtained a chain of points whose density is proportional to their posterior assuming some prior, we could in principle just re-weight the probability of each point by some new prior and re-plot all of the results. In the present case, we would for example use the flat-prior sample, but reweight the probability of each point by $1/c$, then re-bin any results. The only disadvantage to this approach is that if previously unlikely regions of parameter space now have high probabilities, they will not have been sampled very often in the initial Markov chain run. Thus the statistical fluctuations in this region will be large. If the prior doesn't make a lot of difference to the posterior, this degradation of the statistics may not matter much. Here though, we will find a significant difference in the results due to different priors, implying that a simple re-weighting procedure would suffer from poor statistics. We have actually re-run the MCMC with the new prior, which circumvents any potential statistics problem. ",
        "conclusions": "By employing a simple MCMC, we have investigated the effect of a naturalness prior on CMSSM combined fits. The prior has a significant effect upon the fits, indicating that current data are either not sufficiently precise or not numerous enough to render posterior probabilities insensitive to the prior. Our results bear some similarity to those of ref.~\\cite{Giusti:1998gz}, which also effectively used a naturalness prior, but did not use a likelihood, only LEP 1,2 exclusion limits. However, features due to the dominant dark matter annihilation mechanism are prominent in our results. The naturalness prior enhances the posterior probability of the light CP-even higgs annihilation region of CMSSM parameter space, which comes back inside the 95$\\%$ C.L. The relative probabilities of $A^0-$pole and stau co-annihilation regions decrease. Probability distributions for sparticle masses are skewed somewhat towards lower values, except for the sleptons, which show a slight skew towards heavier masses. The naturalness prior implies an enhanced probability for the lightest neutralinos and charginos to be accessible at a future linear collider. The prospects for evidence of $B_s \\rightarrow \\mu^+ \\mu^-$ at the Tevatron remain good at 32$\\%$ whichever prior is used for 8 fb$^{-1}$ of integrated luminosity. Rather light CP-even Higgs masses are preferred, the 95$\\%$ upper C.L. limit being only 120 GeV. This is unfortunately a somewhat difficult regime for $h^0$ discovery by the LHC, and implies that several years of data-taking will be required~\\cite{ATLASTDR}. Ref.~\\cite{gunion} found that fine-tunings of at least 180 are required by the LEP2 bound $m_{h^0}>114$ GeV, $\\tan \\beta=10$. Our results indicate fine-tunings in the range of 20-30 are still feasible. The results in the present paper include wider variations of more parameters (particularly in $A_0$), but the dominant difference to Ref.~\\cite{gunion} is that here we allow a 3 GeV theoretical error upon the $m_{h^0}$ prediction, which extends the fine-tuning range of the fitted points downwards. The random scan of Ref.~\\cite{gunion} was also very sparse in the parameters and so there is a distinct possibility of just not finding the less fine-tuned points.  The amount of skew of the probability distributions will be modified by changing the prior assumption. For example, choosing a prior that is proportional to $1/c^2$ instead of $1/c$ (thus effectively counting the naturalness twice rather than just once) would produce more skew in the distributions.  There are many possible future directions for the use of MCMC algorithms in fits to supersymmetric models. One obvious direction is the possibility of considering different supersymmetry breaking models, for example relaxing some of the universality assumed in the CMSSM\\@. Another possible direction would be to include electroweak observables in the fit. Yet another would be to explore the $\\mu<0$ part of parameter space and work out its probability normalisation with respect to the $\\mu>0$ part. The MCMC procedure carried out here was at the limit of CPU capabilities available to the author, but finding a more efficient algorithm could allow for an enlargement of the parameter space considered. In particular, going to high values of $m_0$ would be useful in order to fully investigate the focus point regime~\\cite{leszek}."
    },
    "0601/astro-ph0601269_arXiv.txt": {
        "abstract": "IceCube is a 1 km$^3$ neutrino observatory being built to study neutrino production in active galactic nuclei, gamma-ray bursts, supernova remnants, and a host of other astrophysical sources. High-energy neutrinos may signal the sources of ultra-high energy cosmic rays.  IceCube will also study many particle-physics topics: searches for WIMP annihilation in the Earth or the Sun, and for signatures of supersymmetry in neutrino interactions, studies of neutrino properties, including searches for extra dimensions, and searches for exotica such as magnetic monopoles or Q-balls.  IceCube will also study the cosmic-ray composition.  In January, 2005, 60 digital optical modules (DOMs) were deployed in the South Polar ice at depths ranging from 1450 to 2450 meters, and 8 ice-tanks, each containing 2 DOMs were deployed as part of a surface air-shower array.  All 76 DOMs are collecting high-quality data. After discussing the IceCube physics program and hardware, I will present some initial results with the first DOMs. ",
        "introduction": "Neutrinos are strongly penetrating particles that can be used to search for the sources of high-energy cosmic rays and study many energetic astrophysical phenomena.  Except at the very highest energies (above 50 EeV), high-energy protons and nuclei are bent in the interstellar magnetic fields, and high-energy photons are absorbed by interactions with the $3^0$K microwave background and infrared photons from red-shifted starlight.  Only neutrinos can provide a clear view of the sky above 50 TeV.  The primary purpose of IceCube is to map out this sky \\cite{PDD}; many calculations using a wide variety of approaches find that a volume of at least 1 km$^3$ detector is required to observe neutrino sources \\cite{icsense}.  Neutrinos are expected as `by-products' of hadronic acceleration, and may be observed from active galactic nuclei (AGNs), gamma-ray bursters (GRBs), and possibly supernova remnants.  The completed IceCube detector will collect data on about 100,000 atmospheric $\\nu$/year \\cite{atmosphere}.  These neutrinos can be used to measure the neutrino-nucleon cross sections (through absorption in the earth) and search for new forms of neutrino oscillations, such as those predicted by some models of quantum gravity \\cite{qg}).  IceCube will also search for signs of supersymmetry (some parameter sets with high mass scales lead to pairs of upward going particles in the detector) \\cite{SUSY}, neutrinos from WIMP annihilation in the Earth or the Sun \\cite{WIMPs}, and for exotic particles like magnetic monopoles, Q-balls, and the like.  IceCube also serves as a supernova monitor; a supernova in our galaxy should increase the singles rates of all of the in-ice phototubes \\cite{supernova}. ",
        "conclusions": "In 2004/5 the IceCube collaboration deployed a string of 60 DOMs at the South Pole, along with 8 ice-tanks containing an additional 16 DOMs.  All 76 DOMs are working well, and the system shows time resolutions of $\\approx 2$ nsec.  In 2005/6, we expect to deploy an additional 8-12 strings, moving toward completion of the 80-string array by 2010.  \\begin{theacknowledgments}  This research was supported by the Deutsche Forschungsgemeinschaft (DFG); German Ministry for Education and Research; Knut and Alice Wallenberg Foundation, Sweden; Swedish Research Council; Swedish Natural Science Research Council; Fund for Scientific Research (FNRS-FWO), Flanders Institute (IWT), Belgian Federal Office for Scientific, Technical and Cultural affairs (OSTC), Belgium. UC-Irvine AENEAS Supercomputer Facility; University of Wisconsin Alumni Research Foundation; U.S. National Science Foundation, Office of Polar Programs; U.S. National Science Foundation, Physics Division; U.S. Department of Energy.  \\end{theacknowledgments}"
    },
    "0601/astro-ph0601575_arXiv.txt": {
        "abstract": "{The Vela supernova remnant (SNR) is a complex region containing a number of sources of non-thermal radiation. The inner section of this SNR, within 2 degrees of the pulsar \\object{PSR B0833$-$45}, has been observed by the \\hess\\ \\gr\\ atmospheric Cherenkov detector in 2004 and 2005. A strong signal is seen from an extended region to the south of the pulsar, within an integration region of radius $0.8\\dg$ around the position ($\\alpha = 08^{h} 35^{m} 00^{s}$, $\\delta\\ = -45\\dg\\ 36\\arcmin$ J2000.0).  The excess coincides with a region of hard X-ray emission seen by the ROSAT and ASCA satellites. The observed energy spectrum of the source between 550 GeV and 65 TeV is well fit by a power law function with photon index $\\Gamma = 1.45 \\pm 0.09\\stat\\ \\pm 0.2\\sys$ and an exponential cutoff at an energy of $13.8 \\pm 2.3\\stat\\ \\pm 4.1\\sys$ TeV. The integral flux above 1 TeV is $(1.28 \\pm 0.17\\stat\\ \\pm 0.38\\sys) \\times 10^{-11}\\ \\iflux$.  This result is the first clear measurement of a peak in the spectral energy distribution from a VHE \\gr\\ source, likely related to inverse Compton emission. A fit of an Inverse Compton model to the \\hess\\ spectral energy distribution gives a total energy in non-thermal electrons of $\\sim$$2 \\times 10^{45}$ erg between 5 TeV and 100 TeV, assuming a distance of 290 parsec to the pulsar. The best fit electron power law index is $2.0$, with a spectral break at 67 TeV. ",
        "introduction": "The region surrounding the Vela supernova remnant (SNR) is very well studied across the electromagnetic spectrum and contains a number of complex objects, including the SNR \\object{RX J0852.0$-$4622}, which has been previously detected in the very high energy \\gr\\ (VHE) range \\citep{katagiri05,aharonian05b}.  The Vela SNR itself, at a distance estimated to be $\\sim$290~pc \\citep{dodson03,caraveo01}, extends over a diameter of $\\sim$$8\\dg$. It is the nearest SNR to contain a young active pulsar, \\object{PSR B0833$-$45}, with a period of 89 ms and a period derivative \\.P of $1.25 \\times 10^{-13}\\ \\textrm{s}\\ \\textrm{s}^{-1}$. This implies a spin-down luminosity of $7 \\times 10^{36}\\ \\textrm{erg}\\ \\textrm{s}^{-1}$ and age of 11,000 years.  Observations by Chandra \\citep{helfand01} clearly show the torus-like morphology of the compact X-ray nebula surrounding the pulsar and have allowed its rotation axis to be inferred. The plane of the torus (i.e. pulsar equator) appears to intersect the plane of the sky at a position angle of $40.6\\dg$ relative to North \\citep{ng04}, and its X-ray spectral index of $\\sim$1.5 \\citep{mangano05} is typical of pre-cooled pulsar wind torii.  The SNR also contains a number of regions of non-thermal emission, including those labelled by \\citet{rishbeth58} from radio observations as Vela X, Vela Y and Vela Z (which is part of the shell of \\object{RX J0852.0$-$4622}).  A diffuse emission feature has been detected by ROSAT \\citep{markwardt95} in hard X-rays ($0.9$--$2.0$~keV), coinciding with the centre of the Vela X region.  It was first suggested that this feature, which is aligned closely with a filament detected at radio wavelengths, corresponds to the outflow jet from the pole of the pulsar \\citep{frail97}. However, the Chandra observations showed that this feature lies along the extension of the pulsar equator, although bending to the southwest.  The study of Vela X is important from two perspectives: It was the first middle-aged pulsar wind nebula (PWN) to be detected, thereby introducing the concept of PWN evolution \\citep{weiler80}. Also, it served as the first prototype for offset PWN as a result of SNR expansion into an inhomogeneous medium \\citep{blondin01}, of which \\object{G18.0-0.7} served as a second example \\citep{gaensler03}. This second PWN was also recently identified with VHE source \\object{HESS J1825--137} \\citep{aharonian05c}.  Here we discuss observations of the Vela pulsar and the centre of the Vela SNR with \\hess, including the Vela~X feature. A region of extended \\gr\\ emission is detected, to the south of the pulsar position (previously reported by \\citet{khelifi05}.  This region was previously observed by the CANGAROO collaboration \\citep{yoshikoshi97} and emission of VHE \\grs\\ was claimed from a source $0.13\\dg$ to the southeast of the pulsar. However doubt was later cast on the significance of this detection \\citep{dazeley01}. A recent paper from the CANGAROO collaboration \\citep{enomoto05} retracts the claimed detection, and suggests some evidence for an excess coincident with the \\hess\\ signal.   \\section {Observations with \\hess}  \\hess\\ is a system of four atmospheric Cherenkov telescopes designed to look for astrophysical \\gr\\ emission above $\\sim$100 GeV; a review of the detector is given by \\citet{hinton04}. The Vela region was observed between January and March 2004 with the complete \\hess\\ array. In total 10.3 hours of data were obtained, after standard data quality selection, at zenith angles between 20\\dg\\ and 40\\dg. These observations were taken using a method (\\emph{wobble} mode) whereby the source is offset by a small angular distance from the centre of the field of view, alternating between 28 minute runs in the positive and negative declination (or right ascension) directions.  The observations were made at offsets of 0.5\\dg\\ in declination from the position of the Vela pulsar, ($\\alpha = 08^{h} 35^{m} 20^{s}$, $\\delta\\ = -45\\dg\\ 10'\\ 35''$ J2000.0), which is referred to as position I for the purposes of this paper.  Following a detection of extended \\gr\\ emission to the south of the pulsar in the 2004 dataset, further observations were made in 2005 surrounding the position $\\alpha = 08^{h} 35^{m} 00^{s}$, $\\delta\\ = -45\\dg\\ 36'$ (J2000.0), here referred to as position II, which was measured as the centre of gravity of the excess. These were also taken in wobble mode, and the mean offset for the 2005 data is 0.5\\dg. After data quality selection 7.2 hours of observations were available for analysis, giving total observation live time over the two years of 16.4 hours, after dead time correction (approximately 10\\%). The mean zenith angle for the complete set of observations is $30.2\\dg$. ",
        "conclusions": "The new VHE source reported here, \\object{HESS J0835--455}, is situated to the south of the pulsar and the compact X-ray nebula (as seen by Chandra). The integral flux is estimated to be $\\sim$50$\\%$ of that of the Crab nebula above 1 TeV.  As the distance to the pulsar is well measured, one can estimate the size of the emission region seen by \\hess\\ to be 5.1 parsec (full length for 68\\% containment) along the major axis by 3.8 parsec (full width).  The luminosity of the emission region in the energy range from 550 GeV to 65 TeV can be estimated to be $L = 9.9 \\times 10^{32}\\ \\textrm{erg}\\ \\textrm{s}^{-1}$, using the power law fit to the spectrum with the exponential cutoff.  \\object{HESS J0835--455} appears to be spatially coincident with the X-ray ($0.4$--$2.4$ keV) emission as seen by ROSAT (shown in Figure \\ref{fig:skymap_narrow}). It has been suggested \\citep{blondin01} that the Vela X feature corresponds to the pulsar wind nebula, displaced to the south by the unequal pressure of the reverse shock from the SNR. This hypothesis is consistent with the \\hess\\ observations which demonstrate conclusively for the first time that this feature emits non-thermal radiation. A similar explanation has been proposed for the emission seen by \\hess\\ (\\object{HESS J1825--137}) close to the pulsar wind nebula G18.0--0.7 \\citep{aharonian05c}.  The size of the feature described by \\citet{markwardt95} has been measured as 45 $\\times$ 12 arcmin$^{2}$, whereas the VHE intrinsic size is 58 $\\times$ 43 arcmin$^{2}$. An extension of the X-ray feature to the southwest was suggested by \\citet{lu00}, however this may be unrelated thermal emission, a suggestion which is supported by the fact that no excess is seen in this region by \\hess  \\begin{figure} \\centering \\includegraphics[width=0.5\\textwidth, height=7cm]{Hl202_fig4.eps} \\caption{Spectral energy distribution for the \\hess\\ and ASCA spectral measurements \\citep{markwardt97}. The two alternative X-ray spectra are described in the text.  The fitted inverse Compton emission (solid line) from the Vela X region is shown, given the electron energy distribution described in the text. The predicted synchrotron flux is shown for three possible magnetic field levels, between $2\\ \\mu$G and $8\\ \\mu$G.} \\label{fig:sed} \\end{figure}  The larger VHE size probably indicates that synchrotron cooling in the X-ray domain was (or still is) important.  This is confirmed by the detection of a cooled X-ray photon index of $\\sim$2.1 \\citep{markwardt97} in the emission region, which is consistent with the uncooled photon index of 1.5 closer to the source. Figure \\ref{fig:sed} shows a spectral energy distribution (SED) for Vela X, including the \\hess\\ result reported here. Two possible X-ray spectra are shown, the lower including only the 'Cocoon' feature defined in the above paper, while the upper flux includes the emission from the 'SNR' component extrapolated over the \\hess\\ integration region.  A simple one-zone model of inverse Compton emission from interactions with the cosmic microwave background (CMBR) is also shown.  The \\hess\\ observations show the expected peak in the IC emission; this result is the first clear measurement of such a peak at VHE energies. The total energy content in non-thermal electrons between 5 TeV and 100 TeV is $2.2 \\times 10^{45}$ erg, for a distance of 290 parsec.  The best fit electron power law index is $2.0$, with a break at 67 TeV and a post-break index of 9. The \\rchisq\\ of this fit is 10.2/9.  In Figure \\ref{fig:sed} three possible synchrotron spectral energy distributions are shown, for magnetic field levels in the emission region from $2\\ \\mu$G to $8\\ \\mu$G. The synchrotron distributions are not constrained by the X-ray data, a combined fit on the X-ray and \\gr\\ data would more closely reproduce the measured X-ray spectral slope. A magnetic field of this magnitude may be expected in the scenario of a nebula displaced from the pulsar by an asymmetric shock.  This result demonstrates the usefulness of \\hess\\ observations in clarifying the complex morphology of this source. VHE observations of inverse Compton scattering of the CMBR allow direct inference of the spatial and spectral distribution of non-thermal electrons in a PWN, independent of contamination by thermal emission or variations in the local magnetic field."
    },
    "0601/astro-ph0601096_arXiv.txt": {
        "abstract": "In recent years evidence has accumulated suggesting that the gas in galaxy clusters is heated by non-gravitational processes. Here we calculate the heating rates required to maintain a physically motived mass flow rate, in a sample of seven galaxy clusters. We employ the spectroscopic mass deposition rates as an observational input along with temperature and density data for each cluster. On energetic grounds we find that thermal conduction could provide the necessary heating for A2199, Perseus, A1795 and A478. However, the suppression factor, of the clasical Spitzer value, is a different function of radius for each cluster. Based on the observations of plasma bubbles we also calculate the duty cycles for each AGN, in the absence of thermal conduction, which can provide the required energy input. With the exception of Hydra-A it appears that each of the other AGNs in our sample require duty cycles of roughly $10^{6}-10^{7}$ yrs to provide their steady-state heating requirements. If these duty cycles are unrealistic, this may imply that many galaxy clusters must be heated by very powerful Hydra-A type events interspersed between more frequent smaller-scale outbursts. The suppression factors for the thermal conductivity required for combined heating by AGN and thermal conduction are generally acceptable. However, these suppression factors still require `fine-tuning` of the thermal conductivity as a function of radius. As a consequence of this work we present the AGN duty cycle as a cooling flow diagnostic. ",
        "introduction": "The radiative cooling times of the hot, X-ray emitting gas in the centres of many galaxy clusters may be as short as $10^{6}$ years. In the classical model \\citep[see e.g.][for a review]{fab94}, where radiative losses occur unopposed, a cooling flow develops in which the gas cools below X-ray temperatures and accretes onto the central cluster galaxy where it accumulates in molecular clouds and subsequently forms stars. However, studies have demonstrated both a lack of the expected cold gas at temperatures well below 1-2keV \\citep[e.g.][]{edge01} and that spectoscopically determined mass deposition rates are a factor of 5-10 less than the classically determined values \\citep{voigt04}. These findings suggest that the cluster gas is being reheated in order to produce the observed minimum temperatures and the reduction in excess star formation compared to expectations. As yet it is unclear whether this heating occurs continuously or periodically.   There is observational evidence that the temperature profiles of galaxy cluster atmospheres are well described by the same mathematical function across a range in redshift \\citep[e.g.][]{allen01}. Given such similarity it is possible that the atmospheres of galaxy clusters are in a quasi-steady state and that observable parameters such as the radial temperature profile do not vary signicantly over the lifetime of a cluster. This universal temperature profile would be difficult to sustain if the density profiles varied significantly with time. In addition, if large quantities of gas were deposited in the central regions one might expect that the density profiles would be much more centrally peaked than observations suggest. The culmination of this argument is that, at least in the central regions, the inward flow rate of mass must be roughly independent of radius to ensure that mass is not deposited at any particular location which would significantly alter the density profile.  One particular problem with this is that all of the mass flowing in this region must be deposited at the cluster centre. Therefore, unless the mass flow rate is small this would result in large quantities of gas at the cluster centre.   In contrast, spectroscopically determined mass deposition rates suggests that mass $\\it{is}$ deposited at a roughly constant rate throughout the cluster, at resolved radii \\citep[][]{voigt04}. However, a roughly constant mass deposition rate is indicative of a mass flow rate that is roughly proportional to radius, rather than constant as the density profiles suggest. Yet, if such a flow were persistent this proportionality would result in density profiles which contradict the observational evidence by being more centrally peaked.   If we are to reconcile the implications of the density profiles and mass deposition rates then one possible explanation is that the mass flow rate exhibits two different asymptotic properties. That is, near the cluster centre or within the central galaxy the mass flow rate should be roughly constant so that the density profiles are essentially left unchanged, while the mass deposited at the cluster centre is sufficiently small to agree with observations. At larger radii the mass flow rate should still be proportional to radius and satisfy the observations on resolvable scales which imply a constant mass deposition rate. The relationship between the mass flow rates and the mass deposition rates are defined section 3.   The two main candidates for heating cluster atmospheres are: Active Galactic Nuclei (AGN) \\citep[e.g.][]{tabor93,bub01,nature,brueggen03} and thermal conduction \\citep[e.g.][]{gaetz, zakamska03, voigt02, voigt04}. Heating by AGN is thought to occur through the dissipation of the internal energy of plasma bubbles inflated by the AGN at the centre of the cooling flow. Since these bubbles are less dense than the ambient gas, they are buoyant and rise through the intracluster medium (ICM) stirring and exciting sound waves in the surounding gas \\citep[e.g.][]{ruszbegel02,fabpers02}. This energy may be dissipated by means of a turbulent cascade, or viscous processes if they are significant. Deep in the central galaxy other processes such as supernovae and stellar winds will also have some impact on the ambient gas.  However, AGN are only periodically active which could result in similarly periodic heating rates, thus making the possibility of a totally steady state unlikely. In this case one could imagine a scenario in which a quasi steady-state is possible where temperature and density profiles oscillate around their average values. It is also possible that dissipation of the plasma bubbles may occur over timescales longer than the AGN duty cycle thus providing almost continous heating \\citep[e.g.][]{reynolds05}.   Thermal conduction may also play a significant role in transferring energy towards central regions of galaxy clusters given the large temperature gradients which are observed in many clusters. In fact, several authors \\citep[e.g.][]{voigt02,voigt04} have shown that on energy grounds alone it may be possible for thermal conduction to provide the necessary heating in some clusters.  However, the exact value of the thermal conductivity remains uncertain. The theoretical value for the thermal conductivity of a fully ionised, unmagnetised plasma is calculated by \\cite{spitzer}. Magnetic fields, which are thought to exist in the cluster gas, based on Faraday rotation measures of clusters \\citep[e.g.][]{carilli02}, may greatly alter this value. The standard method by which the unknown effects of magnetic fields are taken into account is to include a suppression factor, $f$, in the Spitzer formula, which indicates the actual value as a fraction of the full Spitzer value.  Several authors have found temperature profiles in some galaxy clusters that are compatible with thermal conduction, and values of the suppression factor that are physically meaningful \\citep[e.g.][]{zakamska03}. In contrast, more detailed work in which the Virgo cluster was simulated using complete hydrodynamics, including radiative cooling and heating by thermal conduction has demonstrated that thermal conduction cannot prevent the occurence of a cooling catastrophe, in this particular example \\citep[][]{pope05}. Furthermore, the mass flow rates under such circumstances are not constant with radius so that mass builds up in particular areas and therefore changes the density distribution, making it much more centrally peaked than observed in real clusters.    To answer the questions regarding the necessary heating rates and mechanisms involved in galaxy clusters we construct a simple model based on two main assumptions: galaxy clusters are in approximate steady state, and the mass flow rates fulfill the constraints provided by observations at both large and small radii. To ensure this, the radial behaviour of the mass flow rates is modelled using an observationally motivated, but empirical function consistent with the observations. Using mathematical functions fitted to the temperature and density data of seven galaxy clusters we derive the required heating rates within regions which are consistent with those for which the mass flow rates were determined. We also compare the observed energy currently available in the form of plasma bubbles with the heating requirements for each cluster. From the heating rate profiles we calculate the thermal conduction suppression factor as a radial function for each cluster in order that we may determine which clusters could be heated by this process.  The plan for this paper is as follows. In Section 2 we give the details of the parameters used to fit observational temperature and density data for seven galaxy clusters. In Section 3 and 4 we describe the model used to estimate the required heating, and thermal conduction suppression factors. The heating rates we have calculated are compared with observational estimates of both AGN heating and a combination of AGN heating and thermal conduction in Sections 5 and 6. In Section 7 we summarize our main findings.   The results are given for a cosmology with $H_{0}=70{}{\\rm km s}^{-1}{\\rm Mpc}^{-1}$, $\\Omega_{M} = 0.3$ and $\\Omega_{\\Lambda} = 0.7$. ",
        "conclusions": "\\label{sect:summary}  In the heating of intracluster gas, the energy input required from AGNs is likely to depend on the strength of any other heating processes. For example, if thermal conduction is present and can reduce the integrated mass deposition rate towards the cluster centre, it is likely that in this case the AGN outbursts will have to be less powerful than in an otherwise identical cluster in which thermal conduction does not occur. There may also be additional heating processes such as stirring by galaxy motions or the gas motion resulting from past cluster-cluster mergers.  We have calculated the heating rates required to maintain a steady state for a sample of seven galaxy clusters with the assumption that the mass flow rates are independent of radius. For our model we use the spectroscopically determined integrated mass deposition rates. Here we summarise our main findings for each cluster in our sample in terms of the mass deposition rates, thermal conduction and required AGN duty cycles.  A2199: It appears that, without AGN heating, thermal conduction is sufficient for the steady-state heating requirements. In the absence of thermal conduction, the required AGN duty cycle of roughly $10^{7}$ yrs is short, but still compatible with predicted AGN lifetimes. When we consider simultaneous heating by both AGN, based on the current outburst, and thermal conduction, the suppression factor necessary to satisfy the heating requirements within the cooling radius is roughly 0.16 which is acceptable.    Virgo: Even allowing for mass deposition, thermal conduction alone is insufficient for providing the heating requirements at all radii. The required AGN duty cycle, in the absence of thermal conduction, of roughly $10^{6}$ yrs is very short compared with predicted AGN lifetimes suggesting that significantly larger AGN outflows are required at certain intervals to prevent a cooling catastrophe. This is interesting since Virgo has the smallest heating requirement in our sample. Within the cooling radius, the combined effect of the current AGN outburst and thermal conduction can achieve the required heating with a suppression factor of 0.7.   Perseus: It seems that in the absence of an AGN, thermal conduction may be insufficient compared to heating requirements at roughly roughly the central 30 kpc, but may be sufficient outside this radius. Without thermal conduction, the required AGN duty cycle of roughly $3 \\times 10^{7}$ yrs is also compatible with predicted AGN lifetimes. The suppression factor for the combined effect of AGN and thermal conduction to provide the required heating is 0.05.    Hydra: Thermal conduction alone cannot provide the necessary heating for maintaining a steady-state. The required AGN duty cycle in the absence of thermal conduction is long compared with predicted AGN lifetimes which reflects the powerful nature of the current outburst. This is the most powerful central AGN outflow in this small sample. Since the current AGN outburst in Hydra is so powerful, it more than satisfies the heating requirements within the cooling radius. Hence we derive a negative value for the suppression factor.  A2597: It appears that thermal conduction on its own cannot provide sufficient heating except between roughly 4-10 kpc. Without thermal conduction, the required AGN duty cycle is also compatible with predicted AGN lifetimes. A2597 has the largest heating requirements in this sample, followed by A478. Both require more than two orders of magnitude more heating power within the cooling radius than Virgo. It is interesting to note that even a combination of the current AGN heating rate and full Spitzer thermal conduction provide roughly only one fiftieth of the required heating rate within the cooling radius. It therefore seems as if the only means by which a cooling catastrophe can be averted in this system is by an extremely powerful AGN outburst, or a cluster-scale merger.    A1795: On its own, thermal conduction is sufficient to match the heating requirements throughout most of the cluster, although not near 5 kpc. The required AGN duty cycle, in the absence of thermal conduction, of roughly $4.3\\times 10^{7}$ yrs is also compatible with predicted AGN lifetimes. For a combination of heating by AGN and thermal conduction a suppression factor of roughly 0.2 is required to achieve the necessary heating.  A478: Thermal conduction alone should be sufficient for matching the heating requirements. In addition, the required AGN duty cycle in the absence of thermal conduction is very short compared with predicted AGN lifetimes. The central AGN outburst in this system is the least powerful in our sample.   These results suggest that heating from AGN, with duty cycles of $10^{8}$ yrs, alone is unable to maintain a steady-state, but may result in a quasi steady-state if more powerful outbursts are interspersed with the more common lower power outbursts. Alternatively, it may be that a combination of AGN heating and thermal conduction could provide the necessary heating, at least within the cooling radius, in all but one cluster of our sample: A2597. However, although thermal conduction is a possible heating mechanism of the central regions of galaxy clusters it is still not clear whether thermal conduction actually is an efficient heating mechanism in galaxy clusters. For example, from numerical simulations of the Virgo cluster by \\cite{pope05}, taking in to account thermal conduction with different suppression factors, the observed temperature profiles are most consistent with a suppression factor of between 0-0.1. This is at odds with the value of the suppression factor we have derived for most clusters in this sample. In all cases the effectiveness of thermal conduction must be fine-tuned as a function of radius in order to allow a steady-state of the cluster gas.  It is worth noting that galaxy clusters also need to be heated at radii outside the cooling radius if a time-averaged steady-state is to be maintained. Two possible mechanisms for this are sound waves excited by buoyantly rises bubbles and comparatively rare, but very powerful, AGN outbursts.    It may also be possible to predict which clusters are more likely to experience more powerful AGN outbursts in the near future. For example, if a cooling flow has formed, the fractional decrement in gas temperature between the cooling radius and the cluster centre is likely to be larger than in a cluster which has recently been heated by a large AGN outburst. Then, if the duty cycle that we have calculated based on the current heating rate, required to maintain a steady state, is much less than $10^{8}$yrs, then we may predict that a powerful outburst is iminent.  Based on the above criteria it seems that in our sample A478 and A2597 are the clusters most likely to experience a powerful AGN outburst in the near future. Virgo appears also to be a candidate for an iminent ouburst, however, the relatively small temperature difference between the cluster center and the cooling radius indicates that it may recently have experienced sufficient heating.   Our results for the required heating rates indicate that, for our initial estimate of the AGN duty cycles, if buoyant plasma bubbles are the only heating mechanism by which clusters are heated, then, of our small sample, only the AGN at the centre of Hydra is currently injecting sufficient energy to heat the gas within several multiples of the cooling radius. For the remaining six clusters, the energy available in the form of plasma bubbles is only sufficient to provide the required heating within a relatively small fraction (roughly 0.1-0.5) of the cooling radius. This significant energy deficit implies that these clusters may not be in a steady-state, or that the AGNS are active far more frequently at the centres than outside of clusters.    A possible scenario which is consistent with our results is that AGNs in general produce only relatively small or moderately powerful outbursts, such as we observe in the majority of clusters, which delay the occurence of a cooling catastrophe by a small amount each time. These small heating events may occur too infrequently to prevent the formation of a cooling flow and the temperature gradient grows with time. Eventually sufficient material is accreted by the central galaxy to initiate an extremely energetic outburst such as those seen in Cygnus-A and Hydra-A. This would also allow the heating of material at radii which are hard to reach with lower power outbursts. In this way the moderately powerful AGN outflows observed in most clusters only delay more energetic events. Such a scenario for clusters, small-scale AGN heating interpersed occasionally with more energetic events, evolution is complex and would have to be modelled numerically."
    },
    "0601/astro-ph0601605_arXiv.txt": {
        "abstract": "{Using large scale maps in \\cdo\\jtwo\\ and in the continuum at 1.2mm obtained at the IRAM-30m antenna with the Heterodyne Receiver Array (HERA) and MAMBO2, we investigated the morphology and the velocity field probed in the inner layers of the Horsehead nebula. The data reveal a non--self-gravitating ($m/m_{\\rm vir}\\approx 0.3$) filament of dust and gas (the ``neck'', $\\varnothing = 0.15-0.30\\,\\pc$) connecting the Horsehead western ridge, a Photon-Dominated Region illuminated by $\\sigma$Ori, to its parental cloud L1630. Several dense cores are embedded in the ridge and the neck. One of these cores appears particularly peaked in the 1.2\\,mm continuum map and corresponds to a feature seen in absorption on ISO maps around 7\\,$\\mu$m. Its \\cdo\\ emission drops at the continuum peak, suggestive of molecular depletion onto cold grains. The channel maps of the Horsehead exhibit an overall north-east velocity gradient whose orientation swivels east-west, showing a somewhat more complex structure than was recently reported by \\cite{pound03} using BIMA CO\\jone\\ mapping. In both the neck and the western ridge, the material is rotating around an axis extending from the PDR to L1630 (angular velocity $=1.5-4.0\\,\\kms$). Moreover, velocity gradients along the filament appear to change sign regularly (3\\,\\kmspc, period=0.30\\,pc) at the locations of embedded integrated intensity peaks. The nodes of this oscillation are at the same velocity. Similar transverse cuts across the filament show a sharp variation of the angular velocity in the area of the main dense core.  The data also suggest that differential rotation is occurring in parts of the filament. We present a new scenario for the formation and evolution of the nebula and discuss dense core formation inside the filament. ",
        "introduction": "Dense molecular cores are the basic units of isolated low-mass star formation. It is now common that molecular line mapping from dark clouds reveals clumps embedded in filamentary structures \\citep[\\eg][]{onishi98,onishi99,obayashi98,loren89a,chini97}. These filaments can be either self-gravitating or not. The clumps are mostly gravitationally bound, and in many cases, star formation is known to have already started. Therefore, molecular filaments are thought to play a crucial role in low-mass star formation.  However, little is known about the general physical properties of these filaments. For example, the density distribution is very likely not uniform, but instead varies according to a power-law. \\cite{stepnik03} have investigated this observationally in a filament of the Taurus molecular cloud, and concluded that a $r^{-2}$ profile was compatible with the observations. Steeper density profiles can however be observed, as is shown by Hily-Blant \\& Falgarone ({\\it in prep}) in a non--self-gravitating filament connected to the low-mass dense core L1512. Theoretical models with and without magnetic fields suggest power-laws with exponents ranging from $\\approx-2$ \\citep{fiege00a} to $-4$ \\citep{ostriker64, nakamura93, stodol63}.  The importance of the velocity field in the formation of clumps in filamentary clouds has already been noted several years ago by \\cite{loren89b}, who did a systematic study of the velocity pattern of well-known filaments in \\rhop. The longitudinal velocity was shown to exhibit gradients near the locations of the main clumps harboured in the filament. From the theoretical and numerical points of view, recent studies also stress the role of the velocity field. Dealing with self-similar rotating magnetized cylinders, \\cite{hennebelle03} shows that the velocity field strongly depends on the relative intensity of the toroidal component of the magnetic fields with respect to the poloidal one and to the gravitational force: the rotation can be mainly that of a rigid body or instead be differential. The importance of the velocity field has been further stressed by \\cite{tilley03} who show how it could help in distinguishing between collapsing and equilibrium cylindrical distributions and between magnetized and unmagnetized filaments. \\cite{fiege00a} performed numerical studies of the fragmentation of pressure truncated isothermal filaments threaded by helical magnetic fields and found that the toroidal velocity could change sign periodically when magnetic field dominates over gravity. More recently, \\cite{sugimoto04} numerically studied the decay of Alfv\\'en waves in filamentary clouds. They show that the propagating circularly polarized waves generate longitudinal sub-waves that steepen into shocks where the initial energy is being dissipated. While propagating, the circularly polarized Alfv\\'en waves make the filament rotate. In some cases, the filament can fragment into regularly spaced clumps.  Instabilities (gravitational, magnetic, or both) are often invoked to explain the formation of dense cores in filaments. Nonetheless, molecular observations of filamentary structures at high spatial and spectral resolution are still lacking, which would allow both the density and the velocity fields to be constrained. The scales of interest are typically of the order on 0.1\\,pc and velocity gradients of the order on 1\\,\\kmspc. In this paper, we investigate the velocity field and its link to the density distribution in the Horsehead nebula, a close-by (400\\,pc) dark protrusion emerging from its parental cloud L1630 in the Orion B molecular complex.  This condensation is illuminated by the O9.5V star $\\sigma$Ori (distance 0.5$^\\circ$ from the cloud) and presents a Photon-Dominated Region (PDR) on its western side seen perfectly edge-on \\citep{abergel03}. Until some years ago, molecular observations of this source were still scarce and limited to low spatial resolution \\citep[2\\arcmin\\ resolution, ][]{kramer96}. New insights at higher resolution (10--20\\arcsec) have recently been revealed by \\cite{abergel03} and \\cite{pound03}, who observed the Horsehead in various millimetre transitions and isotopes of CO. In particular, \\citeauthor{pound03} (hereafter \\prb) analyse the velocity field of this source, using interferometric CO\\jone\\ data (10\\arcsec$\\times$8\\arcsec~resolution), and reported a general NE-SW velocity gradient of order 5\\,\\kmspc\\ across the nebula. They also note that star formation may have already started in the cloud and discuss possible scenarios that could have driven the formation of such a structure. The use of an optically thick probe may however, have hindered them from obtaining a thorough picture of the structure in the innermost layers of this object. In the present study, we make use of the rarer isotopomer \\cdo, as well as the emission of dust at millimeter wavelengths.  The paper is organized as follows. After a brief description of the observations in Sect.~\\ref{obs}, we analyst in Sect.~\\ref{spatial} the spatial distribution of the \\cdo\\jtwo\\ emission and of support data obtained in the continuum at various wavelengths to derive representative physical parameters. The analysis of the velocity field is then explained in Sect.~\\ref{velo}, and its consequences on the genesis and evolution of the object are then discussed in Sect.~\\ref{discussion}. We finally summarize our conclusions in Sect.~\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  We have studied the morphology and velocity structure of the Horsehead nebula using new observations at high frequency and spatial resolution in the \\jtwo\\ transition of \\cdo. Our conclusions can be summarized as follows: \\begin{enumerate} \\item the emission of the likely optically thin tracer reveals a filamentary morphology similar to the one seen in dust millimetre continuum emission and visual extinction. The cloud appears as a curved ridge (the PDR) connected to L1630 by a thin (0.15\\,pc diameter) pillar hosting several condensations. One of them (Peak2) also appears in absorption in the mid-IR emission at 7\\,$\\mu$m. \\item we note that our thin \\cdo\\jtwo\\ map is much less clumpy than the possibly thick \\twCO\\jone\\ from \\citetalias{pound03}, a phenomenon that could be due to either excitation or photodissociation selective effects or the combination of both. \\item these cores are density peaks ($n_{\\rm H_2}$ in the range $2-4\\tdix{4}$\\,\\ccc) embedded in a medium that is more diffuse by a factor of 2-3 on average. Kinetic temperatures in these cores are found in the range 20-30\\,K, although the brightest of them (Peak2) shows hints of \\cdo~depletion onto grains, pointing towards a somewhat colder environment. \\item the overall shape of the filament is found to be fairly cylindrical with a diameter of about 0.15--0.30\\,pc. The total mass of the \\cdo\\ gas probed here is 20\\,\\msol, of which the largest fraction was found to be in the gaseous phase. \\item the centroid velocity map suggests that the gas in the Horsehead is wrapped around an axis coincident with the pillar. The cloud also presents various velocity gradients along the filaments described above. While the western ridge presents two strong north-south gradients inverting sign at the rise of the mane, the east-west pillar exhibits a more complex shape where gradient oscillations occur on a regular basis and seem bracketed by the location of the embedded condensations. \\item tranverse gradients are prominent almost everywhere across the pillar, but some of them depart from linear shape and might indicate differential rotation. The amplitude of the gradient is on average constant (1.5\\,\\kmspc, corresponding to a rotation period of 4 Myrs, comparable to the 5\\,Myrs lifetime estimated by \\citetalias{pound03}) at either end of the neck, but a sharp increase to 4\\,\\kmspc\\ is observed between two of the embedded peaks. Reasons for this peculiar behaviour are still unclear. \\item this complex dynamic picture raises the question of whether the origin of this protrusion could be linked to a pre-existent velocity field that progressively separated the mane and nose from the neck via centrifugal effects. The presence of several condensations unveiled inside the filament also supports the idea that dense cores are forming in the neck, and their characteristics share several similarities with modelling works on that topic. Their origin is very likely due to the combination of longitudinal (infall) and radial motions, but the actual phenomena at work may even be more complex, as the magnetic field is expected to also play a role in this process. Moreover, the origin of the initial velocity field invoked in this work remained unknown. \\end{enumerate}  Firmer conclusions on the formation scenario of such a protrusion should, however, await comparison with numerical models, and the relatively simple geometry of the Horsehead neck should be a particularly appealing application."
    },
    "0601/astro-ph0601433_arXiv.txt": {
        "abstract": "Multiply connected space sections of the universe on a scale smaller than the horizon size can leave an imprint on cosmic microwave background polarization maps, in such a way that the so-called ``circles-in-the-sky'' method can be used to detect or constrain the topology. We investigate some specific cases, namely toroidal and sixth-turn spaces, in order to show the influence of topology on CMB polarization. The correlation between matched points happens to be always positive and higher than 75\\% regardless of the angular scale and of the cosmological parameters, except for reionization. This figure is better than what occurs in temperature maps, but is achieved only in the absence of noise.  It is only slightly reduced by the filtering scheme. ",
        "introduction": "The question of a possible multiply connected topology of our cosmic space goes back to the pioneering works of Schwarzschild and Friedmann more than 75 years ago. In the past two decades, various strategies and methods have been devised to probe a non trivial topology of the space sections of the universe, using current or forthcoming data from cosmological observations~\\cite{revmeth}. Since the topological length scales are a priori not known, it is useful to survey the largest observable scale in order to increase the odds of detecting, or at least constraining, the space topology. The Cosmic Microwave Background (CMB) emitting region, the so-called last scattering surface (LSS) is situated at a redshift of $z \\simeq 1100$ and represents the deepest region that can be studied in the electromagnetic domain. Thus its usefulness for cosmology in general is tremendous~\\cite{revcmb}. Since more than one decade~\\cite{cobe}, temperature anisotropies in the CMB radiation have been mapped as a function of the direction. They result from density fluctuations in the primordial plasma, which are the seeds of evolved structures in the universe. Their study provides crucial information on the matter-energy contents of the universe, as well as some hints on the physical process which could have generated the first density perturbations at a much earlier epoch~\\cite{revcmb}.  CMB anisotropies can also be used to constrain the topology.  The main imprint of a non trivial topology on the CMB is well-known in the case when the characteristic topological length scale (called the injectivity radius) is smaller than the radius of the last scattering surface: the crossings of the LSS with its topological images give rise to pairs of matched circles of equal radii, centered at different points on the CMB sky, and exhibiting correlated patterns of temperature variations~\\cite{topocmb}. Such ``circles-in-the-sky'' searches are currently in progress using the WMAP data. They are however computationally very expensive, and present results are not completely clear. On the one hand, a massive search for matching circles with radii larger than $25^\\circ$ gave negative results~\\cite{circglenn}. But this result may well be overstated since, on the other hand, two other more specific searches gave hints for a positive detection. Roukema et al.~\\cite{circrouk} extended the search to smaller radii and claimed to have found six pairs of antipodal matched circles in a dodecahedral pattern; Aurich et al. \\cite{aurich2005} also found a marginal hint for spherical spaces, both searches being consistent with the Poincar\\'e dodecahedron space model recently proposed by some of us~\\cite{lum03} to account for the observed anomalies of the CMB angular power spectrum on large scales. The statistical significance of such results still has to be clarified. In any case, a lack of nearly matched circles does not exclude a multiply connected topology on scale less than the horizon radius: detectable topologies may produce circles of small radii which are statistically hard to detect and current analysis of CMB sky maps could have missed even antipodal matching circles because various effects may damage or even destroy the temperature matching~\\cite{flatspaces,aurich04}. Moreover, even if it might already seem severely constrained by observational data, there are still unexpected and still unexplained features on the large angular scales of the CMB temperature map~\\cite{anomisocmb} which may hint for some sort of breaking in statistical isotropy of the CMB, a feature which is not shared by many models other than multiply connected topologies, and which is present even when the injectivity radius is larger than the size of the observable universe~\\cite{flatspaces}.  Recent work by Aurich {\\it et al.}~\\cite{aurich04}, Gundermann~\\cite{gunder05} and Caillerie {\\it et al.}~\\cite{caillerie05} support the dodecahedron model. Since the above mentioned anomalies detected in the first year WMAP data~\\cite{anomisocmb} suggest the presence of some statistical anisotropy in the CMB radiation, it appears thus necessary to better constrain topology, by using tests of different nature. It is now widely understood that the polarization of the CMB, predicted long ago~\\cite{Rees68}, can provide a lot of additional informations for reconstructing the cosmological model~\\cite{HuWhite97}. The information encoded in polarization is expected to become a necessary input for precision cosmology, when it will be mapped in detail in the next release of the WMAP data~\\cite{wmaphttp} and later by the Planck satellite mission~\\cite{planckhttp}. In this article we show how the polarization can also be used to put additional constraints on space topology.  This paper is organized as follows: in section~\\ref{Seccorr}, we remind the key ingredients for detecting topological signatures in a CMB temperature map, using the circles-in-the-sky method. In section~\\ref{Secqr}, we recall the basic properties of CMB polarization that are useful for our purpose.  We then compute in section~\\ref{SecPol} the expected correlation for polarization maps, taking into account the specificities of polarization. Our main result is that the correlation in polarization maps is always larger than 75\\%, assuming an ideal case where no noise is present in the maps. In section~\\ref{SecSim} we present simulated maps for the 3-torus and the sixth-turn space in order to illustrate the validity of our method. In section~\\ref{SecBlur}, we look for various sources of blurring of the correlation when considering two effects: reionization and finite resolution. ",
        "conclusions": "High precision experiments such as WMAP now give a clear outline of the cosmological model of our universe. Despite this tremendous success of modern cosmology, some old questions such as the shape of space are still unanswered. In this paper, we have described a new test to search for the topology of the universe using polarization maps. This test is rather ambitious as it necessitates clean data on the CMB polarization map.  Given the advantages of the polarization maps which have been presented here, it appears highly likely that future missions with sufficient good signal-to-noise ratio (Planck, or some future polarization dedicated missions) will help bringing definitive conclusions about the question of cosmic topology."
    },
    "0601/cond-mat0601307_arXiv.txt": {
        "abstract": "We use Hanbury-Brown-Twiss interferometry (HBTI) to study various quantum phases of  hard core bosons (HCBs) and ideal fermions confined in a one-dimensional quasi-periodic (QP) potential. For HCBs, the QP potential induces a cascade of Mott-like band-insulator phases in the extended regime, in addition to the Mott insulator, Bose glass, and superfluid phases. At critical filling factors, the appearance of these insulating phases is heralded by a peak to dip transition in the interferogram, which reflects the fermionic aspect of HCBs.   On the other hand, ideal fermions in the extended phase display various complexities of incommensurate structures such as devil's staircases and Arnold tongues. In the localized phase, the HCB and the fermion correlations are identical except for the sign of the peaks. Finally, we demonstrate that HBTI provides an effective method to distinguish Mott and glassy phases. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601510_arXiv.txt": {
        "abstract": "{ We present a new EROS-2 measurement of the microlensing optical depth toward the Galactic Bulge. Light curves of $5.6\\times 10^{6}$ clump-giant stars distributed  over  $66\\,\\deg^2$ of the Bulge were monitored during seven Bulge seasons. 120 events were found with apparent amplifications greater than 1.6 and Einstein radius crossing times in the range $5\\,{\\rm d}<t_\\e <400\\,{\\rm d}$. This is the largest existing sample of clump-giant events and the first to include northern Galactic fields. In the Galactic latitude range $1.4\\degr<|b|<7.0\\degr$, we find $\\tau/10^{-6}=(1.62\\,\\pm 0.23)\\exp[-a(|b|-3 \\,{\\rm deg})]$ with $a=(0.43\\,\\pm0.16)\\deg^{-1}$. These results are in good agreement with our previous measurement, with recent measurements of  the MACHO and  OGLE-II groups, and  with predictions of Bulge models. ",
        "introduction": "\\label{section:introduction}   Gravitational microlensing of stars is an important tool for constraining the quantity and characteristics of faint compact objects between the stars and the observer. In microlensing events, the lensing object passes near the line of sight towards the background star, causing a transient magnification of the star's primary image as well as creating a secondary image. At Galactic scales, neither the image separation nor the image size are normally resolvable, so the only easily observable effect during a microlensing event is an apparent transient amplification of the star's flux. The characteristic timescale giving the effective duration of the amplification is proportional to the square root of the lens mass and inversely proportional to the transverse velocity. For Galactic stellar lenses, the timescales are on the order of a few weeks.  Microlensing searches were originally proposed \\citep{Pac1986} as a tool for detecting dark matter in galactic halos. Searches for the lensing of stars in the Magellanic Clouds by the MACHO \\citep{MACHOMC} and EROS-2 \\citep{EROSMC,tisserand} projects have placed constraints on the  fraction of the Milky Way halo that can be comprised of faint compact objects. Searches for objects in the halo of M31 are also being performed with candidate events reported by the VATT-Columbia \\citep{Ugle04}, WeCAPP \\citep{Riff03}, AGAPE  \\citep{PH2005}, MEGA \\citep{mega2005} and Nainital \\citep{nainital} collaborations. The AGAPE and MEGA collaborations presented efficiency calculations allowing them to constrain the content of the  M31 and Milky Way halos.   Microlensing has also been used as a tool to investigate compact objects in the visible regions of the Milky Way. Besides  the Galactic Bulge that is the subject of this paper, microlensing of stars in the spiral arms of the Galaxy has also been studied \\citep{DER01}.   Microlensing surveys toward the Galactic Bulge were first proposed \\citep{PAC91,GRI91} as a probe of ordinary stars in the Galactic Disk, though it was soon realized \\citep{KIRA94} that lensing by stars in the Bulge itself is of comparable importance. The optical depth, i.e. the probability that at a given time a  star at a distance $D_{\\rm s}$ is amplified by more than a factor 1.34 is \\begin{equation} \\tau \\;=\\; \\frac{4\\pi G D_{\\rm s}^2}{c^2}\\int_0^1 \\rho(x)x(1-x)dx \\;, \\label{optdepthint} \\end{equation} where $\\rho$ is the mass density of lenses and $x=D_{\\rm l}/D_{\\rm s}$ is the ratio between the lens and source distances. Very qualitatively, the integral for the Disk contribution to the optical depth is \\begin{equation} \\tau_{\\rm disk} \\;\\sim\\; \\frac{GM_{\\rm disk}}{c^2 h_{\\rm disk}} \\;, \\label{optdepthdisk} \\end{equation} where $M_{\\rm disk}$ and $h_{\\rm disk}$ are the total mass and scale height of the Disk lensing population. Lenses in the Bulge are near the source star  so the Bulge contribution is \\begin{equation} \\tau_{\\rm bulge} \\;\\sim\\; \\frac{GM_{\\rm bulge}}{c^2 R_{\\rm bulge}} \\;\\sim\\; \\frac{G\\rho_{\\rm bulge}R^2_{\\rm bulge}}{c^2 } \\;, \\label{optdepthbulge} \\end{equation} where $M_{\\rm bulge}$, $\\rho_{\\rm bulge}$ and $R_{\\rm bulge}$ are the mass, density and line of sight thickness of the Bulge lensing population. For Bulge stars near the direction of the Baade window ($\\ell=1\\degr ,b=-3.9\\degr$), the most recent calculations \\citep{EVA02,Bissantz,HanGould03,woodmao} give total optical depths in the range $1<\\tau/10^{-6}<2$, with about 60\\% of the rate being due to lensing by Bulge stars and the remainder by  Disk stars. The expected optical depth naturally has a strong  dependence on the Galactic latitude, $b$, falling typically from $\\tau\\sim5\\,\\times10^{-6}$ at $(l\\sim0,|b|=1\\degr) $ to $\\tau\\sim5\\,\\times10^{-7}$ at $(l\\sim0,|b|=6\\degr )$.  The first measurements of the optical depth towards the Galac\\-tic Bulge by the OGLE collaboration \\citep{UDA94b, UDA00, WOZ01}, the MACHO collaboration \\citep{ALC97, ALC00} and the MOA collaboration \\citep{ sumi2003} yielded optical depths significantly higher than these estimates, suggesting that a fundamental revision of Galactic models might be  necessary \\citep{BIN00}. However, the interpretation of the first results was difficult because of  two effects. First, stars in the direction of the Galactic Bulge are not necessarily located in the Bulge itself so foreground and background contamination must be taken into account. Second, and more importantly, photometry of the numerous faint stars in the crowded Galactic Bulge fields is complicated by ``blending'' where the reconstructed stellar flux receives contributions from more than one star. This makes the effective number of monitored stars greater than the number of cataloged stars and generates a complicated relationship between real and reconstructed amplifications.   These problems can be largely avoided by considering only microlensing of clump giant stars.  Such stars are identified by their well-defined position in the color-magnitude diagram. This position ensures that they are most likely in the Galactic Bulge.  Their large flux also makes blending problems relatively unimportant.  The first published optical depth using clump giants was based on only 13 events observed by the MACHO group \\citep{ALC97}. The optical depth, $3.9^{+1.8}_{-1.2}\\times 10^{-6}$ at $\\langle l,b\\rangle=(2.55\\degr ,-3.64\\degr )$, was still considerably higher than expectations. Since then, measurements using clump giants have tended to give results in better agreement with models. Based on 16 events, the EROS-2 collaboration \\citep{Afonso} gave an optical depth of $0.94\\pm0.29\\times 10^{-6}$ at $\\langle l,b\\rangle =(2.5\\degr ,-4.0\\degr)$. The MACHO group \\citep{machocgpop} recently presented their results using 62 events with clump giant sources.  Their optical depth is $\\tau=2.17^{+0.47}_{-0.38}\\times\\,10^{-6}$ at $\\langle l,b\\rangle =(1.5\\degr,-2.68\\degr)$, in good agreement with calculations and with their earlier preliminary results \\citep{POP00}. The OGLE-II group \\citep{sumiogle} recently used 33 events to obtain  $\\tau=2.55^{+0.57}_{-0.46}\\times 10^{-6}$ at $\\langle l,b\\rangle =(1.16\\degr,-2.75\\degr)$, again in agreement with models.  In this paper we present a measurement of the optical depth for microlensing toward the Galactic Bulge based on the totality of the EROS-2 clump-giant data. More details on this analysis can be found in \\citet{hamadache}. The 120 events used comprise the largest sample of clump-giant events studied to date. Section \\ref{section:data} describes the data collection and processing required to produce the light curves. Section \\ref{clumpsec} describes the definition of the clump-giant sample to be used to measure the optical depth. Section \\ref{section:event_selection}  lists the criteria used to isolate a sample of microlensing candidates. Section \\ref{section:efficiencies} describes the calculation of the detection efficiency, the optical depth and the timescale distribution. Section \\ref{section:effect_blending} discusses the possibility that stellar blending may affect the results. Finally, Section \\ref{conclusionsec} compares our results with those of other groups and with the predictions of Galactic models. ",
        "conclusions": "\\label{conclusionsec}   The EROS-2 optical depth measurement presented here is in remarkably good agreement with the other clump-giant measurements shown in Figure \\ref{profoptallfig}, i.e. those of the MACHO group \\citep{machocgpop} and OGLE-II group \\citep{sumiogle}. Only the original low statistics MACHO measurement \\citep{ALC97} is a bit too far  from our fitted optical depth (\\ref{profoptfit}).   \\begin{figure} \\centering \\includegraphics[width=7.8cm]{4893f20.eps} \\caption{ Optical depth as a function of Galactic latitude. The filled circles are this work while the open circle is the first EROS-2 analysis \\citep{Afonso}. The solid line shows the fit (\\ref{profoptfit}). The measurements of the MACHO group are the open squares \\citep{machocgpop} and the crossed square of the early analysis of \\citet{ALC97}. The crosses are from the OGLE-II group \\citep{sumiogle}. } \\label{profoptallfig} \\end{figure}   Figure \\ref{profoptfig} shows that our optical depth measurements are in reasonable agreement with some recent calculations based on models of the Galactic Bar and Disk. The dotted line in Figure \\ref{profoptfig}  is the prediction of the non-parametric model of \\citet{Bissantz}. The strong latitude dependence is correctly predicted as is the nearly flat longitude dependence. The three dashed lines correspond to calculations of \\citet{EVA02} based on the three models of the inner Galaxy.  Curve A is based on the  model of \\citet{Binmodel},  curve B on the model of of \\citet{Dwek}, and curve C on the model of of \\citet{Freu}. (For these three models, we show the predictions without corrections due to spiral structure; these  would give at most a 20\\% increase to the calculated optical depth.) All three models have a common total mass of $1.5\\times10^{10}M_{\\odot}$ within $2.5\\,{\\rm kpc}$ of the Galactic Center. The data in Figure \\ref{profoptfig} indicate that the ``swollen bar'' model of \\citet{Freu} is not favored by the data, unless the assumed mass is reduced by $\\sim30\\%$.   Finally, Figure \\ref{profoptfig} shows the Baade window prediction of \\citet{HanGould03}.  The calculated optical depth is within one standard deviation of the curve given by (\\ref{profoptfit}).   We make no attempt here to separate the Disk (\\ref{optdepthdisk}) and Bulge (\\ref{optdepthbulge}) contributions so our conclusion is simply that the model of \\citet{Bissantz} or models A and B of \\citet{EVA02} give a good estimate of the total optical depth (\\ref{optdepthint}). We note that these models place the bulk of the mass in normal stars and gas so the agreement with our measurements indicate that there is no need for additional mass in a non-lensing form. In models with cold dark matter, the dark component may make a significant contribution to the mass in the inner parts of the Milky Way.  For example, in the models of \\citet{KZS}, the cold dark matter contribution to the rotation curve at $1\\,{\\rm kpc}$ from the Galactic Center is between 30\\% and 50\\%. Such a contribution would, by itself, lower the optical depth by roughly the same factor. Our optical depth measurement, with its 15\\% uncertainty, therefore provides a useful constraint on such models.    Galactic models generally suppose a smooth distribution of lenses and therefore predict a smooth longitude and latitude dependence of the optical depth. The MACHO group \\citep{machocgpop} observed 9 events near $(\\ell=2.9\\degr,b=-2.9\\degr)$ (their field 104), yielding an optical depth about 2 standard deviations above the expectation based on their other fields. As seen in Figure \\ref{fig:cont_candidates}, we do not see an excess in this direction. Of their 9 events, 4 events occurred after the beginning of EROS-2. Of these 4 events, two are not in an Eros field and two are in our field 611.   Both of these events were found (in Table \\ref{eventtable1}, MACHO events 104.20515.498 and 104.20640.8423).  No other EROS events were found in this region in the three bulge seasons after the shutdown of MACHO. The fact that EROS-2 saw only two events for a solid angle and time period equivalent to that in which MACHO found 9 events suggests that the MACHO excess is a statistical fluctuation.   While our optical depth measurement is in good agreement with other measurements and with Galactic models, agreement on  the $t_\\e$ distribution is less satisfactory. Our value of $\\langle t_\\e \\rangle$ is in excellent agreement with the value,  $(28.1\\,\\pm4.3)\\,{\\rm d}$, found by  OGLE-II \\citep{sumiogle}. The EROS-2 and OGLE-II values are, however, significantly higher than that calculated using the 62 MACHO events (Table 3 of \\citep{machocgpop}): $\\langle t_\\e\\rangle=(21.6\\,\\pm 3) \\,{\\rm d}$. The discrepancy lies entirely in the large number of MACHO events with $t_\\e < 10\\,{\\rm d}$. The number of MACHO events with ($t_\\e<5$, $5<t_\\e<10$, $t_\\e>10\\day$) is $(6,14,42)$ whereas the corresponding numbers for EROS-2 are $(0,10,110)$.  Part of the differences in event numbers is due to the fact the the MACHO efficiency is relatively greater than the EROS-2 efficiency for $t_\\e<10\\,{\\day}$. Using the 14 MACHO events with $5\\,\\day<t_\\e<10\\,\\day$ and the relative efficiencies of the two experiments, we estimate that EROS-2 should see $24\\pm7$ events in this range compared to the 10 that are seen.  The very low EROS-2 efficiency for  $t_\\e<5\\,{\\day}$ makes it unsurprising that we found no events in this range.   Whatever the source of the discrepancy, these short events have relatively little effect on the measured optical depths. The events with $t_\\e<10\\,{\\rm d}$ constitute only about 12\\% of the MACHO optical depth and only 4\\% of the EROS-2 optical depth.   The Galactic models are  capable of reproducing the $t_\\e$ distribution by adopting an appropriate stellar mass function. The model of \\citet{woodmao} is in good agreement with our distribution and with that of OGLE-II, both with $\\langle t_\\e \\rangle\\sim28\\,\\day$. The model of \\citet{BI2004} is in good agreement with the distribution reported by MACHO with $\\langle t_\\e \\rangle\\sim21\\,\\day$. According to \\citet{woodmao}, the different $\\langle t_\\e\\rangle$ predictions of the two models is mostly due to different adopted mass functions leading to a mean lensing mass of $0.35M_\\odot$ for \\citet{woodmao} and $0.11M_\\odot$ for \\citet{BI2004}.  Kinematical differences between the two models may also have a non-negligible effect.  The present comparison with models is still primarily limited by the low number of observed events.  Since the EROS-2 survey uses most of the fields in the Bulge region that are easily observed in optical bands, a large increase in the number of monitored clump giants would require infrared observations of highly obscured regions near the Galactic Center \\citep{gouldkband}. Of course the number of events can also be increased by using dim stars but, in this case, reliable results would require a good understanding of blending.   \\bigskip"
    },
    "0601/astro-ph0601660_arXiv.txt": {
        "abstract": "We present images of the jets in the nearby radio galaxy NGC\\,315 made with the VLA at five frequencies between 1.365 and 5\\,GHz with resolutions between 1.5 and 45\\,arcsec FWHM. Within 15\\,arcsec of the nucleus, the spectral index of the jets is $\\alpha = 0.61$. Further from the nucleus, the spectrum is flatter, with significant transverse structure. Between 15 and 70\\,arcsec from the nucleus, the spectral index varies from $\\approx$0.55 on-axis to $\\approx$0.44 at the edge. This spectral structure suggests a change of dominant particle acceleration mechanism with distance from the nucleus and the transverse gradient may be associated with shear in the jet velocity field. Further from the nucleus, the spectral index has a constant value of 0.47. We derive the distribution of Faraday rotation over the inner $\\pm$400\\,arcsec of the radio source and show that it has three components: a constant term, a linear gradient (both probably due to our Galaxy) and residual fluctuations at the level of 1 -- 2~rad\\,m$^{-2}$. These residual fluctuations are smaller in the brighter (approaching) jet, consistent with the idea that they are produced by magnetic fields in a halo of hot plasma that surrounds the radio source. We model this halo, deriving a core radius of $\\approx$225\\,arcsec and constraining its central density and magnetic-field strength. We also image the apparent magnetic-field structure over the first $\\pm$200\\,arcsec from the nucleus. ",
        "introduction": "\\label{intro}   The giant FR\\,I \\citep{FR74} radio source NGC\\,315 was first imaged by \\citet{Brid76}, who showed that it has an angular size of nearly 1$^\\circ$.  The extended radio structure is described in more detail by \\citet{Brid79}, \\citet{Fom80}, \\citet{Willis81}, \\citet{Jaegers}, \\citet{Venturi93}, \\citet{Mack97} and \\citet{Mack98}. The main jet has also been imaged extensively on parsec scales \\citep{Linfield81,Venturi93,Cotton99,Xu00}. X-ray emission from the first 10\\,arcsec of the main jet was detected by \\citet{WBH}, but no optical emission from this region has yet been reported.   The source is associated with a giant elliptical galaxy at a redshift of 0.01648 \\citep{Trager}, giving a scale of 0.335\\,kpc/arcsec for our adopted cosmology (Hubble constant $H_0$ = 70\\,$\\rm{km\\,s^{-1}\\,Mpc^{-1}}$, $\\Omega_\\Lambda = 0.7$ and $\\Omega_M = 0.3$).  NGC\\,315 is a member of a group or poor cluster of galaxies \\citep{Nolthenius,Miller02} located in one of the filaments of the Pisces-Perseus supercluster \\citep{Ensslin01,Huchra}. HST images \\citep{Verdoes99} show a 2.5-arcsec diameter dust lane and a nuclear point source. The dust lane is associated with a disk of ionized gas which is probably in ordered rotation \\citep{Noel-Storr}. CO emission, also with a line profile indicating rotation, was detected by \\citet{Leon}. The inferred mass of molecular hydrogen is $(3.0 \\pm 0.3) \\times 10^8$\\,M$_\\odot$ and the cold gas is likely to be cospatial with the dust. HI absorption against the nucleus was detected by \\citet{vanG89}. There is evidence for a weak, polarized broad H$\\alpha$ line in the nuclear spectrum \\citep{Ho97,Barth99,Noel-Storr}. Hot gas associated with the galaxy has been imaged using {\\sl ROSAT} and {\\sl Chandra} \\citep{WB,WBH}.  Within $\\approx$\\,90 arcsec of the nucleus, the jets in NGC\\,315 are initially narrow, then expand rapidly (``flare'') and re-collimate \\citep{Brid82,CLBC}. We have modelled the inner $\\pm$70\\,arcsec of this {\\em flaring region} as a two-sided, symmetrical, relativistic flow, fitting to deep, high-resolution VLA observations at 5\\,GHz in order to derive the three-dimensional distributions of velocity, proper emissivity and magnetic-field structure \\citep{CLBC}. Our main conclusions are as follows. \\begin{enumerate} \\item The jets are inclined by $38^\\circ \\pm 2^\\circ$ to the line of sight. \\item Where they first brighten, their on-axis velocity is $\\beta = v/c \\approx 0.9$. They decelerate to $\\beta \\approx 0.4$ between 8 and 18\\,kpc from the nucleus (15 -- 33\\,arcsec in projection) and the velocity thereafter remains constant. \\item The ratio of the speed at the edge of the jet to its value on-axis ranges from $\\approx$0.8 close to the nucleus to $\\approx$0.6 further out. \\item The longitudinal profile of proper emissivity is split into three power-law regions separated by shorter transition zones and the emission is intrinsically centre-brightened. \\item  To a first approximation, the magnetic field evolves from a mixture of longitudinal and toroidal components to predominantly toroidal by 26\\,kpc (48\\,arcsec in projection). \\item Simple adiabatic models fail to fit the emissivity variations. \\end{enumerate}  In the present paper, we investigate the energy spectrum of the relativistic particles in the jets of NGC\\,315 in the context of the models developed by \\citet{CLBC}. We use VLA observations at frequencies between 1.365 and 5\\,GHz\\footnote{The 5-GHz observations are those discussed by \\citet{CLBC}} to derive the spectrum of the jets at resolutions of 5.5 and 1.5\\,arcsec and relate the observed spectral gradients to velocity, emissivity and field structure. A separate paper (Worrall et al., in preparation) will describe the radio structure of the main jet at high resolution and its relation to new {\\sl Chandra} images.  We also determine the variations of Faraday rotation over the jets and test the hypothesis that these result from magnetic-field irregularities in hot, X-ray emitting plasma associated with the surrounding group of galaxies. Finally, we determine the apparent magnetic-field structure of the jets on scales larger than those covered by \\citet{CLBC}.  In Section~\\ref{obs}, we describe the observations and their reduction. The total-intensity images are presented in Section~\\ref{Images} and we use them to derive distributions of spectral index in Section~\\ref{Spectra}. We then discuss the distributions of Faraday rotation (Section~\\ref{Faraday}) and apparent magnetic-field structure (Section~\\ref{Field}) derived from observations of linear polarization.  Section~\\ref{Summary} summarizes our main results. ",
        "conclusions": "\\label{Summary}  We have imaged the jets in the nearby FR\\,I radio galaxy NGC\\,315 with the VLA at five frequencies in the range 1.365 -- 5\\,GHz and at resolutions ranging from 45 -- 1.5\\,arcsec FWHM. Our total intensity observations reveal new details of the structure, particularly around the sharp bend in the main jet.  The flaring regions of both jets, where they initially expand rapidly and then recollimate, show a complex and previously unknown spectral structure. Within 15\\,arcsec of the nucleus, the spectral index has a uniform value of $\\alpha =$ 0.61 in both jets. This region is associated with strong X-ray emission in the main jet, high radio emissivity, complex filamentary structure, and fast flow with $\\beta \\approx 0.9$ \\citep{CLBC,WBH}.  Between 15 and 70\\,arcsec, the spectrum is steeper on-axis than at the edges of the jet. We have developed a novel deprojection technique which allows us to isolate two spectral components. The first (on-axis) forms a continuation of the jet base, its spectral index flattening gradually from 0.61 to 0.55. The second (at the edge of the jet) has $\\alpha \\approx$ 0.44 and is associated with a region where strong shear is inferred \\citep{CLBC}. We speculate that two different acceleration mechanisms are involved, one associated with fast flow, dominant close to the nucleus and capable of accelerating electrons to the very high energies required to produce X-ray emission ($\\gamma \\sim 10^8$), the other being driven by shear and generating the flatter spectral indices seen at the edges of the jet.  Both mechanisms must efficiently generate the electrons with $\\gamma \\sim 10^4$ which radiate at cm wavelengths. At distances $\\ga$70\\,arcsec, the spectral index is consistent with $\\alpha \\approx 0.47$ everywhere.  We have imaged the variations of Faraday rotation over the jets. All of the rotation is resolved and must originate mostly in foreground material. There is no detectable depolarization. The largest contributions -- a constant term and a linear gradient -- are probably Galactic in origin. We have also detected residual fluctuations of $\\approx$1 -- 2 rad\\,m$^{-2}$ rms on scales $\\sim$ 5 -- 100\\,arcsec.  The amplitude of fluctuations on scales $\\ga$30\\,arcsec is larger by a factor $\\approx$2 for the counter-jet, consistent with an origin in magnetoionic material around the source, but not in the known X-ray-emitting halo, whose core radius is too small. We model the Faraday-rotating medium as a spherical halo with a core radius $\\approx$225\\,arcsec and derive an approximate value for the product $(n_0/{\\rm m}^{-3})^2 (B_0/{\\rm nT})^2 (l/{\\rm kpc}) \\approx 700$, where $n_0$ is the central density, $B_0$ the central magnetic field and $l$ is the magnetic-field correlation length. Our analysis is therefore consistent with models of Faraday rotation proposed for rich clusters (e.g.\\ \\citealt{CT}), but requires much lower densities and field strengths.  We predict that a tenuous, group-scale halo should be detectable in sensitive X-ray observations; measurement of its density will allow us to estimate the magnetic-field strength.  We have derived the apparent magnetic field direction (corrected for Faraday rotation) and degree of polarization at distances between 70 and 160\\,arcsec from the nucleus. The structure is qualitatively similar to that seen in the flaring region, with transverse field on-axis and longitudinal field at the edges of both jets, but the degree of polarization on-axis is larger.  The difference in polarization structure between the main and counter-jets observed in the flaring region by \\cite{CLBC} persists at larger distances. This can be explained fully as an effect of differential aberration on radiation from intrinsically identical jets, as long as their velocities remain significantly relativistic on the relevant scales.  The asymmetry in RM fluctuation amplitude is consistent with the jet orientation required by this analysis and the presence of a tenuous, magnetized group halo.  The large angular size of the flaring region in NGC\\,315 and our use of deep observations at several frequencies has allowed us to image spectral variations at a level of detail not yet achieved in any other jet. Taken together with X-ray imaging and modelling of the jet velocity field, this has given important insights into the particle acceleration mechanisms. It will be interesting to see whether our results apply to other FR\\,I jets and to study the spectral variations in NGC\\,315 over a wider frequency range."
    },
    "0601/astro-ph0601383_arXiv.txt": {
        "abstract": "The formation of dust with temperature-dependent non-grey opacity is considered in a series of self-consistent model atmospheres at different phases of an O-rich Mira variable of mass 1.2~$M_\\odot$. Photometric and interferometric properties of these models are predicted under different physical assumptions regarding the dust formation. The iron content of the initial silicate that forms and the availability of grain nuclei are found to be critical parameters that affect the observable properties. In particular, parameters were found where dust would form at 2-3 times the average continuum photospheric radius. This work provides a consistent physical explanation for the larger apparent size of Mira variables at wavelengths shorter than 1\\,$\\mu$m than that predicted by fundamental-mode pulsation models. ",
        "introduction": "Dust shells around Mira variables often dominate the mid- and far-infrared regions of their spectra and strongly influence the observed brightness distributions of their stellar and circumstellar material. In contrast to C-type Miras for which substantial theoretical progress has been made in recent years (e.g. \\citealt{Helling00,Schirrmacher03,Hofner03,GautschyLoidl04}), for oxygen-rich M-type Miras the physical conditions and the geometric region of formation of dust grains is poorly understood.  This situation is first of all due to the much more complex physics of formation of typical dust grains, which may contain metal oxides such as corundum (Al$_2$O$_3$), silicates or even solid iron \\citep{Gail03}. These heterogenous grains form around nuclei that are in general composed of different compounds again, such as TiO$_2$ \\citep{Jeong03}. Furthermore, in order to understand the occurrence of dust in an M-type Mira, the formation of grains has to be studied within the physical environment of the time-dependent dynamic atmosphere of the pulsating star. Interferometric observations indicate that the inner radii of dust shells may be as small as 2 to 4 stellar radii where the stellar radius refers to the position of the continuum-forming layers (e.g. \\citealt{Danchi94,Danchi95,Lopez97,Monnier04,Ireland04a,Ireland05}). This is is well within the upper regions of the stellar atmosphere in which typical molecular features of the stellar spectrum are formed (cf. \\citealt{Hofmann98,Tej03b,Tej03,Ohnaka04a,Perrin04b,Ohnaka05})  So far existing dynamic model atmospheres of M-type Mira stars either are dust-free (e.g. \\citealt{Hofmann98,Woitke99,Hofner03}), or they use grey, composition independent dust opacities (e.g. \\citealt{Winters00,Jeong03}). \\citet{Bedding01} studied schematically the conditions of dust formations in the models of \\citet{Hofmann98} and concluded that corundum and silicate grains are likely to occur, depending on the phase of pulsation, in the upper layers of the star's atmosphere. However, only a small number of observable features of a small number of models were considered, with a very simple opacity treatment.  In this paper we investigate the formation of dust in a series of dynamic models of a Mira variable, and discuss which observable effects would result from different physical assumptions regarding the dust formation. In particular, we focus on effects of non-grey opacities of silicate dust with a variable Fe content and the effect of grain size on observable properties. As hetrogenous O-rich dust formation is complex, our approximations are designed to be most valid for only the innermost dust that has a significant influence on observable properties. We aim to parametrize uncertainties in this first significant dust formation with the smallest number of physical free parameters. ",
        "conclusions": "Condensation of corundum and silicate dust in a model of a 1.2~$M_\\odot$ self-excited Mira variable has been examined. Four dust types were modelled, representing approximated physical dust formation parameters. It was found that a dust type between A and B from Table~\\ref{tblDustTypes} fits existing observations well \\footnote{Predictions of existing and additional models are available upon request to anyone interested for specific observational programs.}. In particular, the effect of scattering by this type of dust explains the systematically large observed apparent diameters at wavelengths shorter than 1\\,$\\mu$m.  This dust type has approximately $10^{-13}$ grain nuclei per H atom available at the onset of silicate condensation, and requires $\\alpha$, the ratio of the exchange coefficient for Mg and Fe ion to the sticking coefficient of SiO, to be at least the order of unity. A model where only corundum forms was found to be a poor fit to observations of $o$~Cet and R~Leo. The use of composition-dependent opacities enabled the survival of the initial Mg-rich silicate condensate at radii of 2-3 times the model parent star radius."
    },
    "0601/astro-ph0601106_arXiv.txt": {
        "abstract": "We have studied the X-ray point source population of the 30~Doradus star-forming complex in the Large Magellanic Cloud using high-spatial-resolution X-ray images and spatially-resolved spectra obtained with the Advanced CCD Imaging Spectrometer (ACIS) aboard the {\\em Chandra X-ray Observatory}.  Here we describe the X-ray sources in a $17\\arcmin \\times 17\\arcmin$ field centered on R136, the massive star cluster at the center of the main 30~Dor nebula.  We detect 20 of the 32 Wolf-Rayet stars in the ACIS field.  R136 is resolved at the subarcsecond level into almost 100 X-ray sources, including many typical O3--O5 stars as well as a few bright X-ray sources previously reported.  Over two orders of magnitude of scatter in $L_X$ is seen among R136 O stars, suggesting that X-ray emission in the most massive stars depends critically on the details of wind properties and binarity of each system, rather than reflecting the widely-reported characteristic value $L_X/L_{bol} \\simeq 10^{-7}$.  Such a canonical ratio may exist for single massive stars in R136, but our data are too shallow to confirm this relationship. Through this and future X-ray studies of 30~Doradus, the complete life cycle of a massive stellar cluster can be revealed. ",
        "introduction": "}  Stars of virtually all masses and stages emit X-rays in their youth, although the mechanisms for X-ray emission vary with stellar mass. T-Tauri and protostars have magnetic flares with $L_X$ up to $10^{31-32}$~ergs~s$^{-1}$ \\citep{Feigelson05}.  Individual OB stars emit X-rays at levels $L_X \\sim 10^{-7}$~L$_{bol} \\sim 10^{32-34}$~ergs~s$^{-1}$, probably arising from shocks within their unstable $>1000$~km~s$^{-1}$ winds \\citep[e.g.][]{Berghoefer97}. Colliding winds of binary O and Wolf-Rayet (WR) stars can produce shocks with $L_X$ up to $10^{35}$~ergs~s$^{-1}$ variable on timescales of days \\citep[e.g.][]{Pollock95}.  For OB stars excavating an H{\\sc II} region within their nascent molecular cloud, diffuse X-rays may be generated as fast winds shock the surrounding media \\citep{Weaver77}; we have recently discovered such parsec-scale diffuse emission with {\\it Chandra} observations of the Galactic high-mass star-forming regions M17 and the Rosette Nebula \\citep{Townsley03}.  30~Doradus (30~Dor) is the largest H{\\sc II} region in the Local Group, hosting several young, massive stellar clusters, several well-known supernova remnants (SNRs), and a vast network of superbubbles created by current and past generations of massive stars and their supernovae. At the center of 30~Dor is the massive compact stellar cluster R136 \\citep[][called ``RMC~136'' in SIMBAD]{Feast60}, the richest resolved OB association with dozens of 1--2~Myr-old $> 50$~M$_\\odot$ O and WR stars \\citep{Massey98}, including several examples of the recently-defined earliest spectral type, O2 \\citep{Walborn02}. Including stars down to $0.1$~M$_{\\odot}$, R136 has a mass of $\\sim 6 \\times 10^4$~M$_{\\odot}$ \\citep{Brandl05}; there are $>$3500 stars within its 10-pc diameter \\citep{Massey98}.  30~Dor contains other stellar clusters 1--10 million years old, the $\\sim 20$~Myr old stellar cluster Hodge~301 \\citep{Grebel00}, and a new generation of deeply embedded high-mass stars just now forming \\citep[][and references therein]{Walborn02b}, making it a prime example of sequential and perhaps triggered star formation.  Recent radio observations have revealed water \\citep{Loon01} and OH \\citep{Brogan04} masers in 30~Dor, also signifying ongoing star formation.  The presence of these young objects as well as a large population of widely-distributed WR stars \\citep{Moffat87} and evolved supergiants $\\sim 25$~Myr old \\citep{Walborn97} shows that 30~Dor is the product of multiple epochs of star formation.  The new generation of embedded stars currently forming may be the result of triggered collapse from the effects of R136 \\citep{Brandner01}.  We observed 30~Dor with {\\em Chandra's} Advanced CCD Imaging Spectrometer (ACIS) camera for $\\sim 21$~ks in 1999 September; see Figure~\\ref{fig:bin4data}.  R136 was positioned at the aimpoint of the $17\\arcmin \\times 17\\arcmin$ ACIS Imaging Array (ACIS-I).  The reader is referred to the accompanying paper \\citep[][henceforth Paper I]{Townsley05} for a review of previous X-ray studies of the 30~Dor complex, a description of our {\\em Chandra} observations and data analysis, and a detailed treatment of the superbubbles, SNRs, and other diffuse structures in the field.  Here we consider the point sources located in the stellar clusters and distributed across the field.  Our methods for detecting and extracting sources are described in \\S\\ref{sec:srcdetect}.  The X-ray properties of these sources and counterparts identified from other studies are tabulated in \\S\\ref{sec:ptsrcs}.  All luminosities are calculated assuming a distance of 50 kpc.  Many of the 180 ACIS-I point sources identified in this observation are early-type stars or stellar systems associated with the very dense, massive stellar cluster R136 (seen as a large concentration of sources at the center of Figure~\\ref{fig:bin4data}).  Our efforts to resolve and identify sources in R136 are detailed in \\S\\ref{sec:R136}.  Several point sources are spatially associated with the N157B SNR; some may be clumps in the diffuse emission, but the cospatial OB association LH~99 may account for true point sources in this region.  Several dozen faint sources are scattered across the field, notably the well-known WR stars R130, R134, R139, R140a, R140b, R144, and R145.  Most X-ray sources without visual or infrared (IR) counterparts are likely background active galactic nuclei (AGN), but some appear to have unusually hard spectra and may be low-luminosity X-ray binaries associated with the 30~Dor complex.  These distributed sources are described in \\S\\ref{sec:otherpoints}. ",
        "conclusions": "We have analyzed early {\\em Chandra}/ACIS observations of 30~Doradus and catalog here a wide variety of X-ray-emitting point sources: individual high-mass stars, colliding-wind binaries, possibly evolved X-ray binaries, foreground stars, and background AGN.  Using publicly-available tools, we describe both the spatial and spectral characteristics of these X-ray sources.  In this $\\sim 21$~ks observation of 30~Dor, 180 X-ray point sources are found in the $17\\arcmin \\times 17\\arcmin$ {\\em Chandra}/ACIS-I field. Most (85\\%) of these sources are associated with high-mass stellar systems.  While many of the remaining sources are likely background AGN, a handful of white dwarf X-ray binaries may be present and one known foreground star is seen.  Nearly 100 of the sources lie in the R136 massive star cluster; several of these may form a small subcluster around the WR star Mk~34.  A few sources are collected in the LH~99 cluster around the supernova remnant N157B.  Individual O stars emit X-rays at levels typical of their Galactic counterparts; we detect a number of early-type (O3--O7) stars by using image reconstruction to reduce the confusion in the R136 core.  The X-ray luminosity of R136, however, is dominated by a few sources known to be WR-WR or WR-O star binaries.  Their spectra show higher plasma temperatures than are typical for single early-type stars; as first noted by \\citet{PPL02}, these are most likely colliding-wind binaries. None of them show variable lightcurves in this short observation.  The individual systems display a wide variety of X-ray luminosities, ranging from Mk~34 with absorption-corrected $L_X = 2.2 \\times 10^{35}$~ergs~s$^{-1}$ to many with $L_X < 4 \\times 10^{32}$~ergs~s$^{-1}$ (assuming a typical stellar spectral model).  Thus we find no evidence for a characteristic $L_X/L_{bol}$ ratio for the massive members of R136, such as the widely reported relation $L_X/L_{bol} \\simeq 10^{-7}$ \\citep[e.g.][]{Harnden79, Pallavicini81, Berghoefer97} that holds for later-type stars.  Rather, the X-ray emission seems to depend critically on the details of the wind properties and binarity for high-luminosity WR and O stars.  For example, we infer that Mk~34, R136c, and R136a3 are likely close binaries due to their high X-ray luminosities and/or spectral hardness, while R134 and Mk~33Sa probably are not.  Although our {\\em Chandra} exposure is too short to penetrate deeply into the stellar X-ray luminosity function of the region, this is still the richest field of very-high-mass stellar sources ever characterized in the X-ray band.  The high stellar density in R136 makes this field a showcase for {\\em Chandra's} excellent spatial resolution.  This study marks the beginning of our attempts to understand the high-energy processes at work in the massive stars of 30~Dor; partnered with Paper I, we are beginning to see how the life cycle of high-mass stars shapes the overall view of giant H{\\sc II} regions and their surroundings through their powerful winds and supernovae.  These results from early {\\em Chandra} observations lay the groundwork for more sensitive {\\em Chandra} studies planned for the near future."
    },
    "0601/astro-ph0601276_arXiv.txt": {
        "abstract": "\\noindent We derive an extended version of the well-known Lyth Bound on the total variation of the inflaton field, incorporating higher order corrections in slow roll. We connect the field variation $\\Delta\\phi$ to both the spectral index of scalar perturbations and the amplitude of tensor modes.  We then investigate the implications of this bound for ``small field'' potentials, where the field rolls off a local maximum of the potential. The total field variation during inflation is {\\em generically} of order $m_{\\rm Pl}$, even for potentials with a suppressed tensor/scalar ratio. Much of the total field excursion arises in the last e-fold of inflation and in single field models this problem can only be avoided  via fine-tuning or the imposition of a symmetry.  Finally,  we discuss the implications of this result for inflationary model building in string theory and supergravity. ",
        "introduction": "Inflation \\cite{Guth:1980zm} has proven to be a fantastically successful paradigm for explaining why the universe is so big, so old, and so flat, and provides a mechanism for generating the primordial cosmological perturbations. However, a fundamental description of the physics responsible for inflation remains elusive: no convincing model for inflation yet exists. We can, however, discern the broad characteristics of the  inflationary epoch from observational data. The absence of a detectable gravitational wave contribution to the CMB (Cosmic Microwave Background) anisotropies tells us that the energy density of the universe during inflation was something less than $10^{16}\\ {\\rm GeV}$. The fact that the primordial perturbation spectrum is close to scale invariant tells us that the potential function of the field or fields responsible for inflation (the {\\it inflaton}) must be very flat. Finally, the highly suppressed amplitude of the perturbation spectrum ($\\delta \\sim 10^{-5}$) requires the introduction of a small parameter to the inflationary Lagrangian, either as a small self-coupling or as a hierarchy of mass scales.  Probably the best guess as  to the fundamental origin of the inflationary potential is that the inflaton (if there is only one) is an effective degree of freedom in the low-energy  limit of quantum gravity \\cite{Lyth:1998xn}. The last few years have seen substantial progress toward realizing this hypothesis within string theory \\cite{Dvali:1998pa,Kachru:2003sx,Silverstein:2003hf,Burgess:2004kv,Alishahiha:2004eh,Burgess:2005sb,Dimopoulos:2005ac,Aazami:2005jf,Easther:2005zr}, but no dominant and compelling stringy model of inflation has emerged, and string phenomenology is far from fully understood.  In this context, effective field theory is a  powerful tool.   Since the inflaton potential is tightly constrained by data, if the inflaton field takes on a vacuum expectation value (vev) at which higher order operators contribute significantly to the potential, it is very unlikely that the delicate balance between the height and slope of the potential will survive intact. Therefore effective field theory arguments favor a vev for the inflaton which is small in Planck units: $\\Delta \\phi \\ll m_{\\rm Pl}$. Lyth \\cite{Lyth:1996im} derived a lower bound on the variation in the inflaton field during inflation in terms of the ratio $r$ between tensor and scalar perturbations generated during inflation, known as the {\\em Lyth Bound}. Combined with a theoretical prejudice based on effective field theory, the Lyth Bound predicts a  strong suppression of primordial gravitational waves relative to the observed amplitude of the scalar perturbations.   The significance of the Lyth Bound has been the subject of much debate. In particular, Linde has argued that only the {\\em energy density} during inflation must be sub-Planckian and that field values greater than $m_{\\rm Pl}$ are physically consistent \\cite{Linde:2004kg}, but this argument does not apply to generic supergravity or superstring inspired models \\cite{Easther:2005zr}. Lyth's argument is restricted to single field models, and multifield models can evade it in two ways. Firstly, one might imagine that several separate inflationary epochs are concatenated together.\\footnote{See \\cite{Easther:2004ir,Burgess:2005sb} for specific proposals along these lines, and \\cite{Aazami:2005jf} for a more general argument about the implementation of multifield models in the string landscape.}   Secondly, the large number of scalar fields introduced by the string landscape makes it natural to look for implementations of {\\em assisted inflation\\/} \\cite{Liddle:1998jc} where a large number of similar (if not identical) scalar fields act cooperatively to drive inflation.  The resulting cosmology mimics that produced by a single field with a large vev, but the individual fields never take on trans-Planckian expectation values.  A concrete example of this approach is provided by  ``N-flation'' \\cite{Dimopoulos:2005ac,Easther:2005zr}.  Questions surrounding the robustness and generality of the Lyth Bound will only grow in importance as measurements of the perturbation spectrum put ever tighter constraints on the amplitude of primordial  tensor fluctuations \\cite{Kinney:2003gi,Alabidi:2005qi}. In this paper, we do not attempt to resolve the debate about the self-consistency of effective field theory in inflation, but simply pose the question: If we suppose that future observations show a strongly suppressed amplitude of primordial tensor fluctuations, what will that tell us about the inflaton potential?  To investigate this question, we derive an extended version of the Lyth Bound accurate to higher order in slow roll. This allows us to connect the field variation $\\Delta\\phi$ to  the spectral index of scalar perturbations, as well as to the amplitude of tensor modes. We then investigate the implications of this bound for ``small field'' potentials, where the field rolls off a local maximum of the potential. We find that, in the absence of fine-tuning, the field variation during the inflationary period is {\\em generically} of order $m_{\\rm Pl}$, even for models with a very small tensor/scalar ratio.  Specifically,  near the end of inflation, the field variation is rapid enough to ensure that  $\\Delta\\phi \\sim m_{\\rm Pl}$ during the last e-fold of expansion. This is true even when the field variation is small during the epoch in which the observable perturbations are generated. We show that this ``last e-fold problem'' can only be solved by suppressing the inflaton mass term -- either by fine-tuning or the imposition of a symmetry. The remainder of the paper is structured as follows: Section \\ref{sec:extendedlythbound} derives the higher-order extension of the Lyth Bound. Section \\ref{sec:smallfieldmodels} discusses the issue of field variation in generic small-field models and the ``last e-fold'' problem. Section \\ref{sec:conclusions} contains a summary and conclusions. ",
        "conclusions": "\\label{sec:conclusions}  We consider the consequences of requiring  the  change in the inflaton field during its evolution  to be  sub-Planckian, or, $\\Delta\\phi < m_{\\rm Pl}$. We derive a higher order version of Lyth's relationship between the field variation and the tensor/scalar ratio, $r$. This is valid to second order in slow roll, and is given by Eq.~(\\ref{extended}). In addition to $r$, the field variation also depends on the scalar spectral index $n$. Naively interpreted, this result appears to indicate that a spectrum close to scale invariance ($n \\simeq 1$) makes it easier to ensure that $\\Delta\\phi$ is small, but we find that this is a model-dependent statement.   Arguments based on effective field theory suggest that a self-consistent inflation model must have $\\Delta\\phi < m_{\\rm Pl}$ during the interval in which cosmological perturbations are laid down. To satisfy this constraint with a single field model, the tensor/scalar ratio is strongly suppressed: $r \\ll 1$.  This argument amounts to putting a strong theoretical prior on the inflationary model space -- namely that we are working with a model derived from supergravity or string theory. In this case, an effective field theory description of the inflaton potential is very unlikely to be reliable when the inflaton acquires a trans-Planckian vev, as the delicate balance between the height and slope of the potential required for successful inflation will almost certainly be disrupted.  In this analysis, we do not take a position on whether it is reasonable to restrict the inflationary model space in this way. Rather, we take the view that observations will eventually either determine the value of $r$ or place a tight upper bound on its value, and we investigate the theoretical implications of these different possibilities. Lyth's  original discussion  considered only the field variation during the interval during which perturbations relevant to the lowest multipole moments CMB were being laid down, relevant to the detectable signal from gravitational waves. We extend this result in two ways. Firstly, we include next-to-leading order terms in the slow-roll expansion, which yields an expression for $\\Delta \\phi$ which depends on both $r$, and the scalar spectral index $n$.  Secondly,  we consider the field variation over the final 60 e-foldings of inflation which gives a tighter constraint on the overall model space. In particular, we found that the field point generically evolves by an amount close to $m_{\\rm Pl}$ during the final moments of inflation, leading to what we have dubbed the ``last e-fold problem'': almost all {\\em minimally coupled, single field\\/} inflationary models will have $\\Delta  \\phi \\apprge m_{\\rm Pl}$.  The simplest way to build an inflationary model with $\\Delta\\phi \\ll m_{\\rm Pl}$ is hybrid inflation \\cite{Linde:1993cn}. In a hybrid scenario, the field evolves toward a minimum with nonzero vacuum energy. On such a potential, the inflationary solution asymptotes to $\\dot\\phi = 0$ as $\\phi$ approaches the minimum of the potential, after an {\\em infinite} amount of inflation \\cite{Kinney:2005vj}.  Consequently, in these models inflation ends because of an instability in an orthogonal direction, induced by an auxiliary field. Clearly, this scenario yields an arbitrarily large amount of inflation with an arbitrarily small field variation. Such models are therefore well-behaved from the standpoint of effective field theory.  It might be argued that hybrid inflation is fine-tuned,  as it relies on having an auxiliary field with an instability exponentially close to the origin. In practice, hybrid models exploit the presence of a hierarchy of mass-scales in nature, and so do not involve any {\\em extra\\/} tunings.  In the case of ``inverted'' potentials characteristic of spontaneous symmetry breaking, inflation ends dynamically when the slow roll parameter $\\epsilon$ passes through unity, and no extra parameters are required to fix the endpoint of inflation. Near the origin, the potential is dominated by the lowest order term in its Taylor series, so we consider ``small-field'' models where $V(\\phi) \\sim 1 - (\\phi /\\mu)^p$.  Without tuning or a symmetry, the inflaton mass term dominates, and $p = 2$. In this case, we derive an expression for the field variation over the entire inflationary period, Eq~(\\ref{phirun}), or $\\Delta \\phi / m_{\\rm Pl}  \\simeq {1 / (2 \\sqrt{\\pi} \\left\\vert\\eta\\right\\vert)} \\sim 1/(1-n)$. Therefore as we approach the scale-invariant limit, $n \\rightarrow 1$, the field variation becomes large in Planck units. Interestingly, most of the field evolution happens near the {\\em end} of inflation, long after observable perturbations have been generated. The field variation can be small during the first four e-folds of inflation, \\begin{equation} \\Delta\\phi(60 \\rightarrow 56) = 0.05 m_{\\rm Pl}, \\end{equation} thereby nominally satisfying $\\Delta\\phi \\ll m_{\\rm Pl}$.  The field variation is, however, {\\em large} during the last e-fold, \\begin{equation} \\Delta\\phi(1 \\rightarrow 0) = 0.2 m_{\\rm Pl}. \\end{equation} Therefore, such models still suffer from the  ``last e-fold'' problem, and generically yield large field variations near the end of inflation. This behavior can be seen in the approximate flow relation \\begin{equation} \\Delta \\phi \\approx \\frac{m_{\\rm Pl}}{2\\sqrt{\\pi}}\\sqrt{\\epsilon} |\\Delta N|, \\end{equation} valid when $\\epsilon \\sim {\\rm constant.}$ Near the end of inflation, $\\epsilon \\simeq 1$, and the field variation is generically of order $m_{\\rm Pl}$, even for models which generate a negligible tensor amplitude on observable scales. This analytic result is consistent with conclusions based on a stochastic analysis using the flow equations, discussed here and in Ref. \\cite{Efstathiou:2005tq}. Whether or not a field variation $\\Delta\\phi \\sim 0.2 m_{\\rm Pl}$ is actually a problem is a model-dependent statement \\cite{Boyanovsky:2005pw}. One of us (WHK) has argued that field variations of this magnitude are perfectly acceptable in the context of natural inflation \\cite{Freese:2004un}. Other models might incorporate such field variations in a natural way (see, for example, Ref. \\cite{Berkooz:2005sf}).  In order for the field variation during inflation to be very small in Planck units, the mass term for the inflaton must be suppressed. This might be accomplished by fine-tuning: for example a very flat potential with a sharp ``drop off'', so the potential resembles a mesa rather than  a rounded hilltop.  In this case, a Taylor expansion of the potential near the end of inflation is dominated by higher-order terms. (A similar tuning problem was identified in Ref. \\cite{Boyle:2005ug} using a counting argument on a generic dimensionless paremeter set.) A less contrived alternative is to introduce a symmetry which forces the mass term to vanish, so that the potential is dominated by quartic (or higher-order) interactions during the inflationary period. In this case,  $\\epsilon$ varies rapidly near the end of inflation, and $\\Delta\\phi \\ll m_{\\rm Pl}$ over the entire inflationary period, without the need for further tuning.  In Ref. \\cite{Kinney:1995cc} WHK and Mahanthappa constructed a detailed model based on an explicitly broken local symmetry, demonstrating that one can achieve the desired suppression of the inflaton mass term within a realistic Lagrangian. We argue that this model satisfies the ``challenge'' issued by Boyle, Steinhardt, and Turok in Ref. \\cite{Boyle:2005ug} to construct a deterministic, complete inflationary model which is consistent with $r < 10^{-2}$. The model is technically natural: no fine-tuning is required even when loop corrections are included, and the potential \\begin{eqnarray} V\\left(\\theta\\right)&=& {3 v^4 \\over 64 \\pi^2} g^4 \\times \\cr && \\left\\lbrace \\sin^4\\left(\\theta \\over v\\right) \\ln\\left[g^2 \\sin^2\\left(\\theta \\over v\\right)\\right] - \\ln\\left(g^2\\right)\\right\\rbrace \\quad \\end{eqnarray} has a smooth minimum, bounded below. The inflaton can be coupled to fermionic degrees of freedom, allowing reheating, without spoiling the form of the inflationary potential. A discussion of how to achieve a similar suppression of the mass term in extra-dimensional theories can be found in Ref. \\cite{Arkani-Hamed:2003mz}. We therefore conclude that tuning arguments {\\em cannot} be used as suggested by Boyle, {\\it et al.} to place a lower-bound on the tensor/scalar ratio for inflation models.  Conversely, there is no watertight argument that can be used to rule out inflationary models with $\\Delta \\phi > m_{\\rm Pl}$, and the value of $r$ will ultimately be determined observationally. However, the analysis here shows that in typical models of inflation, the variation in the field value during the last e-fold of inflation can be significantly larger than the change during the epoch in which cosmological perturbations are laid down. Consequently, if one is considering a class of models for which $\\phi \\sim m_{\\rm Pl}$ is disfavored, the analysis here can be used to put tight constraints on the models' parameter space.  Observationally, we are making rapid strides. Forthcoming observations will be capable of detecting a tiny deviation from scale-invariance. In the absence of running, a blue spectrum, ($n > 1$) will strongly favor hybrid-type inflation models. A red spectrum ($n<1$) favors ``inverted'' potentials typical of small-field inflation.\\footnote{The WMAP3 data \\cite{Spergel:2006hy}, released after the first version of this paper was written, strongly favor a red spectrum.} Current upper bounds on the tensor amplitude are much weaker than the ultimate theoretical limits on their detectability. If tensor modes are detected by future measurements, we will acquire a wealth of phenomenological information about the inflationary epoch \\cite{Lidsey:1995np,Easther:2002rw}. If the tensor/scalar ratio turns out to be unobservably small, we can still learn much about inflation. However, despite recent progress, there is no guarantee that the theoretical questions surrounding the ``naturalness'' of inflation will be quickly resolved."
    },
    "0601/gr-qc0601007_arXiv.txt": {
        "abstract": "We study spatial variations of the fine-structure constant in the presence of static straight cosmic strings in the weak-field approximation in Einstein gravity. We work in the context of a generic Bekenstein-type model and consider a gauge kinetic function linear in the scalar field.  We determine an analytical form for the scalar field and the string metric at large distances from the core. We show that the gravitational effects of $\\alpha$-varying strings can be seen as a combination of the gravitational effects of global and local strings. We also verify that at large distances to the core the space-time metric is similar to that of a global string. We study the motion of test particles approaching from infinity and show that photons are scattered to infinity while massive particles are trapped in bounded trajectories. We also calculate an overall limit on the magnitude of the variation of $\\alpha$ for a GUT string, by considering suitable cosmological constraints coming from the Equivalence Principle. ",
        "introduction": "$ $\\par  There has been a renewed interest in Cosmology in modeling and constraining possible space-time variations of the so called $``$constants\" of Nature. On the observational side, the analysis of quasar absorption spectra, by means of the many-multiplet method, reveals anomalies which could be interpreted in terms of a cosmological variation of $\\alpha \\equiv e^2/(4\\pi\\hbar\\,c)$ at low redshift \\cite{Webb,Murphy}. However, the situation is presently unclear since other more recent works also based on quasar absorption spectra seem to be consistent with no variation of $\\alpha$ \\cite{Quast,Srianand1,Srianand2}. Meanwhile other studies constrain variations of $\\alpha$ at low redshift such as those based on atomic clocks \\cite{Marion}, meteorites \\cite{Olive} and the Oklo \\cite{Fujii} natural nuclear reactor. At high redshifts variations of $\\alpha$ are constrained by the Cosmic Microwave Background and Big Bang Nucleossynthesis \\cite{hannestad,Avelino,Martins,Rocha,wmap1,sigurdson,wmap2}.  On the theoretical side, most work has been based on Bekenstein-type models \\cite{Bekenstein} which provide a simple framework for studying $\\alpha$ variability. Although a variation of $\\alpha$ should in principle be accompanied by a change in the value of the others fundamental constants as well as the grand unification scale, one usually takes a phenomenological approach (see for example Ref.~\\cite{OlivePos,Anchordoqui,Copeland,Nunes,Sandvik,vslcos,vslcos1,vs1,cosmic,linear,Lee,Parkinson, Bennet,Nosso3,Barrow,Barrow2,Barrow3,Nosso,Magueijo,Nosso2}) neglecting possible variations of the other fundamental constants.  This paper is an extension of Ref.~\\cite{Nosso} where the spatial variations of $\\alpha$ in the vicinity of static Abelian vortex in the context of Bekenstein-type models were investigated. There we considered a scalar field with a generic gauge kinetic function and computed the spatial variations of the fine-structure constant predicted in various models. We studied two special cases (a gauge kinetic function taking an exponential form which reproduces the original Bekenstein model \\cite{Bekenstein}, and polynomial gauge kinetic function) and verified that, for certain classes of models, the classical Nielsen-Olesen vortex is still a valid solution. However, we showed that, in general, static strings solutions depart from the standard Nielsen-Olesen solution with the electromagnetic energy concentrated along the string core seeding the spatial variations of $\\alpha$. We have also shown that Equivalence Principle constraints impose tight limits on the allowed variations of $\\alpha$ induced by cosmic strings networks and calculated a conservative overall limit on the spatial variations of the fine-structure constant seeded by cosmic strings on cosmological scales, $\\Delta\\alpha/\\alpha \\lesssim 10^{-12}$, too small to have any significant cosmological impact. All these results were obtained in Minkowsky space-time.  In the present paper, we generalize this previous study to include the gravitational field of the string. The paper is organized as follows. In section II we describe the Abelian-Higgs model coupled to Einstein gravity in the context of Bekenstein-type models and obtain field equations of motion. In section III we adopt an ansatz to describe static vortex solutions and simplify the field equations. We write the components of the energy-momentum tensor of the various components and discuss the boundary conditions. In section IV we describe the vortex solutions for the metric and for the scalar field that describes the variation of $\\alpha$ far away from the string core. We analyse the non-conical nature of this metric and its similarity with that of a global string. In the following section we study the effect that $\\alpha$ variability has in the geodesics of test particles. In section VI we estimate an overall limit on cosmological variations of the fine-structure constant seeded by these cosmic strings. We compare this result with that obtained in Ref.~\\cite{Nosso}. Finally in section VII we summarize and discuss our results. Throughout this paper we shall use units in which $\\hbar=c=1$ and the metric signature $+\\,-\\,-\\,-$. ",
        "conclusions": "$ $\\par  In this article we studied string vortex solutions in Bekenstein-type models with a gauge kinetic function linear in the scalar field. The spatial variations of $\\alpha$ produce an asymptotic form for the string metric which is a reminiscent of that of global strings. We showed that for GUT strings the non-conical effects in the metric become evident only for very large scales (well beyond the present Hubble radius). We have shown that photons approaching the string from infinite distances to the core will be scattered while massive particles are trapped in bounded trajectories. For GUT strings one gets $\\Delta\\,\\alpha/\\alpha\\,\\lesssim\\, 10^{-12}$ in agreement with the result presented in Ref.~\\cite{Nosso} which is too small to be of any cosmological relevance. We shall leave for future work the study of dilatonic $\\alpha$-varying strings. However, we anticipate that these should be similar to strings in axion-dilatonic gravity \\cite{axion} with a generic kinetic gauge coupling. In this case it should be possible to relate the variations of fine-structure constant with those of the gravitational constant."
    },
    "0601/astro-ph0601089_arXiv.txt": {
        "abstract": "We present near-infrared observations of 71 newly discovered L and T dwarfs, selected from imaging data of the Sloan Digital Sky Survey (SDSS) using the $i$-dropout technique. Sixty-five of these dwarfs have been classified spectroscopically according to the near-infrared L dwarf classification scheme of Geballe et al. and the unified T dwarf classification scheme of Burgasser et al.  The spectral types of these dwarfs range from L3 to T7, and include the latest types yet found in the SDSS.  Six of the newly identified dwarfs are classified as early- to mid-L dwarfs according to their photometric near-infrared colors, and two others are classified photometrically as M dwarfs.  We also present new near-infrared spectra for five previously published SDSS L and T dwarfs, and one L dwarf and one T dwarf discovered by Burgasser et al.\\ from the Two Micron All Sky Survey. The new SDSS sample includes 27 T dwarfs and 30 dwarfs with spectral types spanning the complex L--T transition (L7--T3).  We continue to see a large ($\\sim$ 0.5 mag) spread in $J$--$H$ for L3 to T1 types,  and a similar spread in $H$--$K$ for all dwarfs later than L3.  This color dispersion is probably due to a range of grain sedimentation properties, metallicity, and gravity.  We also find L and T dwarfs with unusual colors and spectral properties that may eventually help to disentangle these effects. ",
        "introduction": "Over the last decade, the search for rare astronomical objects has undergone an explosion of productivity, thanks in part to the specialized labor associated with modern digital sky surveys. Innovative programmers and instrumentalists designing highly automated software pipelines and large imaging arrays have laid the groundwork for massive databases that can now be mined for objects once only thought to exist.   The result of these efforts has been the most productive period in history for the discovery of rare objects, which range from the brightest and most distant quasars \\citep{fan99,fan00a,fan01a,fan01b,fan03,fan04,zheng00} to the faintest and nearest stars and brown dwarfs (\\citealt[hereafter G02]{k99,burg99,k00,fan00b, leggett00,burg02a,g02}; \\citealt[hereafter K04]{k04}).  \\begin{figure*}[t] \\epsscale{0.8} \\figurenum{1} \\plotone{f1.b.ps} \\caption{SDSS $z$ finding charts for the 73 new M, L, and T dwarfs found in this work. North is up, East is left.  The fields of view are $3' \\times 3'$.  (Full resolution versions of these finding charts will appear in the published journal edition.)} \\end{figure*}   \\begin{figure*}[t] \\epsscale{0.8} \\figurenum{1} \\plotone{f2.b.ps} \\caption{(continued) SDSS $z$ finding charts for the 73 new M, L, and T dwarfs found in this work. North is up, East is left.  The fields of view are $3' \\times 3'$.   (Full resolution versions of these finding charts will appear in the published journal edition.) } \\end{figure*}   \\begin{figure*}[t] \\epsscale{0.8} \\figurenum{1} \\plotone{f3.b.ps} \\caption{(continued) SDSS $z$ finding charts for the 73 new M, L, and T dwarfs found in this work. North is up, East is left.  The fields of view are $3' \\times 3'$. (Full resolution versions of these finding charts will appear in the published journal edition.)} \\end{figure*}  \\begin{figure*}[t] \\epsscale{0.8} \\figurenum{1} \\plotone{f4.b.ps} \\caption{(continued) SDSS $z$ finding charts for the 73 new M, L, and T dwarfs found in this work. North is up, East is left.  The fields of view are $3' \\times 3'$.  (Full resolution versions of these finding charts will appear in the published journal edition.)} \\end{figure*}   \\begin{figure}[h] \\epsscale{1.1} \\figurenum{2} \\plotone{f5.ps} \\end{figure}   \\begin{figure}[h] \\epsscale{1.1} \\figurenum{2} \\plotone{f6.ps} \\caption{Representative optical and near-infrared spectra of the L and T dwarfs discovered in this work.  The typical spectral resolution is $\\sim150$.   SDSS designations are given above the dashed line corresponding to the zero flux level.  The locations of significant spectral features are indicated at the top of each panel. } \\end{figure}  Since the first discoveries of L and T dwarfs in the early 1990s, the ranks of these spectral classes have been populated through systematic searches of large-area, optical and near-infrared surveys such as the Deep Near Infrared Survey of the Southern Sky \\citep[DENIS;][]{epchtein97}, the Two Micron All Sky Survey \\citep[2MASS;][]{skrutskie97}, and the Sloan Digital Sky Survey \\citep[SDSS;][]{york}.  Although time intensive, these searches are more productive and efficient than the tedious, pointed surveys that preceded them.  The optical and near-infrared spectra of the many known L and T dwarfs have yielded well-defined classification schemes and valuable information about the chemical compositions, temperatures, and other physical characteristics of ultracool dwarfs (\\citealt{k00,burg02a,burg05}; G02; K04, \\citealt{gol04a}).  L and T dwarfs, which link the lowest mass stars and the highest mass planets, are the subjects of a broad range of observational and theoretical studies \\citep{k05}.  A principal goal of these studies is the determination of the substellar mass function, which is fundamental to understanding the star formation process at the lowest mass end and how it is related to the formation and evolution of planetary systems \\citep{burg04b,allen05}.  The hundreds of L dwarfs and more than 50 T dwarfs discovered from 2MASS and SDSS data compose a statistical sample of brown dwarfs essential to understanding the demographics and eventually the luminosity and mass functions of substellar objects (Fan et al., in prep.).  Since 1999, we have used SDSS imaging data to select L and T dwarf candidates and carry out follow-up observations of their photometric and spectroscopic properties (\\citealt{strauss99, tsvetanov00,fan00b,leggett00,leggett02,schneider02}; G02; \\citealt{hawley}; K04; \\citealt{gol04a}). The SDSS $i$--$z$ color, which is our primary selection criterion, is sensitive to brown dwarfs of all effective temperatures as long as they are bright enough to be detected in the $z$ band. Near-infrared selection criteria, such as those employed by the 2MASS group \\citep[e.g.,][]{burg02a}, are more sensitive to late-T dwarfs which are optically faint, but they are less efficient in detecting dwarfs at the L--T transition because their near-infrared colors are indistinguishable from those of the more numerous M dwarfs.  Therefore, the SDSS selection criterion is better suited for defining a complete, magnitude-limited sample of field brown dwarfs.  In this paper, we report the discovery and properties of 71 L and T dwarfs drawn from over 3500 deg$^2$ of SDSS imaging data.  We describe the techniques for identifying and confirming these dwarfs based on their SDSS and near-infrared colors.  We classify 65 of the L and T dwarfs from their near-infrared spectra, and we identify 6 L dwarfs and 2 M dwarfs from their photometric colors.  The observations are presented in \\S 2, spectral classification of the sample is presented in \\S 3, and variations in the photometric and spectroscopic  properties are discussed in \\S 4.  Our conclusions are given in \\S 5. ",
        "conclusions": "Thanks to the great gains in detection efficiency provided by large digital sky surveys, observers have reached the once-unimaginable position of routinely discovering nearby field brown dwarfs.  In this paper, we report the discovery of 71 L and T dwarfs identified during our ongoing search for high-redshift quasars and brown dwarfs in SDSS imaging data.  Using near-infrared photometry and spectra obtained over the past two years at UKIRT and IRTF, we have classified these dwarfs using the L classification scheme of \\citet{g02} and the new unified T classification scheme of \\citet{burg05}.  The near-infrared colors of the 71 dwarfs exhibit the same trends, scatter, and anomalies noted in previously reported samples.  The unusual spectral features exhibited by some dwarfs may lead to a better understanding of their underlying atmospheric properties.  We plan to apply our results to the revision of the spectral infrared indices used in the near-infrared classification of L dwarfs.  We will also integrate the 71 new L and T into a complete magnitude-limited catalog for the purpose of determining the substellar luminosity function."
    },
    "0601/astro-ph0601054_arXiv.txt": {
        "abstract": "The discovery by {\\em Swift} that a good fraction of Gamma Ray Bursts (GRBs) have a slowly decaying X-ray afterglow phase led to the suggestion that energy injection into the blast wave takes place several hundred seconds after the burst. This implies that right after the burst the kinetic energy of the blast wave was very low and in turn the efficiency of production of $\\gamma$-rays during the burst was extremely high, rendering the internal shocks model unlikely. We re-examine the estimates of kinetic energy in GRB afterglows and show that the efficiency of converting the kinetic energy into $\\gamma-$rays is moderate and does not challenge the standard internal shock model. We also examine several models, including in particular energy injection, suggested to interpret this slow decay phase.  We show that with proper parameters, all these models give rise to a slow decline lasting several hours. However, even those models that fit all X-ray observations, and in particular the energy injection model, cannot account self-consistently for both the X-ray and the optical afterglows of well monitored GRBs such as GRB 050319 and GRB 050401. We speculate about a possible alternative resolution of this puzzle. ",
        "introduction": "\\label{sec:XRF1} The X-ray telescope (XRT) on board {\\em Swift} has provided high quality early X-ray afterglow light curves of many Gamma-ray Bursts (GRBs). One of the most remarkable and unexpected features discovered by {\\em Swift} was that many of these X-ray afterglow light curves are distinguished by a slow decline---The flux $F$ decreases with observer's time $t$ as $F\\propto t^{[0,-0.8]}$, lasting from a few hundred seconds to few hours (Nousek et al. 2005, Campana et al. 2005; Vaughan et al. 2005; Cusumano et al. 2005; de Pasquale et al. 2005).  Such a phase is unexpected in the standard fireball model. A simple explanation is that the slow decline arises due to a significant energy injection (Zhang et al. 2005; Nousek et al. 2005; Panaitescu et al. 2006, Granot \\& Kumar 2006), as suggested previously (For baryon-rich injection: Rees \\& M\\'esz\\'aros 1998; Panaitescu et al. 1998; Kumar \\& Piran 2000; Sari \\& M\\'esz\\'aros 2000; Zhang \\& M\\'esz\\'aros 2002; Granot, Nakar \\& Piran 2003. For Poynting flux dominated injection\\footnote{If the outflow ejected from the central engine after the gamma ray burst phase is highly magnetized, at a radius $\\sim 10^{15}$ cm, the MHD condition breaks down. Significant magnetic field dissipation processes are expected to happen which converts energy into radiation. As long as the highly magnetized outflow is steady enough, strong and slowly decaying X-ray emission is possible (see Fan, Zhang \\& Proga [2005a] and the references therein).}: Dai \\& Lu 1998a; Zhang \\& M\\'esz\\'aros 2001; Dai 2004). It has been argued that consequently the resulted GRB efficiency, i.e., the ratio of the energy emitted in  $\\gamma-$ray energy to the total energy (the sum of the $\\gamma-$ray energy and the kinetic energy of the ejecta powering the afterglow), should be $90\\%$ or higher. Some extreme assumptions are needed (Beloborodov 2000; Kobayashi \\& Sari 2001) to reach such a high efficiency within the framework of the standard internal-shocks model (Paczynski \\& Xu 1994; Rees \\& M\\'esz\\'aros 1994; Sari \\& Piran 1997a , 1997b; Kobayashi, Piran \\& Sari 1997; Daigne  \\& Mochkovitch 1998; Piran 1999).  We re-examine this issue focusing on two critical aspects of the analysis. The estimate of the kinetic energy of the ejecta from the afterglow observations and in particular from the X-ray flux and the need of energy injection.  We show in \\S\\ref{sec:eff} that even for these {\\em Swift} GRBs with long duration X-ray flattening the $\\gamma$-ray conversion efficiency is high but not unreasonable.   We then turn to the puzzling slow decline seen in the first few hours of the X-ray afterglow. We explore in  \\S\\ref{sec:XRF2} several models that may give rise to slowly decaying X-ray afterglows: (i) Energy injection. (ii) A small $\\zeta_e$, in which only a small fraction, $\\zeta_e \\ll 1$ of the electrons are accelerated to high energies and contribute to the radiation process. (iii) Evolving shock parameters, where the microscopic shock parameters $\\epsilon_e$ and/or $\\epsilon_B$ (the fraction of shock energy given to the magnetic filed) vary in time and are inversely proportional to the Lorentz factor of the ejecta. (iv) A very low variable external density model, in which the number density of the medium is not only very low but it also a function of the radius. (v) Highly magnetized outflow where flattening might arise because of a slow conversion of the magnetic energy to kinetic energy of the external matter. We present in \\S\\ref{sec:XRF2} analytical derivation as well as numerical calculations of the expected light curves in all these models except the last one. In \\S\\ref{sec:0319} we compare the models to the observations of GRB 050319 and  GRB 050401. We summarize our results and discuss their implications in \\S\\ref{sec:XRF4}. We conclude with a speculation on the nature of the solution to this puzzle. ",
        "conclusions": "  1. The GRB efficiency of pre-{\\it Swift} GRBs that has been derived directly from the X-ray flux 10 hours after the burst has been overestimated. For these {\\it Swift} GRBs with long time X-ray flattening, the GRB efficiency is also moderate (around 0.5), even when taking into account the possibility of energy injection. Such efficiency can be understood within the standard internal shock model.  2. With a proper choice of parameters, the slow decline slope of X-ray afterglow like the one  detected in GRB 050319 can be well reproduced by several models---the energy injection model, evolving shock parameter model (in which the shock parameters are assumed to increase with the decrease of the shock strength for $t<10^4$ s), the small $\\zeta_e$ model (in which the shock energy has been give to a fraction $\\zeta_e$ of electrons, rather than total) and the very low nonconstant density model. Out of these models, the last two are ruled out by the X-ray data itself. In the last model, the resulting X-ray afterglow is too dim to match most XRT observations. The small $\\zeta_e$ model is also disfavored since (1) In the slow decline phase, the XRT spectrum are usually much softer than $\\nu^{-1/2}$; (2) After the light curve break, no spectral steepening has been detected in most cases, which is inconsistent with the model. The other models, including the energy injection model and the evolving shock parameter model seem to be consistent with the X-ray afterglow observations.  3. While two  models: the energy injection model and  the evolving shock parameter model are consistent with the X-ray data, they fail to reproduce both the X-ray and the optical afterglows of GRB 050319 and GRB 050401. In each burst, the optical flux declines slowly up to $\\sim 10^5$ s. On the other hand, the X-ray light curve decays slowly up to $t\\sim 10^4$ and then turns to the normal faster decay ($F \\propto t^{-1.2}$ or so). The temporal index of the slow decay X-ray phase is close to that of the optical light curve. The XRT spectrum is unchanged before and after the X-ray break. This means that the break is not caused by a cooling break in which $\\nu_c$ crosses the observed frequency. .  The failure of all models that we considered to fit both the X-ray and the optical afterglow light curves suggests that we should look for another alternative. An intriguing possibility is based on fact that the extrapolation backwards of the late X-ray light curve is in agreement with (or 1 to 2 order lower than) the prompt X-ray emission. This suggests that we face a \"missing energy problem\". Namely, during the slow decay phase (in which the X-ray flux is rather low) we miss X-ray emission. Is it possible that during this phase this energy is dissipated into a different channel and not into Synchrotron X-rays and that this different channel becomes ineffective at around ten hours? Put differently, during this phase the electrons within the forward shock emit Synchrotron X-rays inefficiently. A possibility of this kind (that we have considered and found not to work) is if the X-ray emitting electrons are cooled efficiently via inverse Compton (and hence their Synchrotron X-ray emission is weaker). As already mentioned inverse Compton cooling is important in determining the X-ray flux. Furthermore, due to the Klein-Nishina cutoff this cooling becomes unimportant at approximately one day. However, this transition is not sharp enough to produce the observed slowly decaying X-ray light curves. While inverse Compton cooling does not work it is possible that another, yet unexplored, process of this kind is responsible for the observed light curves."
    },
    "0601/astro-ph0601262_arXiv.txt": {
        "abstract": "We study the effect of rotation during the collision between dust aggregates, in order to address a mismatch between previous model calculations of Brownian motion driven aggregation and experiments. We show that rotation during the collision does influence the shape and internal structure of the aggregates formed.  The effect is limited in the ballistic regime when aggregates can be considered to move on straight lines during a collision.  However, if the stopping length of an aggregate becomes smaller than its physical size, extremely elongated aggregates can be produced. We show that this effect may have played a role in the inner regions of the solar nebula where densities were high. ",
        "introduction": "Dust aggregation plays a central role in most theories of planet formation. While for the formation of giant planets, disk instability scenarios continue to be discussed \\citep{1997Sci...276.1836B}, dust aggregation certainly stands at the beginning of the formation of planetesimals and therefore of the terrestrial planets.  Independent of the detailed process that finally forms planetesimals \\citep[e.g.][]{1980Icar...44..172W, 2002ApJ...580..494Y,2001ApJ...546..496C}, dust grains first need to grow and settle towards the midplane of the disk.  Planet formation therefore starts with the first steps of dust aggregation, when dust grains inherited by the disk from the interstellar medium start to collide and grow.  This first step of growth is governed by Brownian motion aggregation.  In the high density regions of the disk, dust/gas coupling is so tight that the main source of relative velocities between grains are the random motions produced by individual collisions between gas atoms and the grains. Understanding the Brownian motion phase of dust aggregation therefore means understanding the first step towards planets in disks.  At the very low collision velocities produced by Brownian motion ($\\sim$cm/s), grains always stick and no restructuring is occurring in the collision \\citep{1997ApJ...480..647D,2000Icar..143..138B}. The structure of aggregates formed in this regime is therefore indeed a pure probe of the physical processes driving growth with Brownian motion of the particles.  Theoretically, growth by Brownian motion was studied for example by \\citet{1999Icar..141..388K}. They calculated the motion of micron sized dust grains enclosed in a box of approximately constant dust number density.  Diffusion, caused by the presence of a gaseous medium, produces relative motion of grains and leads to the growth of dust.  The orientation of the (spherically not symmetric) aggregates was not followed during the computations. Instead, in order to randomize the relative orientation during collisions, the orientations of collision partners just before a collision was selected randomly. With these initial conditions, the aggregates were left to collide, without considering rotation \\emph{during} the collision.  The calculations show a slow growth of the dust aggregates with time.  At any time, the box contains a distribution of aggregate shapes which is characterized by a distribution of fractal dimensions.  The mean fractal dimension found in the numerical study is around $D_f = 1.8$.   On the experimental side, a series of low-gravity experiments has been conducted \\citep{2000PhRvL..85.2426B,2004PhRvL..93b1103K} in order to study the growth of fractal aggregates under Brownian motion conditions.  The results confirm many expected aspects of this growth regime, but also showed unexpectedly elongated aggregates, in contrast with the predictions by \\citet{1999Icar..141..388K}.  The authors qualitatively argue that rotation of the aggregates during the collision might lead to more elongated aggregates, because the probability of achieving contact between two aggregates at larger separations increases if the aggregates rotate.  To compare experimental and theoretical results, it is necessary to quantify the visual impression of \\emph{elongation}.  In this study we will be using two different ways to do so:  \\begin{enumerate} \\item\\label{item:1} One can define an \\textbf{elongation factor} as the ratio of maximum to minimum diameter of aggregate \\begin{equation} \\label{eq:1} f_\\mathrm{el}=\\frac{d_\\mathrm{max}}{d_\\mathrm{min}} \\end{equation} where $d_\\mathrm{min}$ is measured in perpendicular direction to $d_\\mathrm{max}$.  Experimentally, this quantity has to be derived from a few, or even a single picture.  When measuring this value for model aggregates, we therefore choose three different projections of the aggregate and determine $d_\\mathrm{max}$ and $d_\\mathrm{min}$ from these images.  \\item We can also derive the \\textbf{fractal dimension} of the aggregates.  This quantity can be computed by measuring the mass--radius relation of aggregates of different mass $m$ and size $R$. If that relation follows a power law \\begin{equation} m \\propto R^{D_\\mathrm{f}} \\label{D_f} \\end{equation} then $D_\\mathrm{f}$ is called the fractal dimension of the aggregate.  The fractal dimension measured in the experiments is 1.4, and typical values predicted by the numerical calculations are around 1.8, indicating a mismatch. \\end{enumerate}  While the idea of particle rotation is appealing as a mechanism to increase the elongation of dust aggregates, this has never been shown quantitatively.  In this paper we investigate the influence of rotation on the structure of aggregates formed under Brownian motion conditions. ",
        "conclusions": "We have studied the effect of aggregate rotation during collisions. The results show that rotation must definitely be treated correctly when modeling growth due to aggregation, or the geometrical structure of the resulting aggregates will be incorrect.  Rotation plays a role because it enhances the probability that two approaching aggregates make contact early on during the collision, between outer constituent grains. For the most simple case of ballistic collisions, during which the colliding aggregates move on a linear path, ignoring aggregate rotation increases the fractal dimension from 1.46 to 1.7 in the limit of pure cluster-cluster aggregation.  This effect becomes strongly enhanced if the density of the surrounding medium becomes so large that the stopping length of an aggregate becomes shorter than the aggregate size, i.e. in the non--ballistic limit.  We find that reducing the stopping length to the monomer radius results in a fractal dimension of 1.25.  When the stopping length is reduced to one tenth of the monomer radius, the fractal dimension drops to 1.1.  For the solar nebula, we find that non--ballistic collisions play a role in the innermost regions of the disk for even very small aggregates made of a few $\\mu m$-sized grains.  For smaller monomers and/or further away from the Sun, enhanced elongation can be expected for larger aggregates, made of a few 100 or 1000 grains. In the outer disk, all collisions may be considered ballistic."
    },
    "0601/astro-ph0601112_arXiv.txt": {
        "abstract": "We present a deep photometric ($B$- and $R$-band) catalog and an associated spectroscopic redshift survey conducted in the vicinity of the Hubble Deep Field South. The spectroscopy yields 53 extragalactic redshifts in the range $0<z<1.4$ substantially increasing the body of spectroscopic work in this field to over 200 objects.  The targets are selected from deep AAT prime focus images complete to $R<24$ and spectroscopy is 50\\% complete at $R=23$. There is now strong evidence for a rich cluster at $z\\simeq 0.58$ flanking the WFPC2 field which is consistent with a known absorber of the bright QSO in this field. We find that photometric redshifts of $z<1$ galaxies in this field based on HST data are accurate to $\\sigma_z/(1+z)=0.03$ (albeit with small number statistics).  The observations were carried out as a community service for Hubble Deep Field science, to demonstrate the first use of the `nod \\& shuffle' technique with a classical multi-object spectrograph and to test the use of `microslits' for ultra-high multiplex observations along with a new VPH grism and deep-depletion CCD.  The reduction of this new type of data is also described. ",
        "introduction": "\\def\\todo#1{{[\\bf TODO: #1]}}  \\def\\REF{\\todo{REF}}  The Hubble Deep Field South (HDF-S) is one of the deepest imaging fields in the sky. It was observed in 1998 (Williams et al. 2000) by the Hubble Space Telescope (HST) as a counterpart to the northern Hubble Deep Field (HDF-N). However in contrast to the northern field it was selected to contain a bright QSO J2233-606 at $z=2.24$ in order to facilitate studies of the connection between foreground galaxies and absorbing systems in the QSO spectrum.  The determination of precise redshifts for extra-galactic sources has been important since the time of Hubble (1929). While much spectroscopy has been done in the other Hubble Deep Field, HDF-S has lagged behind. As a community service and in order to test new instrumentation techniques, in particular `nod \\& shuffle' (Glazebrook \\& Bland-Hawthorn 2001; GB01) with very small `microslits', we carried out a spectroscopic campaign of galaxies in, and in the vicinity of, the HDF-S in order to obtain redshifts.  The plan of this paper is as follows: Section 2 details the prime focus pre-imaging, the procedures used to construct the photometric catalog and the selection of the spectroscopic targets. In Section 3 we describe our novel spectroscopic observations and the special data reduction procedures used. In Section 4 we present our spectroscopic catalog and its basic parameters and compare with other work in this field and discuss the possible galaxy cluster at $z\\simeq 0.58$.  The finding charts, images, spectra, photometric and spectroscopic catalogs presented in this paper are all available from the Anglo-Australian Observatory web site at:  {\\tt http://www.aao.gov.au/hdfs/Redshifts} ",
        "conclusions": "To summarize the paper:  \\begin{enumerate} \\item We present a significant new set of photometric and spectroscopic data on the HDF-S. \\item We have demonstrated a new mode of spectroscopic data taking with high-multiplex per unit area on the sky using the technique of `microslits' in conjunction with nod \\& shuffle. \\item The spectroscopic catalog presented here is the third large catalog of redshifts in the HDF-S and flanking fields. 53 redshifts are obtained of which 45 are new. \\item Comparison with existing photometric redshifts based on WFPC-2 data show these to be in reasonable agreement ($\\sigma_z /(1+z)=0.03$) for $z<1$ galaxies, albeit with small number statistics. \\item We present strong evidence for a rich cluster, compact in all three spatial dimensions, at $z\\simeq 0.58$ flanking the WFPC2 pointing and which very likely contains the known absorbing galaxy of the QSO at this redshift. \\end{enumerate}  Preliminary versions of this catalog have already been used for scientific studies using the HDF-S data. Mann et al. (2002) used it for a study of the relationship between far-infrared and other estimates of galaxy star-formation rates. Vanzella et al. (2002) used this catalog together with theirs to estimate cosmological star-formations rate history. The reader is encouraged to use this catalog for further studies of the HDF-S and its associated QSO."
    },
    "0601/astro-ph0601438_arXiv.txt": {
        "abstract": "{ Simple geometrical ring models account well for near-infrared interferometric observations of dusty disks surrounding pre-main sequence stars of intermediate mass. Such models demonstrate that the dust distribution in these disks has an inner hole and puffed-up inner edge consistent with theoretical expectations. } { In this paper, we reanalyze the available interferometric observations of six intermediate mass pre-main sequence stars (CQ Tau, VV Ser, MWC 480, MWC 758, V1295 Aql and AB Aur) in the framework of a more detailed physical model of the inner region of the dusty disk. Our aim is to verify whether the model will allow us to constrain the disk and dust properties.} { Observed visibilities from the literature are compared with theoretical visibilities from our model. With the assumption that silicates are the most refractory dust species, our model computes self-consistently the shape and emission of the inner edge of the dusty disk (and hence its visibilities for given interferometer configurations). The only free parameters in our model are the inner disk orientation and the size of the dust grains.} { In all objects with the exception of AB Aur, our self-consistent models reproduce both the  interferometric results and the near-infrared spectral energy distribution. In four cases, grains larger than  $\\sim$1.2\\um, and possibly much larger are either required by or consistent with the observations. The inclination of the inner disk is found to be always larger than $\\sim 30^{\\circ}$, and in at least two objects much larger.} ",
        "introduction": " ",
        "conclusions": "In this paper we have analyzed the near-infrared interferometric observations of the six best observed HAe stars using the rim models developed by Isella \\& Natta (2005). Our  aim was to explore the potential of near-infrared interferometry to constrain the properties of the grains in the inner disks of these stars.  The basic assumptions of the IN05 rim models are that the inner disk structure is controlled by the evaporation of dust in the unattenuated stellar radiation field, as expected if the gaseous disks have low optical depth, and that the most refractory grains are silicates. The IN05 self-consistent models for the ``puffed-up'' inner rim reproduce both the interferometric observations and the near-infrared spectral energy distribution of all the objects we have studied, with the exception of AB Aur, which we have briefly discussed.  For the five stars where we are able to obtain a good fit to the data, we can estimate the grain sizes in the rim, i.e., in the midplane of the inner disk. We find that in  four cases grains larger than $\\sim 1.2$\\um\\ are either required by or consistent with the data. Only in one case do we find that the existing data require $a\\sim 0.2-0.3$\\um. Note that  this value of $a=1.2$\\um\\  is  a lower limit to the grain size: grains can be much larger, since the rim location and shape do not change significantly if the grains grow further.  As a result of the model-fitting, one derives also the inclination and position angle of the disk on the plane of the sky. We find that, in general, these parameters are not well constrained by the existing data. However, in all cases we can fit, inclinations lower than $30^{\\circ}$ are not consistent with the observations and the surface brightness distribution can not be circularly symmetric.  This rules out a spherical envelope as the dominant source of the near-infrared emission.  For some objects, estimates of the inclination of the outer disk  have been obtained from millimeter interferometric observations; within the uncertainties, they agree with the values obtained for the inner disk.  Our analysis shows that near-infrared interferometry is a very powerful tool for understanding the properties of the inner disks, in particular when combined with physical models of these regions. However, at present the existing  data are for many objects still too sparse in their coverage of the u-v plane to allow an accurate determination of  the disk parameters. We expect that this will be improved in the future.  In this context, since near-infrared interferometry is and will remain a very time demanding  technique, we stress the importance of using physical models of the inner region of the disk in planning future observations."
    },
    "0601/astro-ph0601324_arXiv.txt": {
        "abstract": "Renewed interest in the first stars that were formed in the universe has led to the discovery of extremely iron-poor stars. Since several competing scenarios exist, our understanding of the mass range that determines the observed elemental abundances remains unclear. In this study, we consider three well-studied metal-poor stars in terms of the theoretical supernovae (SNe) model. Our results suggest that the observed abundance patterns in the metal-poor star BD $+80^\\circ$245 and the pair of stars HD 134439/40 agree strongly with the theoretical possibility that these stars inherited their heavy element abundance patterns from SNe initiated by thermonuclear runaways  in the degenerate carbon-oxygen cores of primordial asymptotic giant branch stars with masses of $\\sim 3.5-5$ \\msp. Recent theoretical calculations have predicted that such SNe could be originated from metal-free stars in the intermediate mass range. On the other hand, intermediate mass stars containing some metals would end their lives as white dwarfs after expelling their envelopes in the wind due to intense momentum transport from outgoing photons to heavy elements. This new pathway for the formation of SNe requires that stars are formed from the primordial gas.  Thus, we suggest that stars of a few solar masses were formed from the primordial gas and that some of them caused thermonuclear explosions when the mass of their degenerate carbon-oxygen cores increased to the Chandrasekhar limit without experiencing efficient mass loss. ",
        "introduction": "Elemental abundance of recently discovered extremely iron-deficient stars has raised questions about the mass range of the first stars that were formed in the universe \\citep{Christlieb_02, Umeda_03, Schneider_03, Shigeyama_03, Suda_04, Frebel_05, Iwamoto_05}. General theoretical considerations seem to suggest that, due to inefficeint cooling of the primordial gas, the first stars were at least as massive as 100 \\ms  \\citep[e.g.,][]{Abel_02,Bromm_02}. However, other studies claim that molecular or atomic hydrogen cooling is sufficient to form low-mass stars from the primordial gas \\citep{Nakamura_01,Omukai_03}. Therefore, the detection of signatures constraining the mass range of the first stars limits the possible formation mechanisms of these objects.  The key parameters controlling the structure and evolution of stars are mass and metallicity. A given combination of these two parameters establishes new pathways for supernova (SN) formation from  {\\it ``primordial intermediate mass\"} stars. According to recent calculations \\citep{Fujimoto_00, Suda_04}, stars with [Fe/H] \\ltsim $-5$ and masses in the range of 3.5\\ms \\ltsim $M$ \\ltsim 7 \\ms do not experience mixing episodes including the third dredge-up, in their convective envelope in the thermally pulsing asymptotic giant branch (TPAGB) phase. These stars will end up with a small amount of carbon as well as $s$-process elements in their envelopes. As a result, these stars are less opaque and more compact than stars that have experienced the third dredge-up. These features also result in stars with lower radiation pressures and stronger surface gravities. Therefore, it is expected that the mass loss from an extremely metal-poor star will be substantially suppressed \\citep{Fujimoto_84}. Thus, a degenerate carbon-oxygen core can reach the Chandrasekhar limit  ($M_{\\rm ch}$) while the star is in the TPAGB phase.  Eventually the star undergoes a thermonuclear explosion to become an SN. This type of SN can be classified as a type Ia in terms of nucleosynthesis but as a type II in terms of spectral type. If the mass of a proto-galactic cloud is $10^6$ \\ms corresponding to the Jeans mass soon after recombination, a typical Type II SN (SN II)  enriches the cloud up to [Fe/H]$\\sim -4$. Therefore only the primordial (population III) stars are likely to end up as such SNe. Hence, we define them as SNe IIIa.  \\begin{figure*}[t] \\begin{center} \\includegraphics[width=13.5cm,clip=true]{f1.eps} \\end{center} \\caption{{\\it Left panel}: Correlation of [Na/Mg] with [Mg/H] for BD+80$^\\circ$245 and HD 134439/40 \\citep[filled circles;][Fulbright 2000]{Ivans_03} together with other Galactic stars \\citep[open triangles;][crosses; Fulbright 2000]{Cayrel_04}. No abundance anomaly is seen for the three stars. {\\it Right panel}: The same as left panel but for [Ba/Mg]. The data shown by open triangles are taken from \\citet{McWilliam_98}.} \\end{figure*}  The lifetime of an SN IIIa progenitor is much shorter than that of an SN Ia progenitor, though these two types of SNe supply similar amounts of nuclear products. In a white dwarf, an SN Ia event needs substantial replenishment of the gas by accretion from a low-mass companion. As a result, SN Ia progenitors have much longer evolutionary time-scales than those of SNe II $-$, typically longer than $10^9$ yrs $-$. This prevented SNe Ia from contributing to chemical enrichment in the early epoch of galaxy evolution. On the other hand, an SN IIIa event occurs when an intermediate mass star evolves into the AGB phase in a period ranging from several $10^7$ yrs to a few $10^8$ yrs, depending on the progenitor mass. Such short time scales could enable us to detect evidence of the above in the elemental abundances of old stars. Recent studies of the chemical composition of metal-poor stars have revealed that they have inherited the abundance pattern of the ejecta of the preceding SNe \\citep[e.g.,][]{Audouze_95, Shigeyama_98, Umeda_03}. If this is the case, there might be some later stars that inherited the nuclear products of SNe IIIa, though their abundance patterns could be modified to some degree by products from SNe II.  The most characteristic abundance feature of stars inheriting the nuclear products of an SNe IIIa event is expected to be the deficiency of $\\alpha$-elements, compared with the enhanced $\\alpha$/Fe ratios commonly seen in halo stars. In fact, there exist so-called low-$\\alpha$ stars in the Galactic halo with low $\\alpha$/Fe ratios. It has been proposed \\citep{Nissen_97, Gilmore_98} that these low-$\\alpha$ stars were born from accreted dwarf galaxies or protogalactic fragments, where a prolonged low star formation rate enables SNe Ia events to occur in a low-metallicity regime. The specific mechanism which explains the low $\\alpha/$Fe ratios in an individual star is easily determined by the abundance of $s$-process elements, since the time scales associated with these two types of SNe are substantially different. The long lifetime of SNe Ia allows AGB stars to substantially pollute the gas with $s$-process elements before low-$\\alpha$ stars are formed, whereas SNe IIIa preceded most AGB stars with masses \\ltsim$3.5\\,M_\\odot$.  In the literature, we found three stars (BD+80$^\\circ$245 and HD 134439/40) that exhibit the abundance patterns expected from SNe IIIa. Elemental features of these three stars are discussed in the subsequent two sections. Their origin is investigated in the context of a chemical evolution model developed by \\citet{Tsujimoto_99} in \\S \\ref{ee}. ",
        "conclusions": "\\begin{figure*}[t] \\begin{center} \\includegraphics[width=14.2cm,clip=true]{f4.eps} \\end{center} \\caption{{\\it Left panel}: The same as Figure 3 but for [Ba/Fe]; solid and dashed curves correspond to the swept-up mass of $6.5\\times10^4$\\ms and $8\\times10^4$\\msp, respectively. Dashed-dotted line shows the evolution of the abundance ratios in the gas.  {\\it Right panel}: The same as left panel but for [Eu/Fe]. The observational data for BD+80$^\\circ$245 are taken from \\citet{Fulbright_00} instead of \\citet{Ivans_03} since \\citet{Fulbright_00} gives a consistent $r$-process Ba/Eu ratio.} \\end{figure*}  The unusually low abundances of $\\alpha$- and neutron-capture elements relative to Fe seen in the metal-poor star BD $+80^\\circ$245 and the pair of stars HD 134439/40 are shown to be fossil records of SN explosions triggered by a thermonuclear runaway in a metal free intermediate mass star. These SNe have two distinct characteristics. First, they can be classified as type Ia in terms of nucleosynthesis, that is, they synthesize a large amount of Fe as a result of the thermonuclear explosions. Second, for their cores to grow to the Chandrasekhar limit during the TPAGB phase, the progenitors need to be metal-free stars. Since the three stars exhibit these two characteristics, ``a sufficient Fe supply on a short time-scale\" and ``primordiality\" in their elemental abundance patterns, this strongly suggests that there must be metal-free stars of a few solar masses.  We also predict that stars inheriting heavy elements from SNe IIIa populate distinct branches in the [$\\alpha$/Fe]$-$[Fe/H] and [$n$-capture/Fe]$-$[Fe/H] planes. The manner in which these stars populate this branch should give valuable information on the initial mass function of metal-free stars. Stars descended from SNe IIIa are predicted to have metallicities in the range of $-2.3$\\ltsim[Fe/H]\\ltsim$-1.3$. Thus, the eight low-$\\alpha$ halo stars with $-1.3\\leq[{\\rm Fe/H}] \\leq -0.8$ \\citep{Nissen_97}, and D119 with [Fe/H] =$-2.95$ exhibiting extreme deficiencies of Ca, Sr, and Ba in the Draco dwarf spheroidal galaxy \\citep {Fulbright_04}, are unlikely to be associated with the proposed SNe IIIa.  Theoretically during the formation of galaxies, stars form in the protogalactic clouds assembling to construct the Galactic halo. A spread of several billion years in the stellar age among the halo stars \\citep{Schuster_89} suggests that heavy elements dispersed into the halo from the protoclouds should have polluted other clouds before star formation had started there. Thus, the formation of metal-free stars must have started at a very early epoch and, accordingly, very far from the Galactic center. Thus, stars inheriting heavy elements from Pop III SNe are expected to share kinematic features of the outermost halo population, such as large apo-galactic orbital radii with high eccentricities. Observations reveal that the three stars examined in the current letter do indeed possess such kinematic features \\citep{Carney_97, King_97}.  \\vspace{2cm}"
    },
    "0601/astro-ph0601168_arXiv.txt": {
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\\def\\bl{\\vskip\\normalbaselineskip}% \\def\\np{\\vfill\\eject}% \\def\\nospace{\\nulldelimiterspace=0pt \\mathsurround=0pt}% \\def\\big##1{{\\hbox{$\\left##1 \\vbox to2ex{}\\right.\\nospace$}}}% \\def\\Big##1{{\\hbox{$\\left##1 \\vbox to2.5ex{}\\right.\\nospace$}}}% \\def\\bigg##1{{\\hbox{$\\left##1 \\vbox to3ex{}\\right.\\nospace$}}}% \\def\\Bigg##1{{\\hbox{$\\left##1 \\vbox to4ex{}\\right.\\nospace$}}}% } \\autoload\\twinout{twin.txs}% \\def\\twinprint{% \\preprint \\let\\t@tl@=\\title \\def\\title{\\vskip-1.5in\\t@tl@}% \\let\\endt@tlep@ge=\\endtitlepage \\def\\endtitlepage{\\endt@tlep@ge \\twinformat}% } \\def\\twinformat{% \\tenpoint\\doublespaced \\def\\Tbf{\\twelvebf}\\def\\tbf{\\tenbf}% \\headlineoffset=0pt \\twinout}%  \\catcode`\\@=11 \\let\\NX=\\noexpand\\let\\XA=\\expandafter \\offparens \\newcount\\tabnum        \\tabnum=\\z@ \\newcount\\fignum        \\fignum=\\z@ \\newif\\ifRomanTables    \\RomanTablesfalse \\newif\\ifCaptionList    \\CaptionListfalse \\newif\\ifFigsLast       \\FigsLastfalse \\newif\\ifTabsLast       \\TabsLastfalse \\def\\FiguresLast{\\FigsLasttrue}\\def\\FiguresNow{\\FigsLastfalse} \\def\\TablesLast{\\TabsLasttrue}\\def\\TablesNow{\\TabsLastfalse} \\long\\def\\figure{\\@figure\\topinsert} \\long\\def\\topfigure{\\@figure\\topinsert}% \\long\\def\\midfigure{\\@figure\\midinsert} \\long\\def\\fullfigure{\\@figure\\pageinsert} \\long\\def\\bottomfigure{\\@figure\\bottominsert} \\long\\def\\heavyfigure{\\@figure\\heavyinsert} \\long\\def\\widefigure{\\@figure\\widetopinsert} \\long\\def\\widetopfigure{\\@figure\\widetopinsert} \\long\\def\\widefullfigure{\\@figure\\widepageinsert} \\def\\FigureName{Figure}% \\def\\TableName{Table}% \\def\\@figure#1#2{% \\vskip 0pt \\begingroup \\def\\CaptionName{\\FigureName}% \\def\\@prefix{Fg.}% \\let\\@count=\\fignum \\let\\@FigInsert=#1\\relax \\def\\@arg{#2}\\ifx\\@arg\\empty\\def\\@ID{}% \\else\\LabelParse #2;;\\endlist\\fi \\ifFigsLast \\emsg{\\CaptionName\\space\\@ID. {#2} [storing in \\jobname.fg]}% \\@fgwrite{\\@comment> \\CaptionName\\space\\@ID.\\space{#2}}% \\@fgwrite{\\string\\@FigureItem{\\CaptionName}{\\@ID}{\\NX#1}}% \\seeCR\\let\\@next=\\@copyfig \\else \\emsg{\\CaptionName\\ \\@ID.\\ {#2}}% \\let\\endfigure=\\@endfigure \\setbox\\@capbox\\vbox to 0pt{}% \\def\\@whereCap{N}% \\let\\@next=\\@findcap \\ifx\\@FigInsert\\midinsert\\goodbreak\\fi \\@FigInsert \\fi \\@next} \\def\\@endfigure{\\relax \\if B\\@whereCap\\relax \\vskip\\normalbaselineskip \\centerline{\\box\\@capbox}% \\fi \\endinsert \\endgroup}% \\def\\endfigure{\\emsg{> \\string\\endfigure before \\string\\figure!}} \\def\\figuresize#1{\\vglue #1}% \\def\\@copyfig#1#2\\endfigure{\\endgroup \\ifx#1\\par\\@fgNXwrite{#2\\@endfigure}\\else\\@fgNXwrite{#1#2\\@endfigure}\\fi} \\def\\@FGinit{\\@FileInit\\fgout=\\jobname.fg[Figures]\\gdef\\@FGinit{\\relax}} \\def\\@fgwrite#1{\\@FGinit\\immediate\\write\\fgout{#1}} \\long\\def\\@fgNXwrite#1{\\@FGinit\\unexpandedwrite\\fgout{#1}} \\def\\PrintFigures{\\ifFigsLast\\@PrintFigures\\fi} \\def\\@PrintFigures{% \\@fgwrite{\\@comment>>> EOF \\jobname.fg <<<}% \\immediate\\closeout\\fgout \\begingroup \\FigsLastfalse \\vbox to 0pt{\\hbox to 0pt{\\ \\hss}\\vss}% \\offparens \\catcode`@=11 \\emsg{[Getting figures from file \\jobname.fg]}% \\Input\\jobname.fg \\relax \\endgroup}% \\def\\@FigureItem#1#2#3{% \\begingroup #3\\relax \\def\\@ID{#2}% \\def\\CaptionName{#1}% \\setbox\\@capbox\\vbox to 0pt{}\\def\\@whereCap{N}% \\@findcap}% \\long\\def\\table{\\@table\\topinsert} \\long\\def\\toptable{\\@table\\topinsert}% \\long\\def\\midtable{\\@table\\midinsert} \\long\\def\\fulltable{\\@table\\pageinsert} \\long\\def\\bottomtable{\\@table\\bottominsert} \\long\\def\\heavytable{\\@table\\heavyinsert} \\long\\def\\widetable{\\@table\\widetopinsert} \\long\\def\\widetoptable{\\@table\\widetopinsert} \\long\\def\\widefulltable{\\@table\\widepageinsert} \\def\\@table#1#2{% \\vskip 0pt \\begingroup \\def\\CaptionName{\\TableName}% \\def\\@prefix{Tb.}% \\let\\@count=\\tabnum \\let\\@FigInsert=#1\\relax \\def\\@arg{#2}\\ifx\\@arg\\empty\\def\\@ID{}% \\else\\ifRomanTables \\global\\advance\\@count by\\@ne \\edef\\@ID{\\uppercase\\expandafter {\\romannumeral\\the\\@count}}% \\tag{\\@prefix#2}{\\@ID}% \\else \\LabelParse #2;;\\endlist\\fi \\fi \\ifTabsLast \\emsg{\\CaptionName\\space\\@ID. {#2} [storing in \\jobname.tb]}% \\@tbwrite{\\@comment> \\CaptionName\\space\\@ID.\\space{#2}}% \\@tbwrite{\\string\\@FigureItem{\\CaptionName}{\\@ID}{\\NX#1}}% \\seeCR\\let\\@next=\\@copytab \\else \\emsg{\\CaptionName\\ \\@ID.\\ {#2}}% \\let\\endtable=\\@endfigure \\setbox\\@capbox\\vbox to 0pt{}% \\def\\@whereCap{N}% \\let\\@next=\\@findcap \\ifx\\@FigInsert\\midinsert\\goodbreak\\fi \\@FigInsert \\fi \\@next} \\def\\endtable{\\emsg{> \\string\\endtable before \\string\\table!}} \\def\\@copytab#1#2\\endtable{\\endgroup \\ifx#1\\par\\@tbNXwrite{#2\\@endfigure}\\else\\@tbNXwrite{#1#2\\@endfigure}\\fi} \\def\\@TBinit{\\@FileInit\\tbout=\\jobname.tb[Tables]\\gdef\\@TBinit{\\relax}} \\def\\@tbwrite#1{\\@TBinit\\immediate\\write\\tbout{#1}} \\long\\def\\@tbNXwrite#1{\\@TBinit\\unexpandedwrite\\tbout{#1}} \\def\\PrintTables{\\ifTabsLast\\@PrintTables\\fi} \\def\\@PrintTables{% \\@tbwrite{\\@comment>>> EOF \\jobname.tb <<<}% \\immediate\\closeout\\tbout \\begingroup \\TabsLastfalse \\catcode`@=11 \\offparens \\emsg{[Getting tables from file \\jobname.tb]}% \\Input\\jobname.tb \\relax \\endgroup}% \\newbox\\@capbox \\newcount\\@caplines \\def\\CaptionName{}% \\def\\@ID{}% \\def\\captionspacing{\\normalbaselines}% \\def\\@inCaption{F}% \\long\\def\\caption#1{% \\def\\lab@l{\\@ID}% \\global\\setbox\\@capbox=\\vbox\\bgroup \\def\\@inCaption{T}% \\captionspacing\\seeCR \\dimen@=20\\parindent \\ifdim\\colwidth>\\dimen@\\narrower\\narrower\\fi \\noindent{\\bf \\linkname{\\@TagName}{\\CaptionName~\\@ID}:\\space}% #1\\relax \\vskip 0pt \\global\\@caplines=\\prevgraf \\egroup \\ifnum\\@caplines=\\@ne \\global\\setbox\\@capbox=\\vbox{\\noindent\\seeCR \\hfil{\\bf \\linkname{\\@TagName}{\\CaptionName~\\@ID}:\\space}% #1\\hfil}\\fi \\if N\\@whereCap\\def\\@whereCap{B}\\fi \\if T\\@whereCap \\centerline{\\box\\@capbox}% \\vskip\\baselineskip\\medskip \\fi}% \\def\\Caption{\\begingroup\\seeCR\\@Caption}% \\long\\def\\@Caption#1\\endCaption{\\endgroup \\ifCaptionList \\incaplist{#1}\\fi \\caption{#1}}% \\def\\endCaption{\\emsg{> \\string\\endCaption\\ called before \\string\\Caption.}} \\long\\def\\@findcap#1{% \\ifx #1\\Caption \\def\\@whereCap{T}\\fi \\ifx #1\\caption \\def\\@whereCap{T}\\fi #1}% \\def\\@whereCap{N}% \\def\\ListCaptions{\\@ListCaps\\caplist=\\jobname.cap[List of Captions] {\\let\\FIGLitem=\\CAPLitem}} \\def\\ListFigureCaptions{% \\@ListCaps\\figlist=\\jobname.fgl[List of Figure Captions] {\\let\\FIGLitem=\\CAPLitem}} \\def\\ListTableCaptions{% \\@ListCaps\\tablelist=\\jobname.tbl[List of Table Captions] {\\let\\FIGLitem=\\CAPLitem}} \\def\\CAPLitem#1#2#3\\@endFIGLitem#4{% \\bigskip \\begingroup \\raggedright\\tolerance=1700 \\hangindent=1.41\\parindent\\hangafter=1 \\noindent #1\\ #2 #3 \\hskip 0pt plus 10pt \\vskip 0pt \\endgroup}% \\def\\infiglist{\\begingroup\\seeCR \\@infiglist\\figlist} \\def\\intablelist{\\begingroup\\seeCR \\@infiglist\\tablelist} \\def\\incaplist{\\begingroup\\seeCR \\@infiglist\\caplist} \\def\\FigListWrite#1#2{% \\ifx#1\\figlist\\relax   \\FigListInit\\fi \\ifx#1\\tablelist\\relax \\TabListInit\\fi \\ifx#1\\caplist\\relax   \\CapListInit\\fi \\edef\\@line@{{#2}}% \\write#1\\@line@}% \\def\\FigListInit{\\@FileInit\\figlist=\\jobname.fgl[List of Figures]% \\gdef\\FigListInit{\\relax}} \\def\\TabListInit{\\@FileInit\\tablelist=\\jobname.tbl[List of Tables]% \\gdef\\TabListInit{\\relax}} \\def\\CapListInit{\\@FileInit\\caplist=\\jobname.cap[List of Captions]% \\gdef\\CapListInit{\\relax}} \\def\\FigListWriteNX#1#2{\\writeNX#1{#2}} \\def\\@infiglist#1#2{% \\FigListWrite#1{\\@comment > \\CaptionName\\space\\@ID:}% \\FigListWrite#1{\\string\\FIGLitem{\\CaptionName} {\\@ID.\\space}}% \\@copycap#1#2\\endlist \\FigListWrite#1{{\\NX\\folio}}% \\endgroup}% \\def\\@copycap#1#2#3\\endlist{% \\ifx#2\\space\\writeNX#1{#3\\@endFIGLitem}% \\else\\writeNX#1{#2#3\\@endFIGLitem}\\fi} \\def\\ListFigures{\\@ListCaps\\figlist=\\jobname.fgl[List of Figures]{}} \\def\\ListTables{\\@ListCaps\\tablelist=\\jobname.tbl[List of Tables]{}} \\def\\@ListCaps#1=#2[#3]#4{% \\immediate\\closeout#1 \\openin#1=#2 \\relax \\ifeof#1\\closein#1 \\else\\closein#1\\emsg{[Getting #3]}% \\begingroup \\showsectIDtrue \\ATunlock\\quoteoff\\offparens #4 \\input #2 \\relax \\endgroup \\fi} \\long\\def\\FIGLitem#1#2#3\\@endFIGLitem#4{% \\medskip \\begingroup \\raggedright\\tolerance=1700 \\ifx\\TOCmargin\\undefined\\skip0=\\parindent \\else\\skip0=\\TOCmargin\\fi \\advance\\rightskip by \\skip0 \\parfillskip=-\\skip0 \\hangindent=1.41\\parindent\\hangafter=1 \\noindent \\ifshowsectID #1\\ \\fi #2 #3 \\hskip 0pt plus 10pt \\leaddots \\hbox to 2em{\\hss\\linkto{page.#4}{#4}}% \\vskip 0pt \\endgroup} \\def\\Fig#1{\\linkto{Fg.#1}{Fig.~\\use{Fg.#1}}} \\def\\Figs#1{\\linkto{Fg.#1}{Figs.~\\use{Fg.#1}}} \\def\\Fg#1{\\linkto{Fg.#1}{\\use{Fg.#1}}} \\def\\Tab#1{\\linkto{Tb.#1}{Table~\\use{Tb.#1}}} \\def\\Tbl#1{\\linkto{Tb.#1}{Table~\\use{Tb.#1}}} \\def\\Tb#1{\\linkto{Tb.#1}{\\use{Tb.#1}}} \\autoload\\Tablebody{Tablebod.txs}\\autoload\\Tablebodyleft{Tablebod.txs} \\autoload\\tablebody{Tablebod.txs} \\autoload\\epsffile{epsf.tex}    \\autoload\\epsfbox{epsf.tex} \\autoload\\epsfxsize{epsf.tex}   \\autoload\\epsfysize{epsf.tex} \\autoload\\epsfverbosetrue{epsf.tex}\\autoload\\epsfverbosefalse{epsf.tex} \\obsolete\\topFigure\\figure \\obsolete\\midFigure\\midfigure \\obsolete\\fullFigure\\fullfigure \\obsolete\\TOPFIGURE\\figure \\obsolete\\MIDFIGURE\\midfigure \\obsolete\\FULLFIGURE\\fullfigure \\obsolete\\endFigure\\endfigure \\obsolete\\ENDFIGURE\\endfigure \\obsolete\\topTable\\toptable \\obsolete\\midTable\\midtable \\obsolete\\fullTable\\fulltable \\obsolete\\TOPTABLE\\toptable \\obsolete\\MIDTABLE\\midtable \\obsolete\\FULLTABLE\\fulltable \\obsolete\\endTable\\endtable \\obsolete\\ENDTABLE\\endtable \\def\\FIG{\\@obsolete\\FIG\\Fig\\Fig}% \\def\\TBL{\\@obsolete\\TBL\\Tbl\\Tbl}%  \\catcode`@=11 \\catcode`\\|=12 \\catcode`\\&=4 \\newcount\\ncols         \\ncols=\\z@ \\newcount\\nrows         \\nrows=\\z@ \\newcount\\curcol        \\curcol=\\z@ \\let\\currow=\\nrows \\newdimen\\thinsize      \\thinsize=0.6pt \\newdimen\\thicksize     \\thicksize=1.5pt \\newdimen\\tablewidth    \\tablewidth=-\\maxdimen \\newdimen\\parasize      \\parasize=4in \\newif\\iftableinfo      \\tableinfotrue \\newif\\ifcentertables   \\centertablestrue \\def\\centeredtables{\\centertablestrue}% \\def\\noncenteredtables{\\centertablesfalse}% \\def\\nocenteredtables{\\centertablesfalse}% \\let\\plaincr=\\cr \\let\\plainspan=\\span \\let\\plaintab=& \\def\\ampersand{\\char`\\&}% \\let\\lparen=( \\let\\NX=\\noexpand \\def\\ruledtable{\\relax \\@BeginRuledTable \\@RuledTable}% \\def\\@BeginRuledTable{% \\ncols=0\\nrows=0 \\begingroup \\offinterlineskip \\def~{\\phantom{0}}% \\def\\span{\\plainspan\\omit\\relax\\colcount\\plainspan}% \\let\\cr=\\crrule \\let\\CR=\\crthick \\let\\nr=\\crnorule \\let\\|=\\Vb \\def\\hfill{\\hskip0pt plus1fill\\hbox{}}% \\ifx\\tablestrut\\undefined\\relax \\else\\let\\tstrut=\\tablestrut\\fi \\catcode`\\|=13 \\catcode`\\&=13\\relax \\TableActive \\curcol=1 \\ifdim\\tablewidth>-\\maxdimen\\relax \\edef\\@Halign{\\NX\\halign to \\NX\\tablewidth\\NX\\bgroup\\TablePreamble}% \\tabskip=0pt plus 1fil \\else \\edef\\@Halign{\\NX\\halign\\NX\\bgroup\\TablePreamble}% \\tabskip=0pt \\fi \\ifcentertables \\ifhmode\\vskip 0pt\\fi \\line\\bgroup\\hss \\else\\hbox\\bgroup \\fi}% \\long\\def\\@RuledTable#1\\endruledtable{% \\vrule width\\thicksize \\vbox{\\@Halign \\thickrule #1\\killspace \\tstrut \\linecount \\plaincr\\thickrule \\egroup}% \\vrule width\\thicksize \\ifcentertables\\hss\\fi\\egroup \\endgroup \\global\\tablewidth=-\\maxdimen \\iftableinfo \\immediate\\write16{[Nrows=\\the\\nrows, Ncols=\\the\\ncols]}% \\fi}% \\def\\TablePreamble{% \\TableItem{####}% \\plaintab\\plaintab \\TableItem{####}% \\plaincr}% \\def\\@TableItem#1{% \\hfil\\tablespace #1\\killspace \\tablespace\\hfil }% \\def\\@tableright#1{% \\hfil\\tablespace\\relax #1\\killspace \\tablespace\\relax}% \\def\\@tableleft#1{% \\tablespace\\relax #1\\killspace \\tablespace\\hfil}% \\let\\TableItem=\\@TableItem \\def\\RightJustifyTables{\\let\\TableItem=\\@tableright}% \\def\\LeftJustifyTables{\\let\\TableItem=\\@tableleft}% \\def\\NoJustifyTables{\\let\\TableItem=\\@TableItem}% \\def\\LooseTables{\\let\\tablespace=\\quad}% \\def\\TightTables{\\let\\tablespace=\\space}% \\LooseTables \\def\\TrailingSpaces{\\let\\killspace=\\relax}% \\def\\NoTrailingSpaces{\\let\\killspace=\\unskip}% \\TrailingSpaces \\def\\setRuledStrut{% \\dimen@=\\baselineskip \\advance\\dimen@ by-\\normalbaselineskip \\ifdim\\dimen@<.5ex \\dimen@=.5ex\\fi \\setbox0=\\hbox{\\lparen}% \\dimen1=\\dimen@ \\advance\\dimen1 by \\ht0% \\dimen2=\\dimen@ \\advance\\dimen2 by \\dp0% \\def\\tstrut{\\vrule height\\dimen1 depth\\dimen2 width\\z@}% }% \\def\\tstrut{\\vrule height 3.1ex depth 1.2ex width 0pt}% \\def\\bigitem#1{% \\dimen@=\\baselineskip \\advance\\dimen@ by-\\normalbaselineskip \\ifdim\\dimen@<.5ex \\dimen@=.5ex\\fi \\setbox0=\\hbox{#1}% \\dimen1=\\dimen@ \\advance\\dimen1 by \\ht0 \\dimen2=\\dimen@ \\advance\\dimen2 by \\dp0 \\vrule height\\dimen1 depth\\dimen2 width\\z@ \\copy0}% \\def\\vctr#1{\\hfil\\vbox to 0pt{\\vss\\hbox{#1}\\vss}\\hfil}% \\def\\nextcolumn#1{% \\plaintab\\omit#1\\relax\\colcount \\plaintab}% \\def\\tab{% \\nextcolumn{\\relax}}% \\let\\novb=\\tab \\def\\vb{% \\nextcolumn{\\vrule width\\thinsize}}% \\def\\Vb{% \\nextcolumn{\\vrule width\\thicksize}}% \\def\\dbl{% \\nextcolumn{\\vrule width\\thinsize \\hskip 2\\thinsize \\vrule width\\thinsize}}% {\\catcode`\\|=13 \\let|0 \\catcode`\\&=13 \\let&0 \\gdef\\TableActive{\\let|=\\vb \\let&=\\tab}% }% \\def\\crrule{\\killspace \\tstrut \\linecount \\plaincr\\tablerule }% \\def\\crthick{\\killspace \\tstrut \\linecount \\plaincr\\thickrule }% \\def\\crnorule{\\killspace \\tstrut \\linecount \\plaincr }% \\def\\crpart{\\killspace \\linecount \\plaincr}% \\def\\tablerule{\\noalign{\\hrule height\\thinsize depth 0pt}}% \\def\\thickrule{\\noalign{\\hrule height\\thicksize depth 0pt}}% \\def\\cskip{\\omit\\relax}% \\def\\crule{\\omit\\leaders\\hrule height\\thinsize depth0pt\\hfill}% \\def\\Crule{\\omit\\leaders\\hrule height\\thicksize depth0pt\\hfill}% \\def\\linecount{% \\global\\advance\\nrows by1 \\ifnum\\ncols>0 \\ifnum\\curcol=\\ncols\\relax\\else \\immediate\\write16 {\\NX\\ruledtable warning: Ncols=\\the\\curcol\\space for Nrow=\\the\\nrows}% \\fi\\fi \\global\\ncols=\\curcol \\global\\curcol=1}% \\def\\colcount{\\relax \\global\\advance\\curcol by 1\\relax}% \\long\\def\\para#1{% \\vtop{\\hsize=\\parasize \\normalbaselines \\noindent #1\\relax \\vrule width 0pt depth 1.1ex}% }% \\def\\begintable{\\relax \\@BeginRuledTable \\@begintable}% \\long\\def\\@begintable#1\\endtable{% \\@RuledTable#1\\endruledtable}%  \\def\\E#1{\\hbox{$\\times 10^{#1}$}} \\def\\square{\\hbox{{$\\sqcup$}\\llap{$\\sqcap$}}}% \\def\\grad{\\nabla}% \\def\\del{\\partial}% \\def\\frac#1#2{{#1\\over#2}} \\def\\smallfrac#1#2{{\\scriptstyle {#1 \\over #2}}} \\def\\half{\\ifinner {\\scriptstyle {1 \\over 2}}% \\else {\\textstyle {1 \\over 2}}\\fi} \\def\\bra#1{\\langle#1\\vert}% \\def\\ket#1{\\vert#1\\/\\rangle}% \\def\\vev#1{\\langle{#1}\\rangle}% \\def\\simge{% \\mathrel{\\rlap{\\raise 0.511ex \\hbox{$>$}}{\\lower 0.511ex \\hbox{$\\sim$}}}} \\def\\simle{% \\mathrel{\\rlap{\\raise 0.511ex \\hbox{$<$}}{\\lower 0.511ex \\hbox{$\\sim$}}}} \\def\\gtsim{\\simge}% \\def\\ltsim{\\simle}% \\def\\therefore{% \\setbox0=\\hbox{$.\\kern.2em.$}\\dimen0=\\wd0 \\mathrel{\\rlap{\\raise.25ex\\hbox to\\dimen0{\\hfil$\\cdotp$\\hfil}}% \\copy0}} \\def\\|{\\ifmmode\\Vert\\else \\char`\\|\\fi} \\def\\sterling{{\\hbox{\\it\\char'44}}} \\def\\degrees{\\hbox{$^\\circ$}}% \\def\\degree{\\degrees}% \\def\\real{\\mathop{\\rm Re}\\nolimits}% \\def\\imag{\\mathop{\\rm Im}\\nolimits}% \\def\\tr{\\mathop{\\rm tr}\\nolimits}% \\def\\Tr{\\mathop{\\rm Tr}\\nolimits}% \\def\\Det{\\mathop{\\rm Det}\\nolimits}% \\def\\mod{\\mathop{\\rm mod}\\nolimits}% \\def\\wrt{\\mathop{\\rm wrt}\\nolimits}% \\def\\diag{\\mathop{\\rm diag}\\nolimits}% \\def\\TeV{{\\rm TeV}}% \\def\\GeV{{\\rm GeV}}% \\def\\MeV{{\\rm MeV}}% \\def\\keV{{\\rm keV}}% \\def\\eV{{\\rm eV}}% \\def\\Ry{{\\rm Ry}}% \\def\\mb{{\\rm mb}}% \\def\\mub{\\hbox{\\rm $\\mu$b}}% \\def\\nb{{\\rm nb}}% \\def\\pb{{\\rm pb}}% \\def\\fb{{\\rm fb}}% \\def\\cmsec{{\\rm cm^{-2}s^{-1}}}% \\def\\units#1{\\hbox{\\rm #1}} \\let\\unit=\\units \\def\\dimensions#1#2{\\hbox{$[\\hbox{\\rm #1}]^{#2}$}} \\def\\parenbar#1{{\\null\\! \\mathop{\\smash#1}\\limits ^{\\hbox{\\fiverm(--)}}% \\!\\null}}% \\def\\nunubar{\\parenbar{\\nu}} \\def\\ppbar{\\parenbar{p}} \\def\\buildchar#1#2#3{{\\null\\! \\mathop{\\vphantom{#1}\\smash#1}\\limits ^{#2}_{#3}% \\!\\null}}% \\def\\overcirc#1{\\buildchar{#1}{\\circ}{}} \\def\\sun{{\\hbox{$\\odot$}}}\\def\\earth{{\\hbox{$\\oplus$}}} \\def\\slashchar#1{\\setbox0=\\hbox{$#1$}% \\dimen0=\\wd0 \\setbox1=\\hbox{/} \\dimen1=\\wd1 \\ifdim\\dimen0>\\dimen1 \\rlap{\\hbox to \\dimen0{\\hfil/\\hfil}}% #1 \\else \\rlap{\\hbox to \\dimen1{\\hfil$#1$\\hfil}}% / \\fi}% \\def\\subrightarrow#1{% \\setbox0=\\hbox{% $\\displaystyle\\mathop{}% \\limits_{#1}$}% \\dimen0=\\wd0 \\advance \\dimen0 by .5em \\mathrel{% \\mathop{\\hbox to \\dimen0{\\rightarrowfill}}% \\limits_{#1}}}% \\newdimen\\vbigd@men \\def\\vbigl{\\mathopen\\vbig} \\def\\vbigm{\\mathrel\\vbig} \\def\\vbigr{\\mathclose\\vbig} \\def\\vbig#1#2{{\\vbigd@men=#2\\divide\\vbigd@men by 2 \\hbox{$\\left#1\\vbox to \\vbigd@men{}\\right.\\n@space$}}} \\def\\Leftcases#1{\\smash{\\vbigl\\{{#1}}} \\def\\Rightcases#1{\\smash{\\vbigr\\}{#1}}} \\def\\doublecolumns{\\relax}\\def\\enddoublecolumns{\\relax} \\def\\leftcolrule{\\relax}\\def\\rightcolrule{\\relax} \\def\\longequation{\\relax}\\def\\endlongequation{\\relax} \\def\\newcolumn{\\relax} \\def\\widetopinsert{\\topinsert}\\def\\widepageinsert{\\pageinsert} \\def\\forceleft{\\relax}\\def\\forceright{\\relax} \\def\\SetDoubleColumns#1{% \\imsg{The double column macros are not a part of mTeXsis.} \\imsg{If you want to use double column mode, get TXSdcol.tex} \\imsg{and add \\string\\input\\space TXSdcol.tex to your .tex file.} }   \\def\\addTOC#1#2#3{\\relax}\\def\\Contents{\\relax}  % \\newif\\ifContents                               % \\def\\ContentsSwitchtrue{\\Contentstrue}\\def\\ContentsSwitchfalse{\\Contentsfalse}  \\def\\obsolete#1#2{\\let#1=#2\\relax #2}           %  \\let\\Input=\\input                               % \\newdimen\\colwidth      \\colwidth=\\hsize        % \\def\\ORGANIZATION{}%   \\newhelp\\@utohelp{% loadstyle: The definition of the macro named above is actually contained^^J% in a style file, and so it cannot be used with mTeXsis.  If you really^^J% need to load the definition from that file, you should do so explicitly^^J% at the begining of your manuscript file, with % '\\string\\input\\space stylefilename.txs'^^J}  \\Ignore \\def\\loadstyle#1#2{% \\newlinechar=10                              % \\errhelp=\\@utohelp                           % \\emsg{> Whoops! Trying to load \\string#1\\space from style file #2.}% \\errmessage{You cannot use macro definitions from style files in mTeXsis}} \\endIgnore   \\hbadness=10000         % \\overfullrule=0pt       % \\vbadness=10000         %   \\ATunlock \\SetDate                                % \\ReadAUX                                % \\def\\fmtname{TeXsis}\\def\\fmtversion{2.17}% \\def\\revdate{1 January 1998}% \\def\\imsg#1{\\emsg{\\@comment #1}}% \\imsg{=========================================================== \\@comment} \\imsg{This is mTeXsis, the core macros from TeXsis.} \\imsg{You can get the complete TeXsis package (and avoid this annoying} \\imsg{advertisement) from ftp://lifshitz.ph.utexas.edu/texsis, } \\imsg{or from a CTAN server near you (in macros/texsis).} \\imsg{See the README and INSTALL files there for more information.} \\imsg{============================================================ \\@comment} \\emsg{m\\fmtname\\space version \\fmtversion\\space (\\revdate)  loaded.}% \\ATlock                                 % \\texsis                                 % \\quoteoff \\global\\mathchardef\\DELTA=\"7001 \\global\\mathchardef\\LAMBDA=\"7003 \\global\\mathchardef\\XI=\"7004 \\global\\mathchardef\\SIGMA=\"7006 \\global\\mathchardef\\UPSILON=\"7007 \\global\\mathchardef\\OMEGA=\"700A \\global\\mathchardef\\Delta=\"7101 \\global\\mathchardef\\Lambda=\"7103 \\global\\mathchardef\\Xi=\"7104 \\global\\mathchardef\\Sigma=\"7106 \\global\\mathchardef\\Upsilon=\"7107 \\global\\mathchardef\\Omega=\"710A   \\catcode`\\@=11  \\font\\tenmsa=msam10 \\font\\sevenmsa=msam7 \\font\\fivemsa=msam5 \\font\\tenmsb=msbm10 \\font\\sevenmsb=msbm7 \\font\\fivemsb=msbm5 \\newfam\\msafam \\newfam\\msbfam \\textfont\\msafam=\\tenmsa  \\scriptfont\\msafam=\\sevenmsa \\scriptscriptfont\\msafam=\\fivemsa \\textfont\\msbfam=\\tenmsb  \\scriptfont\\msbfam=\\sevenmsb \\scriptscriptfont\\msbfam=\\fivemsb  \\def\\hexnumber@#1{\\ifnum#1<10 \\number#1\\else \\ifnum#1=10 A\\else\\ifnum#1=11 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\\def\\refFormat{\\catcode`\\^^M=10} \\def\\Refs#1#2{Refs.~\\use{Ref.#1},\\use{Ref.#2}} \\def\\Refsand#1#2{Refs.~\\use{Ref.#1} and~\\use{Ref.#2}} \\def\\Refsrange#1#2{Refs.~\\use{Ref.#1}--\\use{Ref.#2}} \\superrefsfalse % \\def\\Chapter#1{Chapter~\\use{Chap.#1}} \\def ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601442_arXiv.txt": {
        "abstract": "A new technique is presented for producing images from interferometric data. The method, ``smear fitting'', makes the constraints necessary for interferometric imaging double as a model, with uncertainties, of the sky brightness distribution.  It does this by modelling the sky with a set of functions and then convolving each component with its own elliptical gaussian to account for the uncertainty in its shape and location that arises from noise.  This yields much sharper resolution than CLEAN for significantly detected features, without sacrificing any sensitivity.  Using appropriate functional forms for the components provides both a scientifically interesting model and imaging constraints that tend to be better than those used by traditional deconvolution methods.  This allows it to avoid the most serious problems that limit the imaging quality of those methods. Comparisons of smear fitting to CLEAN and maximum entropy are given, using both real and simulated observations.  It is also shown that the famous Rayleigh criterion (resolution = wavelength / baseline) is inappropriate for interferometers as it does not consider the reliability of the measurements. ",
        "introduction": "Interferometers give us a much sharper view than filled aperture telescopes of the same area, by sampling the sky's spatial frequencies with a set of baselines given by the separations between each receiver.  Unfortunately the distribution of samples (``visibilities'') is incomplete, so the Fourier transform of the measured visibilities is \\emph{not} the sky brightness distribution.  The Fourier transform of the visibilities instead yields the ``dirty map'', which is the sky brightness distribution plus noise, convolved with the dirty beam (the Fourier transform of the sampling pattern).  In other terms, a simple Fourier transform produces an image with absolutely no power at unsampled spatial frequencies.  As a result the dirty beam exhibits sidelobes causing the fainter structure in the image to be buried under the diffraction patterns of the brightest objects in the field.  Unfortunately the most straightforward way of deconvolving away the effect of the dirty beam would divide by zero in the $uv$ plane (the Fourier transform of the image plane) wherever no measurement was made, and practical methods must instead attempt to fill unmeasured regions of the $uv$ plane with a more reasonable estimate than zero.  Simply interpolating between the visibilities does not work in general because aliasing of oscillations in the $uv$ plane can erroneously move the emission toward the center of the image plane, and distort its appearance.  Thus there is no unique prescription for extracting the optimum estimate of the true sky brightness, and many methods, most notably CLEAN \\citep{bib:hogbom74} and maximum entropy deconvolution (\\citealt{bib:gullskilling}, \\citealt{bib:cbb}), that vary in their properties have been devised.  They are usually, if not strictly correctly, called deconvolution methods.  They rely on expected or desired properties of the images such as positivity, locality, or smoothness to constrain the images they produce.  This paper presents a new method, called smear fitting, which uses as simple a model as it can as its main constraint, and can optionally use additional constraints such as positivity and/or confinement of the source to a certain region.  It renders the measurements into images with better fidelity and resolution, and fewer image artifacts, than traditional deconvolution methods.  A great benefit of smear fitting is considerably improved resolution for objects with peak brightness greater than 4.20 times the root mean square (rms) noise (it does not change the resolution of fainter features) without a loss of sensitivity from reweighting the data.  This feature is extremely important since better resolution cannot be achieved by simply adding data, and the noise that comes with it, from longer baselines or decreasing the weight of short baselines, without increasing the root mean square (rms) errors in surface brightness over the whole image.  The objects of interest in typical radio astronomical observations span wide ranges of both brightness and size, and smear fitting offers a way to optimally handle both.  Smear fitting also produces a set of components modelling the source, and calculates the uncertainties of the distribution of those components on the sky.  The components often correspond to distinct physical features of the source(s), making their parameters and associated uncertainties of immediate scientific interest.  Smear fitting has been implemented as a modification (patch) to difmap \\citep{bib:difmap1997}.  The patch, known as smerf, is freely available at http:$/\\!$/www.drao-ofr.hia-iha.nrc-cnrc.gc.ca/$\\!\\sim$rreid/smerf/. ",
        "conclusions": "\\label{sec:conclusion}  Smear fitting is an image deconvolution method that fits a near maximally simple model to interferometer visibilities and then broadens the model to account for the uncertainty of those visibilities.  The model construction method avoids several problems that can limit the quality of CLEAN and maximum entropy deconvolutions, and smearing generally yields sharper and fairer images than CLEAN.  As mentioned above, smear fitting has been implemented as a modification (patch) to difmap, a well known program for imaging and selfcalibrating data from radio interferometers.  The patch, known as smerf, is freely available under the GNU license \\citep{bib:gnu} at http:$/\\!$/www.drao-ofr.hia-iha.nrc-cnrc.gc.ca/$\\!\\sim$rreid/smerf/, and includes a manual on its use."
    },
    "0601/astro-ph0601397_arXiv.txt": {
        "abstract": "The {\\bf A}ntarctic {\\bf M}uon {\\bf A}nd {\\bf N}eutrino {\\bf D}etector {\\bf A}rray (\\aman) is a high-energy neutrino telescope. It is a lattice of optical modules (OM) installed in the clear ice below the South Pole Station. Each OM contains a photomultiplier tube (PMT) that detects photons of \\cer light generated in the ice by muons and electrons. \\ice is a cubic-kilometer-sized expansion of \\aman currently being built at the South Pole. In \\ice the PMT signals are digitized already in the optical modules and transmitted to the surface. A prototype string of 41 OMs equipped with this new all-digital technology was deployed in the \\aman array in the year 2000. In this paper we describe the technology and demonstrate that this string serves as a proof of concept for the \\ice array. Our investigations show that the OM timing accuracy is 5~ns. Atmospheric muons are detected in excellent agreement with expectations with respect to both angular distribution and absolute rate. ",
        "introduction": "\\label{sec:intro} The {\\bf A}ntarctic {\\bf M}uon {\\bf A}nd {\\bf N}eutrino {\\bf D}etector {\\bf A}rray (\\aman)~\\cite{Amanda:2001} at the geographic South Pole is a lattice of photomultiplier tubes (PMTs) each enclosed in a transparent pressure sphere to comprise an optical module (OM). The OMs are buried in the polar ice at depths between 1500\\,m and 2000\\,m. The primary goal of \\aman is to discover astrophysical sources of high energy neutrinos. High-energy muon neutrinos penetrating the earth from the Northern Hemisphere are identified by the secondary muons produced in charged current neutrino-nucleon interactions near or within the detector. \\aman was the first neutrino telescope with an effective area in excess of 10000~m$^2$ -- with over 600 optical modules in place.   Large as it is, \\aman is only capable of seeing the brighter (and closer) sources of neutrinos. \\ice will be an array of 4,800 optical modules within a cubic kilometer of ice~\\cite{PDD:2001}. Frozen into vertical holes 2.4~km deep, drilled by hot water, the optical modules will lie between 1450~m and 2450~m below the surface. An instrument of this size is suited to study neutrinos from distant astrophysical sources~\\cite{Ahrens:2003ix}.  The optical modules in \\aman and \\ice are arranged like strings of beads connected electrically and mechanically to long cables leading up to a control room at the surface. The PMT within the OM  detects individual photons of \\cer light generated in the ice by muons and electrons moving with velocities near the speed of light. Signal events consist primarily of up-going muons produced in neutrino interactions in the bedrock of Antarctica or in the ice. In addition, the detector will identify electromagnetic and hadronic showers (``cascades'') from interactions of  $\\nu_e$ and $\\nu_\\tau$ inside the detector volume provided they are sufficiently energetic. Background events are mainly downward-going muons from cosmic ray interactions in the atmosphere above the detector. For \\ice, the background will be monitored for calibration purposes by the \\textsc{IceTop} air shower array~\\cite{PDD:2001} covering a 1~km$^2$ area at the ice surface right above the in-ice OMs of the \\ice detector. This background will also be used as a test beam for commissioning the detector.  Unlike \\aman, \\ice PMT signals are digitized directly in the optical module itself and then transmitted digitally to the surface. This is done such that the arrival time of a photon at an OM can be determined to within a few nanoseconds. The electronics at the surface determines when an event has occurred (e.g., that a muon traversed or passed near the array) and records the information for subsequent event reconstruction and analysis.  In the following we will discuss the performance obtained with a string of 41 Digital Optical Module (DOM) prototypes 12~m apart deployed in January 2000 by the \\aman Collaboration as the 18$^\\text{th}$ string (\\st) of the \\aman array. ",
        "conclusions": "\\st is established as proof of concept for the \\ice array. The all-digital technology with pulse digitization already at the PMT inside the optical module, digital data transmission to the surface and automated time calibration was shown to be fully functional including the benchmark observation of atmospheric muons.   \\begin{ack} This research was supported by the following agencies: Deutsche Forschungsgemeinschaft (DFG); German Ministry for Education and Research; Knut and Alice Wallenberg Foundation, Sweden; Swedish Research Council; Swedish Natural Science Research Council; Fund for Scientific Research (FNRS-FWO), Flanders Institute to encourage scientific and technological research in industry (IWT), and Belgian Federal Office for Scientific, Technical and Cultural affairs (OSTC), Belgium. UC-Irvine AENEAS Supercomputer Facility; University of Wisconsin Alumni Research Foundation; U.S. National Science Foundation, Office of Polar Programs; U.S. National Science Foundation, Physics Division; U.S. Department of Energy; D.F. Cowen acknowledges the support of the NSF CAREER program. I. Taboada acknowledges the  support of FVPI. M.~Ribordy acknowledges the support of the Swiss National Science Foundation. K.~Helbing acknowledges the  support of the Alexander von Humboldt Foundation. \\end{ack}"
    },
    "0601/astro-ph0601674_arXiv.txt": {
        "abstract": "We present 2-degree field spectroscopic observations of a sample of 96 A--type stars selected from the Sloan Digital Sky Survey Data Release 3 (SDSS DR3). Our aim is to identify blue horizontal branch (BHB) stars in order to measure the kinematic properties of the tidal tails of the Sagittarius dwarf spheroidal galaxy. We confine our attention to the 44 classifiable stars with spectra of signal-to-noise ratio $>15$\\AA$^{-1}$.  Classification produces a sample of 29 BHB stars at distances $5-47\\,$kpc from the Sun. We split our sample into three bins based on their distance. We find 10 of the 12 stars at $14-25\\,$kpc appear to have coherent, smoothly varying radial velocities which are plausibly associated with old debris in the Sagittarius tidal stream. Further observations along the orbit and at greater distances are required to trace the full extent of this structure on the sky. Three of our BHB stars in the direction of the globular cluster Palomar (Pal) 5 appear to be in an overdensity but are in the foreground of Pal 5. More observations are required around this overdensity to establish any relation to Pal 5 and/or the Sgr stream.  We emphasize observations of BHB stars have unlimited potential for providing accurate velocity and distance information in old distant halo streams and globular clusters alike. The next generation multi-object spectrographs provide an excellent opportunity to accurately trace the full extent of such structures. ",
        "introduction": "There have been numerous searches for streams of material responsible for the build-up of stellar halos in galaxies. In the Milky Way the best studied is the Sagittarius (Sgr) dwarf spheroidal galaxy (Ibata, Gilmore \\& Irwin 1994) and its stellar stream (e.g. Majewski et al. 2003). An extended stream of stars has also been uncovered in the halo of the Andromeda galaxy (M31), revealing that it too is cannibalizing a small companion (e.g. Lewis et al. 2004). Such streams yield crucial information on the merging and accretion history of galaxy halos and have been used to constrain the mass of the Milky Way halo (e.g.  Johnston et al.  1999) and the mass of the halo in M31 (Ibata et al. 2004).  There are numerous studies reporting stellar streams associated with the tidal debris of the Sgr system either trailing or leading it along its orbit. The stars involved in the tidal disruptions can be used as test particles which can lead to an understanding of the shape (Helmi 2004) and strength of the Milky Way potential. A number of studies have sought to trace Sgr and its stream with a variety of stellar populations. For example, Totten \\& Irwin (1998) found evidence for such streams using intrinsically rare, but very luminous, carbon stars. The calibration of the carbon star distance scale remains controversial and is further complicated by variability, dust and metallicity effects. Numerous RR Lyrae stars have also been found to be associated with the Sgr stream (Ivezi{\\'c} et al. 2000; Vivas et al. 2001; Vivas et al. 2005). The Quasar Equatorial Survey Team (QUEST) survey (Vivas at al. 2004) found 85 RR Lyrae variables covering about 36$^{\\circ}$ in right ascension, of which 16 were identified spectroscopically (Vivas et al. 2005). An all-sky view of the Sgr stream has been investigated using M giants selected from the Two Micron All-Sky Survey (2MASS) database (Majewski et al. 2003, 2004).  There is strength in diversity. Heterogeneous stellar populations not only provide an important check on the distance and velocity measurements of any putative stream but also probe different formation epochs. Blue Horizontal Branch (BHB) stars provide yet another stellar population ideal for exploring Halo structure (Sirko et al. 2004; Brown et al., 2004; Clewley et al., 2005). Newberg at al. (2002) find a distant sample of BHB stars that, they suggest, reside in the trailing arm of the Sgr stream.  Further, Monaco et al. (2003) suggest that a sample of BHB stars discovered in the core of the Sgr dwarf spheroidal galaxy are the counterpart of the stars observed by Newberg et al. (2002).  Both BHB stars and RR Lyraes are excellent standard candles, enabling us to determine their distances to 5-10\\%; by contrast the distances to K and M giants are accurate to 20\\% - 25\\% (e.g. Dohm-Palmer et al. 2001; Majewski et al. 2003). BHB and RR Lyrae stars are also old and metal poor, so they are ideal for tracing old parts ($>$ 5 Gyr) of tidal streams. The age of the Sgr M giants is controversial.  Majewski et al. 2003 considered these stars to be younger than 5 Gyr with a significant fraction of them only 2$-$3$\\,$Gyrs old (Majewski et al. 2003). This poses two problems for their use as tracers. First, they may not be useful halo test particles as such dynamically young tracers do not place very stringent constraints on the halo (Helmi 2004). Second, comparisons between the evolution timescales of the M giants and the dynamical timescale required for the stars to populate the stream suggest that the stars exist in a stream that took more than their age to form. This inconsistency has been addressed by Bellazzini et al. (2005) who finds the M giants were considerably older ($>$ 5 Gyr) than was originally thought and may even probe similar epochs as RR Lyraes or BHB stars.  This paper is the first in a series that discusses the use of BHB stars to directly trace the Sgr stream. BHB stars are A--type giants that are easily identified in the Galactic halo as they lie blueward of the main--sequence turnoff (e.g.  Yanny et al. 2000). Assembling clean samples of remote BHB stars has been stymied by the existence of a contaminating population of high--surface--gravity A--type stars, the blue stragglers, that are around two magnitudes fainter.  Previous analyses required high signal-to-noise ratio ($S/N$) spectroscopy to reliably separate these populations (e.g.  Kinman, Suntzeff \\& Kraft 1994), making identification of BHB stars in the distant halo unfeasible. However, recently Clewley et al. (2002, 2004) developed two classification methods that now enable us to overcome the difficulties of cleanly separating BHB stars from blue stragglers.  In this paper we use these classification techniques, and the Sloan Digital Sky Survey Data Release 3 (SDSS DR3) photometry, to isolate a sample of BHB stars along the Sgr stream. In this new survey we target stars along the equatorial strip in the region $180^{\\circ}<$RA$<249^{\\circ}$, which are sensitive to the different models that distinguish between halo flatness using the Sgr stream (e.g. Mart{\\'{\\i}}nez-Delgado et al. 2004).  In Section 2 of this paper we describe the selection of the BHB candidates from the SDSS data set, and outline the prescription for transforming the SDSS $g,r$ magnitudes of A--type stars to $B,V$ magnitudes detailed in Clewley et al. (2005). Section 3 provides a summary of the 2dF spectroscopic observations and data reduction procedures. In Section 4 we classify these stars into categories BHB and blue straggler. Section 5 presents the results of the classification procedure and provides a summary table of distances and radial velocities of the stars classified as BHB. Three BHB stars in the direction of the globular cluster Pal 5 appear to be in an overdensity but in the foreground of Pal 5; we compare our measurements with those in the literature. In Section 6 we investigate the evidence that the stars reside in a stream and are associated with the Sgr tidal debris. Finally, in Section 7 we provide a summary of the main conclusions of the paper. In this paper we use the coordinate R$_{gal}$ to denote Galactocentric distances and the coordinate R$_{\\odot}$ to denote heliocentric distances. ",
        "conclusions": "Table \\ref{data_sum} summarises the main observational results of the paper; a sample of BHB stars with measured radial velocities in the vicinity of the Sgr stream. We now investigate whether there is evidence that these stars actually reside in this stream.  In Figure \\ref{comp_delgado} we plot Heliocentric radial velocity and distance versus RA for the 29 BHB stars. In Figure \\ref{comp_delgado}(upper) we split the sample into three groups of R$_{\\odot}$: stars at 5 $<$ R$_{\\odot}$ $<$ 14 kpc (open circles); (ii) stars at 14 $<$ R$_{\\odot}$ $<$ 25 kpc (filled circles); and stars at R$_{\\odot}$ $>$ 25 (open triangles). We overplot on this figure simulations of the tidal stream of Sgr for various halo flatness taken from Mart{\\'{\\i}}nez-Delgado et al. (2004). The curves on this plot are for values of the axis ratio of the density distribution, $q_d$, which range from 0.1 to 1.0 (upper to lower curves). However, the axis ratio of the potential, $q_p$, is the quantity that determines the satellite orbit and hence the shape of the potential. This parameter is a function of $q_d$ and Galactocentric distance (see Fig. 11 in Mart{\\'{\\i}}nez-Delgado et al. 2004) and ranges from 0.5 (oblate potentials) to 1.0 (spherical potentials).  In Figure \\ref{comp_delgado}(lower) we plot V$_{\\odot}$ versus RA for the same three groups of BHB stars. The plot reveals a possible correlation between V$_{\\odot}$ and RA which suggests that 10 of the 12 stars in the region 14 $<$ R$_{\\odot}$ $<$ 25 kpc could be associated with a stream. It is clear, however, that the velocity distribution of the stars is systematically lower than any of the Mart{\\'{\\i}}nez-Delgado et al. (2004) models presented here. There could be many reasons for this. One plausible reason is that the Mart{\\'{\\i}}nez-Delgado model considers only the last ($\\sim 5\\,$Gyr) orbit. As BHB stars are old, any putative stream might feasibly belong to a later orbit that is not considered in this model. In the next section we investigate the evidence that this putative stream is plausibly associated with the Sgr tidal debris. \\begin{figure*} \\centering{\\scalebox{0.85}{ \\includegraphics*[40,40][750,750]{fig_6_lowres.ps}} }\\caption{Distances and radial velocities of the best-fit simulations of the Galactic halo potential created by Law, Johnston \\& Majewski (2005) for oblate potentials (top plots), spherical potentials (middle plots) and prolate potentials (bottom plots). The bold dots (green) are old debris stripped four orbits ago whilst faint dots (red) are from earlier orbits. Overplotted are: BHB observations (as in Fig. \\ref{comp_delgado}, i.e. filled and open circles and filled triangles); filled squares are carbon stars selected from Totten \\& Irwin (1998); open stars are M giants from Majewski et al. (2004) and crosses (magenta) are K-giants from Kundu et al. (2002).\\label{law05plot}} \\end{figure*}  \\subsection{Comparisons with previous work} As we briefly discuss in \\S1 there have been numerous detections of the Sgr tidal tail around the sky using tracers that map different epochs, i.e orbits in which the debris was stripped from the satellite. For example, Majewski et al. (2003, 2004) isolated M giants in the Two Micron All-Sky Survey (2MASS) in order to map the position and velocity distribution of tidal debris from Sgr around the entire Galaxy. In Law, Johnston \\& Majewski (2005; hereafter LJM05) these stars are compared with numerous simulations of the tidal debris of the Sgr satellite in differing Galactic potentials. We refer the reader to this paper for details of the parameters used in this model. The simulations are created using the relatively young M giants as the principal observable. We note that BHB stars and carbon stars are expected to trace older debris than M giants. We compare these models with our BHB stars and other stellar types that are available.  Before we are able to compare our data with such models we need to convert our coordinates to the system defined in Majewski et al. (2003). In this coordinate system the zero plane of the latitude coordinate $B_{\\odot}$ coincides with the best-fit great circle defined by the Sgr debris, as seen from the Sun; the longitudinal coordinate $\\Lambda_{\\odot}$ is zero in the direction of the Sgr core and increases along the Sgr trailing stream. Our sample resides in the region $260^\\circ< \\Lambda_{\\odot} <320^\\circ$ and $-17^\\circ < B_{\\odot} < 16^\\circ$. This corresponds to a section of the leading portion of the orbit which is particularly sensitive to the halo shape. Indeed, M stars appear to strongly favour models with prolate rather than oblate halos (Helmi 2004). In contrast, by fitting planes to the leading and trailing debris and measuring the orbital precession of the Sgr debris Johnston, Law \\& Majewski (2005) find oblate halos are favoured.  In Figure \\ref{law05plot} we compare our BHB data with the best-fit numerical simulations. Overplotted are various other tracers which we discuss below. The M giants are represented by open stars. Figure \\ref{law05plot} illustrates the differences for the halo potential models. We plot distances and radial velocities of the best-fit simulations of the Galactic halo potential, which was created by LJM05 for oblate (q=0.9, upper plots), spherical (q=1.0, middle plots) and prolate (q=1.25, lower plots) potentials. The bold points (green) are old debris stripped four orbits ago whilst faint points (red) are from earlier orbits.  The BHB stars are shown as filled circles, open circles and filled triangles as in Figure \\ref{comp_delgado}. The filled circles represent our putative stream. The BHB stars that are shown as open circles do not appear to be associated with the models in either velocity or distance. However, the 10 of 12 filled circles are plausibly associated with the older debris both in velocity and distance in the LJM05 model. More data is required to confirm whether the three distant BHB stars (shown as triangles) are also associated with a more distant younger part of the Sagittarius stream. They do not show a remarkable metallicity which might suggest they are in an earlier orbit of Sgr.  Carbon stars (filled squares) are from the catalogue provided by Totten \\& Irwin (1998), confining ourselves to the same $\\Lambda_{\\odot}$ ranges as the BHB data and to similar distances (improved distance estimates are added from Totten, Irwin \\& Whitelock 2000). Some of the carbon stars (blue squares) do appear to follow the M stars in velocity, however their distances do not appear to be associated at all. This discrepancy might be due to the uncertainty in distance measurements for individual carbon stars. It is also plausible that the carbon stars trace older debris than the M giants as they can have larger ages. This sample however does not show any correlation with the BHB stars.  The crosses in Figure \\ref{law05plot} are eight metal poor K-giants discovered by Kundu et al. (2002) to have coherent, smoothly varying distances and radial velocities. The Kundu et al. (2002) K-giant stars represent relatively old debris stripped from Sgr three pericentric passages ago. The velocities of this sample fit the models reasonably well particularly for the prolate potential shown in Figure \\ref{law05plot}(lower). However, the distances remain uncertain in all three models.  In summary, the M stars and carbon stars predominantly map out the earlier tidal debris. The velocity data for both tracers appears to qualitatively fit the prolate models rather well (Figure \\ref{law05plot}, lower). However the distance information is less clear, particularly for the carbon stars. The velocity information of the metal poor K-giants again appears to fit the prolate models well but again the distance information does not support this case. The BHB data are too scarce in this study to distinguish competing models of halo flatness despite being comparitively accurate standard candles. We emphasise that these observations are over a section of the leading part of the stream. Conclusive evidence for halo flatness will probably require multi-epoch observational data from leading and trailing debris."
    },
    "0601/astro-ph0601504_arXiv.txt": {
        "abstract": "{We construct a data base of  125 post-AGB objects (including R CrB and extreme helium stars) with published photospheric parameters (effective temperature and gravity) and chemical composition. We estimate the masses of the post-AGB stars by comparing their position in the (log \\Teff, log $g$) plane  with theoretical evolutionary tracks of different masses. We construct various diagrams, with the aim of finding clues to AGB nucleosynthesis. This is the first time that a large sample of post-AGB stars has been used in a systematic way for such a purpose and we argue that, in several respects, post-AGB stars should be more powerful than planetary nebulae to test AGB nucleosynthesis.  Our main findings are that:    the vast majority of objects which do not show evidence of N production from primary C   have  a low stellar mass ($M_{\\star}$ $<$ 0.56 M$_{\\odot}$);  there is no evidence that objects which did not experience 3rd dredge-up have a different stellar mass distribution than objects that did; there is clear evidence that 3rd dredge-up is more efficient at low metallicity. The sample of known post-AGB stars is likely to increase significantly in the near future thanks to the   ASTRO-F and follow-up observations, making these objects even more promising as testbeds for AGB nucleosynthesis. ",
        "introduction": "The asymptotic giant branch (AGB) phase, a late stage in the evolution of low and intermediate mass stars, is quite complex. During this phase, various nuclear processes are at work in different zones of the star, and a variety of mixing mechanisms take place (see e.g. Charbonnel\\,\\cite{Charbonnel_2002} or Lattanzio\\,\\cite{Lattanzio_2002} for short reviews). Understanding this phase is important, since the elements manufactured during the AGB contribute significantly to the chemical composition of galaxies.  The first models to give some predictions on the stellar yields from AGB stars are those by Iben \\& Truran\\,(\\cite{Iben_and_Truran_1978}) and Renzini \\& Voli\\,(\\cite{Renzini_and_Voli_1981}). Unfortunately, even nowadays, the state of stellar physics does not allow one to construct models entirely from first principles, and some quantities have to be set as free parameters. The next generation of AGB models (Groenewegen \\& de Jong\\,\\cite{Groenewegen_and_de_Jong_1993}, Marigo et al.\\,\\cite{Marigo_etal_1996}) adjusted those free parameters (essentially mass loss rate and mixing length) to reproduce a few observational constraints on stellar populations. Further models have been constructed since then (e.g. by Forestini \\& Charbonnel\\,\\cite{Forestini_and_Charbonnel_1997}, Boothroyd \\& Sackmann\\,\\cite{Boothroyd_and_Sackmann_1999}, Marigo\\,\\cite{Marigo_2001}, Izzard et al.\\,\\cite{Izzard_etal_2004}). All those models are so-called synthetic models in that, for the thermal pulse phase, they use analytical expressions to extrapolate certain quantities obtained from full evolutionary calculations up to the planetary nebula ejection. With the availability of fast computers, it is now possible to compute complete AGB models (Karakas et al.\\,\\cite{Karakas_etal_2002}, Herwig\\,\\cite{Herwig_2004}) that follow all the pulses in detail. Testing model predictions before using them in chemical evolution models of galaxies is vital, especially because even full evolutionary calculations are computationally difficult and still imply some  ad hoc parameters (mixing length, mass loss rates etc.). Indeed, as shown recently by Ventura \\& D'Antona (2005a, 2005b), the predicted yields of intermediate mass stars depend strongly on the treatment adopted for convection and mass loss and on the nuclear reaction cross-sections.  As mentioned before, synthetic models published  since 1993 reproduce some observables (e.g. the luminosity functions of carbon-stars and of lithium-rich stars) by construction. However, additional tests are needed and are crucial to constrain AGB models. The analysis of  the chemical composition of planetary nebulae can provide some tests. Such an approach has been adopted recently by Marigo et al.\\,(\\cite{Marigo_etal_2003}) using a data base of 10 planetary nebulae (PNe) observed in a wide spectral range, from the far infrared to the ultraviolet. Previously, the chemical composition of planetary nebulae samples had been used in a more empirical way to test AGB nucleosynthesis and get an insight into the relation between planetary nebulae, the final products of  low and intermediate mass stars and their progenitors. For example, Peimbert\\,(\\cite{Peimbert_1978}) defined a category of planetary nebulae, called Type I PNe, as objects having He/H $>$ 0.125 and N/O $>$ 0.5. Those PNe were interpreted as being born from stars that experienced the second dredge-up.  While the designation \"Type I PNe\" stayed in the literature, its definition changed several times (see Stasi\\'nska\\,\\cite{Stasinska_2004}). Kingsburgh \\& Barlow\\,(\\cite{Kingsburgh_and_Barlow_1994}) compared the nebular N/H ratio to the solar (C+N)/H ratio to find out which planetary nebulae exhibit N produced from primary C.  They also identified planetary nebulae whose progenitors had experienced 3rd dredge-up, bringing to the surface carbon produced by the triple-$\\alpha$ reaction. To this end, they constructed a (C+N+O)/H vs C/H diagram. The N/O vs O/H diagram has been interpreted (e.g. Henry\\,\\cite{Henry_1990}) as indicating that in some objects N is produced at the expense of O (during ON cycling). However, this diagram, depending on authors and samples, does not always lead to such a straightforward interpretation (Kingsburgh \\& Barlow\\,\\cite{Kingsburgh_and_Barlow_1994}, Leisy \\& Dennefeld\\,\\cite{Leisy_and_Dennefeld_1996}). Henry et al.\\,(\\cite{Henry_etal_2000}) have plotted PN abundances of C and N as a function of progenitor mass (estimated from the stellar remnant masses by G\\'orny et al.\\,\\cite{Gorny_etal_1997} and Stasi\\'nska et al.\\,\\cite{Stasinska_etal_1997} and converted to progenitor masses using an initial-final mass relation), and compared this with the predictions from synthetic AGB models. This was the first attempt to compare PNe abundances directly with the results of AGB model computations for a given initial mass and it was not very conclusive.  There is another category of objects that may serve for tests of AGB models. These are post-AGB stars, i.e. stars that have already ejected their envelope but are not yet hot enough to ionize it and produce a planetary nebula.  Such stars have several advantages with respect to PNe and constitute an excellent complement for testing AGB nucleosynthesis. The main advantage is that the stellar mass can be obtained directly from the observed stellar spectrum by fitting a model atmosphere that allows one to derive the stellar effective temperature, \\Teff, and gravity $g$ (this can also be done for PNe nuclei, but  as the stars hotter and buried in the ionized gas, this is more difficult). The abundances of quite a variety of elements (He, Li, C, N, O, Na, Mg, Al, Si, S, Ca, Sc, Ti, Fe, Ni, Zn ...) can be obtained from a stellar atmosphere analysis. For carbon, a particularly important element in AGB evolution, the uncertainty in the estimated abundance is of the same order as for the other elements. This is not the case in planetary nebulae, where no strong carbon line exists in the optical, implying that the carbon abundance determination is subject to a higher uncertainty than  the oxygen or the nitrogen abundance determination.   Finally, post-AGB stars are expected to extend to smaller masses than the nuclei of planetary nebulae. Indeed, planetary nebulae with central star masses lower than 0.55\\Ms\\ do not exist since the nebular gas has dissipated in the interstellar medium long before the star has become hot enough to ionize it. On the other hand, there is no such limitation for post-AGB stars.  This paper presents the first attempt to use a large sample of post-AGB stars to test AGB nucleosynthesis. In Sect. 2, we describe the constitution of our sample. In Sect. 3 we show some empirical diagrams and propose simple interpretations. In Sect. 4 we summarize the main findings and outline some prospects. ",
        "conclusions": "\\label{summary}  The main aim of this paper was to show the utility of  post-AGB stars to test theories of AGB nucleosynthesis. So far, the only tests of AGB nucleosynthesis based on large samples have been made using planetary nebulae. Post-AGB stars have several advantages over planetary nebulae: 1) abundances of a large variety of elements can be derived, including of $s$-process elements; 2) the abundance of carbon, an extremely important element for the diagnostics, is known with the same accuracy as the other elements, while in planetary nebulae, the carbon abundance is  significantly less reliable and more difficult to obain than that of O and N; 3) the determination of the atmospheric parameters  \\Teff, and gravity $g$ allows one to estimate the mass of the post-AGB star by comparison with theoretical stellar evolutionary tracks.  Of course, the study of post-AGB stars has its own difficulties. In particular, the abundance analysis is quite difficult  and many effects have to be considered in detail (see e.g. Asplund et al. 2000). The lack of suitable lines for reliable analysis is often a problem:  i) for cooler objects, the oxygen abundance is hard to derive, since useful lines of oxygen only start to show up at temperatures above  6000K; ii) nitrogen is often derived from a few red lines, which are known to suffer non-LTE effects; iii) stars that are not enriched in carbon, sometimes have  only a few suitable carbon lines.  We  considered all those objects from the present version of the catalogue of post-AGB objects (Szczerba et al.\\,\\cite{Szczerba_etal_2001}, Szczerba et al., in preparation) for which photospheric chemical abundances have been determined. We plotted diagrams based on these abundances, similar to the ones built for planetary nebulae studies. The same trends as for planetary nebulae were found, but in a clearer fashion (e.g. N/O vs. $\\epsilon$(O), revealing the effect of the ON cycle, or (C+N+O)/H vs. $\\epsilon$(C) indicating the presence of objects with C produced by the triple $\\alpha$ reaction). This is extremely encouraging and shows the interest of using post-AGB stars to complement planetary nebulae, despite the difficulties in abundance determinations.   Because the  post-AGB stars in our sample do not have all the same metallicities, we argued that a better indicator of third dredge-up and/or hot bottom burning is obtained by considering the photospheric abundances of C, N, O with respect to a metallicity indicator (and not with respect to H). It is therefore these ratios that we compared  with the solar ratios. We show that a convenient metallicity indicator is S (Fe cannot be used for post-AGB objects because dust depletion in former stages may have affected its present photospheric abundance).  Following the definition of Type I planetary nebulae by Kingsburgh \\& Barlow (1994) but accounting for the metallicity, we define a class of Type I post-AGB stars.  We show that non-type I objects are in vast majority of low mass ($M_{\\star}$ $<$ 0.56 M$_{\\odot}$). We also show  clear evidence that 3rd dredge-up is more efficient at low metallicity.     We have thus demonstrated the potential of post-AGB stars to constrain the models of AGB stars and the predicted yields.  The sample of post-AGB stars is likely to grow in the near future, thanks to  ASTRO-F, which is  much more sensitive than IRAS and should allow the discovery of many infrared-excess stars among which post-AGB stars are found. This will make the use of post-AGB stars even more attractive and powerful. In the present paper, we limited ourselves to only a few elements and to simple interpretations without direct comparison to models. We did not address the question of observational biases, which should be investigated by performing simulations on models. Future studies will address other issues related to AGB nucleosynthesis, such as the production of $s$-process elements, which are more easily done with post-AGB stars than with planetary nebulae."
    },
    "0601/astro-ph0601218_arXiv.txt": {
        "abstract": "In this paper, we present Sloan Digital Sky Survey (SDSS) photometry and spectroscopy in the fields of 27 gamma-ray bursts (GRBs) observed by {\\it Swift}, including bursts localized by {\\it Swift}, {\\it HETE-2}, and {\\it INTEGRAL}, after December 2004.  After this bulk release, we plan to provide individual releases of similar data shortly after the localization of future bursts falling in the SDSS survey area.  These data provide a solid basis for the astrometric and photometric calibration of follow-up afterglow searches and monitoring. Furthermore, the images provided with this release will allow observers to find transient objects up to a magnitude fainter than possible with Digitized Sky Survey image comparisons. ",
        "introduction": "Prompt long wavelength follow-up of gamma-ray bursts (GRBs) has revolutionized the study of these energetic and enigmatic objects.  With the successful launch of the recent high-energy missions such as the {\\it High Energy Transient Explorer 2} \\citep[{\\it HETE-2}\\,;][]{lamb2004}, {\\it International Gamma-Ray Astrophysics Laboratory} \\citep[{\\it INTEGRAL}\\,;][]{winkler2003},and {\\it Swift} \\citep{gehrels2004} satellites,  rapid follow-up of gamma-ray bursts has come to maturity.  {\\it Swift} and {\\it HETE-2} not only detect new GRBs but also provide rapid X-ray localization allowing for very precise (uncertainties as small as a few arcseconds) positions, far superior than possible in the past.   The size of  modern GRB samples have become large enough to study the statistical properties of GRBs \\citep[e.g.][]{berger2005c} and rapid follow-up observations have allowed for the first optical, radio, and X-ray localization of afterglows from short-hard gamma-ray bursts \\citep{gehrels2005,hjorth2005,villasenor2005,fox2005, pedersen2005,prochaska2005,covino2005,berger2005b,bloom2005}.  There has been large-scale success in identifying afterglows and conducting prompt spectroscopic follow-up; afterglows extending out to $z\\sim6.3$ \\citep[]{price2005,kawai2005}  have been spectroscopically confirmed providing luminous probes of the Universe even back to the era of reionization.  Recently, high resolution spectroscopy of GRBs has allowed for detailed studies of absorbing systems arising in the local environment of the GRB progenitor \\citep{berger2005,hwc2005}. Absorption line analyses of GRB afterglow spectra show structure similar to that of quasar damped Ly$\\alpha$ Absorbers (DLAs) but extend to higher hydrogen column densities and include several metal lines not found in quasar DLA systems \\citep{watson2005,hwc2005}.  It has been suggested that GRBs can provide a new class of standard candle for future cosmological experiments \\citep{ghirlanda2004,ghirlanda2004b, friedman2005,bloom2003}. As the luminosities of these objects should allow detection to very high redshifts ($z>10$) \\citep{lamb2000},  studies of GRBs could extend current work using SN Ia to high redshifts, thus providing considerable constraints on our understanding of  cosmology.  Before gamma-ray bursts can be used as a cosmological tool, however, a large number of bursts must be observed spectroscopically in order to calibrate the method.  Two keys to studying GRB afterglows are the identification of a transient afterglow coincident with the burst detected at high energies and high-quality information for in-field photometric calibration stars so that all observations can be placed onto a common photometric system. Currently, many afterglow searches compare new data to digitized Palomar Observatory Sky Survey photometric plates (DSS) and use USNO stars \\citep{monnet1998} to calibrate both astrometry and photometry.  As larger telescopes join the search for afterglows, the DSS imaging quickly becomes insufficient, as these images are too shallow for a detailed comparison to deep optical imaging, so afterglow candidates can only be identified through their temporal properties. The USNO catalog is a astrometric catalog, but it was not intended as a source of photometric calibration across the sky; USNO-calibrated photometry may thus be strongly affected by systematic photometric residuals in the USNO catalog itself.  The Sloan Digital Sky Survey \\citep[SDSS;][]{york2000} provides accurate photometry and astrometry for objects over $1/4$ of the sky, making it a viable alternative to the DSS images and USNO catalogs. The imaging from SDSS extends over a magnitude deeper than those from the DSS, making it a valuable resource in identifying transient objects while the stable photometric and astrometric calibration of the survey makes it an ideal source of calibration data in GRB fields.  In order to aid the community, we are releasing SDSS imaging and spectroscopy for the fields of 27 {\\it Swift}-observed GRBs detected after December 2004, including bursts originally localized by {\\it Swift}, {\\it HETE-2}, and {\\it INTEGRAL} observations.  Here, we offer a bulk release of GRB fields observed in the SDSS, but, in the future, we will release similar data through the GRB Coordinate Network (GCN) shortly after the detection of a new burst within the SDSS survey area.  In this document, we provide some basic documentation relating to the Sloan Digital Sky Survey and describe the data products included in each GRB release.  We hope that the information provided here will be a useful primer for those unfamiliar with SDSS data, but this document is by no means intended to be a full description of the survey or survey data products.  The layout of the paper is as follows: in \\S 2, we describe the Sloan Digital Sky Survey itself.  A brief description of photometric quantities measured in SDSS is given in \\S 3. We describe the data products provided with each GRB release in \\S4 and offer some concluding remarks in \\S5. ",
        "conclusions": ""
    },
    "0601/astro-ph0601243_arXiv.txt": {
        "abstract": "{ We discuss generation of magnetic field from cosmological perturbations. We consider the evolution of three component plasma (electron, proton and photon) evaluating the collision term between elecrons and photons up to the second order. The collision term is shown to induce electric current, which then generate magnetic field. There are three contributions, two of which can be evaluated from the first-order quantities, while the other one is fluid vorticity which is purely second order. We compute numerically the magnitudes of the former contributions and shows that the amplitude of the produced magnetic field is about $\\sim 10^{-19} {\\rm G}$ at 10kpc comoving scale at present. Compared to astrophysical and inflationary mechanisms for seed-field generation, our study suffers from much less ambiguities concerning unknown physics and/or processes. ",
        "introduction": "There are convincing evidences that imply existence of substantial magnetic fields in various astronomical objects. Not only galaxies but systems with even larger scales, such as cluster of galaxies and extra-cluster fields, have their own magnetic fields (for a review on cosmological magnetic fields, see e.g., [\\cite{Widrow02}]). Conventionally the magnetic fields in galaxies, and possibly in clusters of galaxies, are considered to have been amplified and maintained by dynamo mechanism. However, the dynamo mechanism needs the seed magnetic field and does not explain the origin of the magnetic fields.  There have been many attempts to generate the seed fields. One of the approaches to this problem is to generate magnetic fields astrophysically, often involving the Biermann mechanism [\\cite{Biermann50}]. This mechanism has been applied to various systems: large-scale structure formation [\\cite{Kulsrud97}], ionizing front [\\cite{Gnedin00}], protogalaxies [\\cite{Davies00}] and supernova remnant of the first stars [\\cite{Hanayama05}]. These studies showed the possibilities of magnetic-field generation with amplitudes of order $10^{-16} \\sim 10^{-21} {\\rm G}$, which would be enough for the required field $\\sim 10^{-20} {\\rm G}$.  On the other hand, cosmological origins, which are often concerned with inflation, can produce magnetic field with coherence lengths of much larger scales, which is typically the horizon scale or possibly super-horizon scale [\\cite{Turner88,Bamba04,Ashoorioon04,Bertolami99,Bertolami05}]. Also they can often produce fields with a wide range of length scales, while astrophysical mechanisms can produce fields only with their characteristic scales. For a constraints on magnetic field with cosmological scale, see [\\cite{Yamazaki04}].  Here, we consider magnetic-field generation from cosmological perturbations after inflation. Around and after the decoupling, coupling among charged particles and photons become so weak that electric current can be induced by the difference in motions of protons and electrons. This electric current leads to generation of magnetic fields. It is well known that vorticity of plasma produces such electric current [\\cite{Harrison70,Hogan00,Berezhiani04,Matarrese04,Gopal04,Lesch95}]. However, because vorticity is not produced at the first order in cosmological perturbations, we must study the second order. We will consider equations of motion for protons, electrons and photons separately up to the second order, although equation of motion for photons does not appear explicitly. To study electric current appropriately, we must treat the three components separately [\\cite{Gopal04}]. Furthermore, we will evaluate the collision term between electrons and photons, which is dominant and essential for the magnetic-field generation, up to the second order. From the equations of motion for protons and electrons, we obtain a generalized Ohm's law, which, combined with the Maxwell equations, leads to an evolution equation for magnetic field. Our study is based on the cosmological perturbation theory, which is highly successful in the anisotropies of the cosmic microwave background, and suffers from much less ambiguities concerning unknown physics and/or processes compared to astrophysical and inflationary mechanisms for seed-field generation. For details of our study, see [\\cite{KT05,KT06}].  \\begin{figure} \\resizebox{\\hsize}{!} {\\includegraphics[width=1cm,clip]{image.eps}} \\caption{All sky map of cosmological microwave background anisotropy obtained by the Wilkinson Microwave Anisotropy Probe (WMAP) satellite (upper) and schematic picture of cosmological magnetic fields generated from density fluctuations discussed here. Red  (blue) regions  are hot (cold) spots with a range of temperature $\\sim 2.725$ K $\\pm 200 \\mu $ K. The magnetic field vectors are shown together with the map. Strong magnetic fields are generated by currents where the gradient of density perturbation in photons is large. \\label{fig:image}} \\end{figure} ",
        "conclusions": ""
    },
    "0601/astro-ph0601075_arXiv.txt": {
        "abstract": "We report on two {\\sl ASCA} observations of the high-mass X-ray binary pulsar \\oao.  A short observation near mid-eclipse caught the source in a low-intensity state, with a weak continuum and iron emission dominated by the 6.4-keV fluorescent line.  A later, longer observation found the source in a high-intensity state and covered the uneclipsed through mid-eclipse phases.  In the high-intensity state, the non-eclipse spectrum has an absorbed continuum component due to scattering by material near the pulsar and 80 per cent of the fluorescent iron emission comes from less than 19~lt-sec away from the pulsar.  We find a dust-scattered X-ray halo whose intensity decays through the eclipse.  We use this halo to estimate the distance to the source as $7.1\\pm1.3\\rm\\ kpc$. ",
        "introduction": "\\oao\\ is one of the most poorly studied eclipsing high mass X-ray binary systems, and yet is potentially one of the most interesting. The reason this system is not so well known is because it took 15 years for its nature as a high mass X-ray binary to be revealed. \\oao\\ was discovered with the {\\sl Copernicus} satellite \\citep{Poli78}. A subsequent observation with {\\sl HEAO-1} revealed a pulse period of 38 s \\citep{WP79}. The 3--30~keV {\\sl Ginga} spectrum \\citep{Kama90} is a typical X-ray pulsar spectrum with a power law ($\\alpha \\sim 0.6$), a high energy cutoff ($E_c \\sim 5$~keV and $E_f \\sim 17$~keV) and an iron line at 6.6~keV (equivalent width $\\sim$ 240 eV).  For many years it was unclear if \\oao\\ was a low-mass or a high-mass system. The mystery was cleared up when the {\\sl CGRO} BATSE discovered a $\\sim 10.4$-d binary orbit based on long term monitoring of the pulse period in the 20--40~keV band \\citep{Chak93}. A $\\sim 1.7$ day eclipse of the X-ray source by its companion was also seen. This conclusively showed \\oao\\ to be a classic eclipsing massive X-ray binary system, probably similar to Vela X-1, making it the seventh eclipsing X-ray pulsar system to be found.  \\citet{Chak93} used the orbital parameters to infer that the companion is a supergiant of spectral class B0--B6. From the observed value of the spin period and its derivative during the spin-up interval they deduced an X-ray luminosity of $\\ga 1.6\\times 10^{37}$~ergs~s$^{-1}$, and thus a distance of $\\ga 11\\rm\\ kpc$. A more precise X-ray location obtained with Chandra allowed \\citet{Chak02} to make an infrared identification of the companion.  The infrared properties of the companion are consistent with a highly-reddened B supergiant at a distance of $6.4\\pm1.5\\rm\\ kpc$, implying an X-ray luminosity of $3\\times10^{36}\\rm\\ erg s^{-1}$.  \\oao's position on the Corbet diagram \\citep{Corb86} suggests that it may not be a typical system. The pulse periods of high-mass X-ray binaries (HMXRB) are distributed between 0.069 and 835~s, with no evidence for a clustering at any particular period. For the Be star systems there is a strong correlation between orbital period, $P_o$, and spin period, $P_s$. The supergiant systems have no strong dependence of orbital period on spin period. \\oao's 38~sec pulsar with its 10.4-d orbital period lies between the Be and supergiant systems on the Corbet diagram. The period measurements from {\\sl RXTE}, BATSE, and previous observations give a steady spin-up timescale of 125 yr \\citep{Bayk97} with short-term fluctuations \\citep{Bayk00}. This system may be in a short-lived phase where it is changing from a wind-accreting system like Vela X-1 to a disk-accreting system like Cen X-3. Thus, \\oao\\ provides an opportunity to test wind and/or disk accretion theory on a system which is in a transition between the two states. ",
        "conclusions": "The eclipse spectrum of the source in its high-intensity state is dominated by the decaying dust-scattered halo.  We estimate the radius of the grains to be $0.095\\pm0.005\\rm\\ \\mu m$ and the distance to the source to be $(7.9\\pm0.9)a_{0.1}^2\\ \\mbox{kpc}$, where $a_{0.1}$ is the grain radius in units of $0.1\\rm\\ \\mu m$.  Using our estimate for the grain size, the distance to the source becomes $7.1\\pm1.3\\rm\\ kpc$.  We find that the non-eclipse high-state continuum spectrum cannot be modeled by a single power law.  A more-absorbed component due to scattering by material near the pulsar is required.  The iron lines underwent partial eclipses during the 1997 observation.  We estimate that 80 per cent of the 6.4-keV fluorescent emission originated less than $19\\pm12$~lt-sec away from the pulsar while the 6.9-keV recombination line was mostly emitted by the extended, ionised stellar wind.  During the 1994 observation, the continuum emission was much lower and the iron emission during eclipse was dominated by the 6.4-keV flourescent line, which suggests that the source was in a low intensity state."
    },
    "0601/astro-ph0601305_arXiv.txt": {
        "abstract": "We present 4.5$\\mu$m and 8$\\mu$m photometric observations of 18 cool white dwarfs obtained with the $Spitzer$ Space Telescope. Our observations demonstrate that four white dwarfs with $T_{\\rm eff}<$ 6000 K show slightly depressed mid-infrared fluxes relative to white dwarf models. In addition, another white dwarf with a peculiar optical and near-infrared spectral energy distribution (LHS 1126) is found to display significant flux deficits in Spitzer observations. These mid-infrared flux deficits are not predicted by the current white dwarf models including collision induced absorption due to molecular hydrogen. We postulate that either the collision induced absorption calculations are incomplete or there are other unrecognized physical processes occuring in cool white dwarf atmospheres. The spectral energy distribution of LHS 1126 surprisingly fits a Rayleigh-Jeans spectrum in the infrared, mimicking a hot white dwarf with effective temperature well in excess of 10$^5$ K. This implies that the source of this flux deficit is probably not molecular absorption but some other process. ",
        "introduction": "The Sloan Digital Sky Survey (Adelman-McCarthy et al. 2005) has increased the number of known field cool white dwarfs from tens of objects to thousands (Kilic et al. 2006; Harris et al. 2006). In addition, Hubble Space Telescope observations of the globular clusters M4 (Hansen et al. 2004) and $\\omega$ Cen (Monelli et al. 2005) resulted in the discovery of more than two thousand white dwarfs. A careful analysis of these large datasets require a complete understanding of the structure and evolution of the white dwarf stars. The main uncertainty in the white dwarf luminosity functions derived from these datasets is caused by the inadequate description of energy transport in white dwarf atmospheres, which affects both the cooling rate and observational appearence of these stars (Hansen 1998).  Cool white dwarfs have atmospheres dominated by hydrogen or helium. Both hydrogen and helium are neutral below 5000 K and the primary opacity source in H-rich cool ($T_{\\rm eff}\\leq5500$ K) white dwarf atmospheres is believed to be collision induced absorption (hereafter CIA; Frommhold 1993) of molecular hydrogen (Bergeron et al. 1995; Hansen 1998; Saumon \\& Jacobson 1999). H-rich white dwarfs are predicted to become redder as they cool until the effects of CIA become significant below 5500 K. CIA opacity is strongly wavelength dependent and is expected to produce broad absorption features in the near-infrared. This flux deficiency can be seen in several stars, designated as ultracool white dwarfs (Oppenheimer et al. 2001; Harris et al. 2001; Gates et al. 2004; Farihi 2004; Farihi 2005).  There are two known opacity mechanisms in pure He-rich white dwarf atmospheres: Rayleigh scattering and He$^-$ free-free absorption. Rayleigh scattering is thought to be the dominant opacity source and therefore He-rich white dwarfs are expected to have blackbody like spectral energy distributions. On the other hand, Kowalski, Saumon \\& Mazevet (2005) argued that He$^-$ free-free absorption may be 2-3 orders of magnitude more significant, which would make it the dominant opacity source. In addition, Kowalski \\& Saumon (2004) showed that the atmosphere of He-rich white dwarfs with $T_{\\rm eff}<$ 10000 K should be treated as a dense fluid rather than an ideal gas, and that refraction effects become important. The density can be as high as 2 g cm$^{-3}$ in these atmospheres and the index of refraction departs significantly from unity. Kowalski, Saumon, \\& Mazevet (2005) used their updated models to match the observed sequence of cool white dwarfs from Bergeron, Ruiz, \\& Legget (1997; hereafter BRL) and suggested that the coolest white dwarfs have mixed H/He atmospheres.  Bergeron \\& Leggett (2002) tried to fit the optical and near-infrared photometry of two ultracool white dwarfs, LHS 3250 and SDSS 1337+00, and ruled out their models for pure hydrogen atmospheric composition for these stars. They found that the overall spectral energy distributions (SEDs) of these stars can be better fitted with mixed H/He models, yet the peak of the SEDs near 6000 \\AA\\ is predicted to be too narrow. There are only three white dwarfs with significant CIA that fits the current white dwarf models in the optical and near-infrared. WD0346+246 (Oppenheimer et al. 2001) and GD392B (Farihi 2004) both require low gravity (log $g\\sim$7), mixed H/He model atmosphere solutions. In addition, Bergeron et al. (1994) found a 5400 K, $log$ g=7.9, and $log$ N(He)/N(H) = 0.8 model atmosphere solution for LHS1126. BRL used new CIA opacity calculations and suggested that the helium abundance in this white dwarf is $log$ N(He)/N(H) = 1.86.  The $Spitzer$ Space Telescope (Werner et al. 2004) opened a new window into the Universe by enabling accurate mid-infrared photometry of faint objects (microjansky-level sensitivity). In order to understand the CIA opacity, and other unrecognized sources of opacity in cool white dwarf atmospheres, we used the $Spitzer$ Space Telescope to observe nearby, relatively bright, cool white dwarfs. In this paper, we present our mid-infrared photometry for 18 cool white dwarfs including LHS 1126. Our observations are discussed in \\S 2, while an analysis of these data and results are discussed in \\S 3. ",
        "conclusions": "Our Spitzer observations showed that all H-rich white dwarfs with $T_{\\rm eff}<7000$ K show slight mid-infrared flux deficits. Having several stars with small deficits makes these deficits significant. Moreover, LHS 1126 shows significantly depressed mid-infrared fluxes relative to white dwarf models. Debes et al. (2005, private communication) have also found 10-20\\% flux deficits ($\\geq3\\sigma$ significance) at 4.5, 5.6, and 8$\\mu$m in two DAZ white dwarfs with $T_{\\rm eff}=$ 6820 K and 7310 K. Combining the facts that WD0552-041 (with a pure-He atmosphere) does not show a clear flux deficit, and all H-rich white dwarfs cooler than 7000 K do exhibit mid-infrared flux deficits, LHS1126 and WD1748+708 are likely to have mixed H/He atmospheres. This is also consistent with Bergeron et al.'s (1994) and Wolff et al.'s (2002) analysis. However, one question remains to be answered: can these flux deficits be explained by CIA?  The atmospheres of M, L, and T dwarfs are also expected to exhibit CIA (Borysow et al. 1997). Roellig et al. (2004) obtained mid infrared spectroscopy of M, L, and T dwarfs in the 5 -- 15 $\\mu$m range. Their Figure 2 shows that the observed spectra are in good agreement with the model atmopshere calculations, with only a few minor deviations. Hence, CIA calculations for low density atmospheres seem to be accurate.  The near-infrared (1--2 $\\mu$m) flux deficit in LHS 1126 was discovered earlier by Wickramasinghe et al. (1982). Bergeron et al. (1994) and BRL explained this deficit as CIA by molecular hydrogen due to collisions with helium and found the H/He ratio to be $\\sim$0.01. Wolff et al. (2002) used Faint Object Spectrograph data plus optical and infrared photometry of LHS1126 to model this star's SED in the 0.2 -- 2.2 $\\mu$m range. They found that the hydrogen abundance reported by Bergeron et al. (1994) would result in an extremely strong Lyman $\\alpha$ absorption, and the SED is best fitted with an abundance ratio of H/He = 3 $\\times$ 10$^{-6}$. Wolff et al. (2002) and Bergeron et al. (1994) models do not give a consistent picture for the H/He ratio and neither model is adequate. We revisit this problem by extending our wavelength coverage to 8$\\mu$m. Figure 4 shows the ultraviolet spectrum (kindly made available to us by D. Koester) and optical and infrared photometry of LHS 1126 along with a 5400 K blackbody (dotted line). Mixed atmosphere white dwarf models (kindly made available to us by D. Saumon) with $T_{\\rm eff}=5400$ K, $log$ g= 7.9, and $log$ N(He)/N(H) = 1.5 (lower solid line) and $log$ N(He)/N(H) = 1 (upper solid line) are also shown. These white dwarf models include CIA opacities but not Lyman $\\alpha$ absorption, therefore they cannot be used to match the ultraviolet data. If we just use the optical and near-infrared photometry (as in Bergeron et al. 1994), we could fit the observations with a $log$ N(He)/N(H) = 1 -- 1.5 model. However, these models cannot explain the observed flux deficits in the Spitzer observations. Bergeron et al. (1995) and Hansen (1998) models cannot explain these flux deficits either (B.M.S. Hansen and P. Bergeron 2005, private communication). CIA is expected to create wiggles and bumps in the spectrum up to 3$\\mu$m, and the SEDs of cool white dwarfs are expected to return to normal in the mid-infrared (Figure 4; see also Figure 6 of Jorgensen et al. 2000).  Figure 4 also shows a power law distribution with $\\alpha=2$, i.e. a Rayleigh-Jeans spectrum (dashed line). The infrared spectral energy distribution of LHS 1126 follows a Rayleigh--Jeans spectrum fairly well. It is striking to note that we can fit the ultraviolet and optical part of the SED with a 5400 K blackbody, and the infrared part of the SED with $T_{\\rm eff}>10^5$ K. The power law is a reasonable but not a perfect fit to the infrared spectral energy distribution of this object, and photometry in the other IRAC channels or spectroscopy with the Spitzer Telescope may reveal more structure in the mid-infrared, but the current data suggest the following:  1 -- The CIA opacity calculations are incomplete, and there is a significant, unexplained flux deficit in the mid-infrared. As the densities in white dwarf atmospheres approach a significant fraction of the liquid state densities, molecular absorption (H$_2$-H$_2$, H$_2$-He), emission, and light scattering are expected to become increasingly important. Current CIA opacity calculations in dense media may be incomplete in certain wavelength regimes, for example the roto-vibrational band of H$_2$ CIA may be enhanced compared to predictions (L. Frommhold and D. Saumon 2005, private communication). On the other hand, the featureless spectra of the ultracool white dwarfs (Gates et al. 2004), and the reasonably good fit of a Rayleigh-Jeans distribution to the LHS 1126 SED in the infrared, suggest another mechanism to explain the mid infrared flux deficits observed in white dwarfs. CIA is expected to disappear  with the dissociation of molecules when $T_{\\rm eff}$ increases past 5500 K. Therefore, white dwarfs with $T_{\\rm eff}>6000K$ that show infrared flux deficits could not be explained by changing the CIA opacities.  2 -- The mid-infrared flux deficits are caused by some as-yet unrecognized physical process(es). The problem may be caused by an unknown or poorly understood absorption process, or  it may be the result of a different source function operating in these dense atmospheres. H$^-$ bound-free absorption is thought to be the most important source of opacity in warmer DA white dwarfs. More work is required to test if possible changes in this absorption could help explain the mid-infrared flux deficits. The input physics used in white dwarf model atmospheres is mostly based on the ideal gas approximation. The extreme conditions in white dwarf atmospheres require a new look at dense medium effects on the equation of state, chemistry, opacities, and radiative transfer. Refractive opacities, presence of heavy elements, or formation of trace species and many other factors can change the opacity sources in white dwarf atmospheres (see Kowalski 2005 for a detailed discussion).  Near/mid-infrared photometry and spectroscopy of ultracool white dwarfs and the so-called C$_2$H stars will be useful to test these ideas. At this time, we only report and do not understand the observed mid-infrared flux deficits, but this mystery needs to be resolved before we can use the cool white dwarfs as accurate chronometers to find the age of Galactic populations."
    },
    "0601/astro-ph0601419_arXiv.txt": {
        "abstract": "{NGC\\,7538 is known to host a 6.7 and 12.2\\,GHz methanol maser cospatial with a Ultra Compact (UC)\\,\\hii region, IRS~1.}{We report on the serendipitous discovery of two additional 6.7\\,GHz methanol masers in the same region, not associated with IRS~1.} {Interferometry maser positions are compared with recent single-dish and interferometry continuum observations.} {The positions of the masers agree to high accuracy with the 1.2\\,mm continuum peak emission in NGC\\,7538~IRS~9 and NGC\\,7538~S. This clear association is also confirmed by the positional agreement of the masers with existing high resolution continuum observations at cm and/or mm wavelengths.}{Making use of the established strong relation between methanol masers and high-mas star formation, we claim that we have accurately positioned the high-mass protostars within the regions where they are detected. The variety of objects hosting a 6.7\\,GHz methanol maser in NGC\\,7538 shows that this emission probably traces different evolutionary stages within the protostellar phase. } ",
        "introduction": "The NGC\\,7538 nebula  is a furnace of intense massive star formation located at some 2.7~kpc from the Sun. At least 11 infrared sources have been identified embedded within it (Fig.~\\ref{fig:ngc_allIR}, \\citealt{ojh04} and references therein). These IR sources lie in four regions hosting a large range of different types of young stellar objects (YSOs) with signs of decreasing age moving from the North-West to the South-East. In the far North-West the first star formation region contains the IR sources IRS~5, 6, 7 which are associated with a developed \\hii region of $\\sim3$~pc in size. IRS~6 has been proposed as its exciting source \\citep{ojh04}. A further star formation region can be identified with the IR sources IRS~1-3, where infrared emission  at K-band and long-ward is dominated by IRS~1 \\citep{deb05}. IRS~1 is known to power a Ultra Compact (UC)\\,\\hii region, an early stage of massive star formation. Masers of different species have also been detected toward this source (OH, H$_2$O, \\meth, see \\citealt{min98} and  references therein). A third star-forming region is centered on NGC\\,7538~S and also contains IRS~11, 10$\\arcsec$ to its northwest (Fig. \\ref{fig:ngc_allIR}). Coincident with NGC\\,7538~S OH and H$_2$O masers have been detected \\citep{kam90}. Finally, furthest to the South-East, the fourth known region of star formation contains the deeply embedded IR source IRS~9, believed to be in a stage prior to the formation of a UC \\hii region. H$_2$O masers were also detected toward this source \\citep{dav98,kam90}.  Methanol masers are indicative of the earliest stages of massive star formation. They arise from cold, embedded molecular clumps and in some cases from dark infrared clouds \\citep{hil05,pur05}. These clumps are often interpreted to be protoclusters of high mass YSOs, i.e. precursors of OB associations (e.g. \\citealt{mot03,min05}). High angular resolution observations of methanol maser environments reveal that these masers are associated with hot cores, UC\\,\\hii regions and very often do not coincide with radio continuum emission from bright Ultra Compact (UC)\\,\\hii regions (\\citealt{wal98,min01}). The detection of a methanol maser in such regions is a clear indication of the presence of a deeply embedded hot core that is characterised by large methanol abundance and 100-200\\,K gas and dust temperature. These physical conditions are in agreement with those derived from maser modelling, as e.g. \\citealt{sob97}. The high astrometric accuracy of VLBI observations of methanol masers gives thus the best estimate of the position of the high-mass protostellar core at a scale of $\\sim$ 10\\,AU. One clear example of this fact is given by the main component of the \\meth{} maser in NGC\\,7538 (component \\A{} in Fig. \\ref{fig:spectrum}): the maser pinpoints the massive  protostellar core within the IRS~1 submm clump to milliarcsecond accuracy (1\\,mas at $\\sim$3\\,kpc is $\\sim$3\\,AU, see \\citealt{pes04a}).  In this letter we report on the serendipitous detection of 6.7\\,GHz methanol masers in the third and fourth star formation regions in NGC\\,7538, cospatial with  NGC\\,7538~S and IRS~9. Supported by what exposed above, we argue that with this discovery we have accurately positioned the two massive protostars in those regions. This is then confirmed by the coincidence of the new masers with the peak of the 1.2\\,mm dust emission coming from the thick cocoon surrounding them. ",
        "conclusions": "\\label{sec:concl}  From the high spatial resolution MERLIN observations of 6.7\\,GHz methanol maser emission toward the massive star forming region NGC\\,7538 we have detected and positioned two new spectral features. We note that this achievement was possible only thanks to the intermediate resolution of the MERLIN with its shortest baselines. We conclude the following: \\begin{itemize} \\item on the basis of the accepted idea that methanol masers are the most efficient tracers of young and embedded massive protostars, the accurate location (to within 50\\,mas or 135\\,AU) of \\F{} and \\G{} is equivalent with the accurate location of the massive protostars in NGC\\,7538~IRS~9 and NGC\\,7538~S; \\item the variety of objects marked by methanol maser emission could be due to different ages or viewing angle geometry. To test these two scenarios, more accurate observations (targeted toward \\F{} \\& \\G) are required.  \\end{itemize}"
    },
    "0601/astro-ph0601133_arXiv.txt": {
        "abstract": "The upcoming generation of SZE surveys will shed fresh light onto the study of clusters.  What will this new observational window reveal about cluster properties?  What can we learn from combining X-ray, SZE, and optical observations?  How do variations in the gas entropy profile, dark matter concentration, accretion pressure, and intracluster medium (ICM) mass fraction affect SZE observables?  We investigate the signature of these important cluster parameters with an analytic model of the ICM.  Given the current uncertainties in ICM physics, our approach is to span the range of plausible models motivated by observations and a small set of assumptions.  We find a tight relation between the central Compton parameter and the X-ray luminosity outside the cluster core, suggesting that these observables carry the same information about the ICM.  The total SZE luminosity is proportional to the thermal energy of the gas, and is a surprisingly robust indicator of cluster mass: $L_{\\rm SZ} \\propto f_{\\rm ICM} M^{5/3}$.  We show that a combination of $L_{\\rm SZ}$ and the half-luminosity radius $r_{\\rm SZ}$ provides a measure of the potential energy of the cluster gas, and thus we can deduce the total energy content of the ICM.  We caution that any systematic variation of the ICM mass fraction will distort the expected $L_{\\rm SZ} - M$ calibration to be used to study the evolution of cluster number density, and propose a technique using kSZ to constrain $f_{\\rm ICM}(M,z)$.\\\\ ",
        "introduction": "\\label{intro} As the largest and most recently formed relaxed objects in the universe, clusters provide cosmological constraints complementary to other observations.  \\citet{press/schechter:1974} showed that their co-moving number density is exponentially sensitive to both cluster mass and the amplitude of the linear power spectrum of density fluctuations.  The number density of massive galaxy clusters in the local universe constrains a model-dependent combination of the amplitude of fluctuations at $8 \\; h^{-1}$ Mpc ($\\sigma_{8}$) and the matter density ($\\Omega_{\\rm m}$).  The redshift evolution of cluster number counts depends on the linear growth factor and and the co-moving volume element, and can thereby potentially break the degeneracy in the family of $\\Lambda$ and quintessence models allowed by primary CMB anisotropy \\citep{wang/steinhardt:1998}.  X-ray and optical cluster counts at low and intermediate redshifts have been used to constrain cosmological parameters \\citep[for recent results, see][]{allen/etal:2003, allen/etal:2004a, bahcall/etal:2003, henry:2004, ikebe/etal:2002, pierpaoli/etal:2003, rapetti/allen/weller:2005, reiprich/bohringer:2002, rosati/borgani/norman:2002, shimizu/etal:2003, viana/nichol/liddle:2002, vikhlinin/etal:2003}.\\\\ The Sunyaev-Zel'dovich effect (SZE) flux is an excellent probe of high redshift clusters, and many SZE surveys are in development, such as ACT\\footnote{http://www.hep.upenn.edu/act/}, APEX\\footnote{http://bolo.berkeley.edu/apexsz/index.html}, Planck\\footnote{http://www.rssd.esa.int/index.php?project=PLANCK\\&page=index}, and SPT\\footnote{http://spt.uchicago.edu/}.  While the information from high redshift available in SZE surveys can potentially provide tight cosmological constraints \\citep{haiman/mohr/holder:2001}, recent work suggests that uncertainties in the mass-observable relation and its scatter can severely degrade their sensitivity.  Fortunately, internal calibration and some follow-up observations can recover much of a survey's sensitivity {\\em if} the mass-observable relation and its scatter can be described accurately in the relevant redshift range by a reasonably small set of parameters (\\citealt{hu:2003, majumdar/mohr:2004, lima/hu:2004, lima/hu:2005}; but see also \\citealt{francis/bean/kosowsky:prep}).  The results presented in this paper support their assumption: we find a tight mass-observable SZE relation that is robust to a wide range of model variations.\\\\ The SZE is a distinct spectral signature in the CMB due to Thomson scattering of CMB photons and hot electrons with a magnitude proportional to the integrated electron pressure along the line of sight \\citep[for discussions of SZE physics, see][]{birkinshaw:1999, carlstrom/holder/reese:2002}.  A unitless measure of the effect is the Compton parameter $y$ \\citep{sunyaev/zeldovich:1972}: \\begin{equation} \\label{compton} y = \\frac{\\sigma_{Thomson}}{m_e c^{2}} \\int P_{e} \\; dl. \\end{equation} We define the SZE luminosity $L_{\\rm SZ}$ as the integration of $y$ over the projected cluster area, so that \\begin{equation} \\label{SZflux} f_{\\rm SZ} d_{A}^{2} = L_{\\rm SZ} = \\int y \\; dA = \\frac{\\sigma_{Thomson}}{m_e c^{2}} \\int P_{e} dV = \\frac{\\sigma_{Thomson}}{m_e c^{2}} N_{e} \\langle T_{e} \\rangle. \\end{equation} Assuming local thermal equilibrium, $L_{\\rm SZ}$ directly measures the total thermal energy content of the intracluster medium (ICM), and $y(\\theta)$ informs us on how the gas is distributed in the cluster.  Because $d_{A}$ flattens at intermediate redshift, the SZE flux $f_{\\rm SZ}$ becomes approximately redshift independent.\\\\ The self-similar collapse model for the ICM considers only gravitational physics and predicts scaling relations between cluster properties.  X-ray observations indicate significant deviations from these predictions, most notably in the X-ray luminosity-temperature relation \\citep{markevitch:1998, arnaud/evrard:1999}.  Many complex non-gravitational processes determine the final thermal state of the ICM.  We discuss many of them in Sec.~\\ref{xraysec} along with related X-ray observations.   In Sec.~\\ref{results} we introduce an analytic model of smooth accretion \\citep{voit/etal:2003} that provides a conceptual foundation for determining the gas properties.  We then relax some of the assumptions of \\citet{voit/etal:2003} and parametrize our uncertainties in ICM physics: the poorly understood effects of heating, cooling, and conduction is encoded in the gas entropy profile and the ICM mass fraction, variation in the dark matter potential is characterized by the concentration parameter, and the boundary condition is set by an accretion pressure.  Motivated by the physics and observations discussed in Sec.~\\ref{xraysec}, our main goal is to explore SZE and X-ray observables over a plausible range of parameter values using a set of phenomenological models under the assumptions of spherical symmetry and hydrostatic equilibrium.   We compare the resulting profiles to recent observations of nearby, relaxed clusters reported in \\citet{vikhlinin/etal:2005}.  We quantify deviations from the expected $L_{\\rm SZ} - M$ relationship and indicate the main sources of scatter accessible to our models.  In Sec.~\\ref{relations} we describe relations between cluster properties and observables that hold across the range of models we consider.  In particular, we show that SZE observables alone provide a robust measure of the total energy content of the ICM.  We discuss the implications and limitations of our models in Sec.~\\ref{conc}, and focus on the important point of constraining the ICM cluster mass fraction $f_{\\rm ICM}$ using the kinetic Sunyaev-Zel'dovich effect (kSZ) in Sec.~\\ref{fbdiscussion}.  Sec.~\\ref{future} states our conclusions and the implications of this work for future SZE surveys.\\\\ ",
        "conclusions": "\\label{future} SZE surveys will soon identify thousands of clusters and measure their photometric redshifts, so we will be able to estimate the gas potential and thermal energy with $L_{\\rm SZ}$ and $r_{\\rm SZ}$, as well as the cluster mass (see Fig.~\\ref{fig:SZ1} and Fig.~\\ref{fig:potential}).  Detection of the kSZ signal in these SZE surveys will also constrain the ICM mass fraction, $f_{\\rm ICM}(M,z)$.  Follow-up observations in X-ray, optical, and radio can provide checks on the survey $L_{\\rm SZ}-M$ relation using independent mass determinations, and could further constrain ICM models through AGN activity-gas energy correlations.  Coupling sufficiently high resolution SZ and X-ray observations could over-constrain the gas density and temperature profiles, and thus provide a consistency check of deprojection techniques.  Sensitive SZE profiles should extend farther into the outer regions of the cluster than X-ray observations allow.\\\\ This paper has focused on SZE observables from $z=0$ clusters, though SZE surveys will find most of their clusters at $z \\sim 0.6$.  However, the models studied here support the self-similar prediction $T_{gas} \\sim T_{dark}$ even under considerable variations of the total gas energy.  We expect $T_{dark}$ to scale with the overdensity $\\left(\\Delta(z) \\rho_{cr}(z)\\right)^{1/3}$, and the dark matter concentration parameter will decrease with redshift as suggested by Eqn.~\\ref{conceqn}: by a factor of $\\sim 1.5$ at $z=1$.  Appendix A suggests that changes in $L_{\\rm SZ}$ induced by a decrease in $c_{\\Delta}$ with redshift is limited to less than $15\\%$.  The range of entropy normalizations examined in this paper is larger than the expected characteristic entropy decrease by $\\sim 20\\%$ at $z=1$ under self-similar evolution.  Because we have found a very weak dependence of $L_{\\rm SZ}$ on the entropy profile, we are optimistic that the redshift evolution of the $L_{\\rm SZ}-M$ relation will be determined by the more easily understood dark matter evolution; this assumption has been verified in numerical simulations including galaxy formation \\citep{nagai:prep}.  In this case the $L_{\\rm SZ}-M$ evolution is characterized by relatively few parameters, and so can hopefully be internally calibrated in SZE surveys (\\citealt{hu:2003, majumdar/mohr:2004, lima/hu:2004, lima/hu:2005}; but see also \\citealt{francis/bean/kosowsky:prep}).  However, we caution that baryon physics may induce a redshift dependent ICM mass fraction $f_{\\rm ICM}(M, z)$ that produces a deviation from simple self-similar redshift evolution.  The observed ICM mass fraction in clusters appears to depend on cluster mass and radius.  In Sec.~\\ref{fbdiscussion} we proposed the use of the kSZ signal to constrain $f_{\\rm ICM}(M,z)$.  Understanding the variation in ICM mass fraction will deepen our understanding of cluster physics and enable the use of clusters as cosmological probes.  Finally, we may hope to probe the gas's thermal history by measuring the thermal and potential energies as a function of redshift through SZE observation.\\\\"
    },
    "0601/astro-ph0601655_arXiv.txt": {
        "abstract": "We show in this paper why molecular millimeter absorption line searches in DLAs have been unsuccessful. We use CO emission line maps of local galaxies to derive the \\htwo\\ column density distribution function $f(\\nhtwo)$ at $z=0$.  We show that it forms a natural extension to \\fnhi: the \\htwo\\ distribution exceeds \\fnhi\\ at $N_{H} \\approx 10^{22} \\icmsq$ and exhibits a power law drop-off with slope $\\sim -2.5$. Approximately 97\\% of the \\htwo\\ mass density $\\rho_{\\rm H_2}$ is in systems above $\\nhtwo=10^{21}\\,\\icmsq$. We derive a value $\\rho_{\\rm H_2} = 1.1 \\times 10^7 h_{70}\\,\\msol {\\rm Mpc}^{-3}$, which is $\\approx 25\\%$ the mass density of atomic hydrogen. Yet, the redshift number density of \\htwo\\ above this \\nhtwo\\ limit is only $\\approx 3\\times 10^{-4}$, a factor 150 lower than that for \\hi\\ in DLAs at $z=0$. Furthermore, we show that the median impact parameter between a $\\nhtwo>10^{21}\\,\\icmsq$ absorber and the centre of the galaxy hosting the \\htwo\\ gas is only 2.5 kpc. Based on arguments related to the Schmidt law, we argue that \\htwo\\ gas above this column density limit is associated with a large fraction of the integral star formation rate density. Even allowing for an increased molecular mass density at higher redshifts, the derived cross-sections indicate that it is very unlikely to identify the bulk of the molecular gas in present quasar absorption lines samples. We discuss the prospects for identifying this molecular mass in future surveys. ",
        "introduction": "At high redshift, the atomic hydrogen is well surveyed through observations of damped \\lya\\ systems (DLAs), quasar absorption line systems characterized by $\\nhi>2 \\times 10^{20}\\,\\icmsq$ \\citep[e.g.,][]{Prochaska2005}.  These observations yield the frequency distribution of \\hi\\ surface density  \\fnhi\\ and the first moment gives the cosmological mass density of predominantly neutral, atomic gas.  Several small surveys have been performed to search for molecular gas associated with DLAs. Yet, there have been no positive detections of CO and other molecules in millimeter wavelength absorption line searches \\citep[][and references therein]{Curran2004b}. Searches for \\htwo\\ absorption in DLAs (via the Lyman and Werner bands) show a success rate of only $\\approx 20\\%$ and even these sightlines have low molecular fractions \\citep{Ledoux2003}.  These low detection rates have drawn into question the relationship between DLAs and star formation at high redshift \\citep[e.g.][]{Hopkins2005}. Other galaxy surveys have shown that star formation is very active at these epochs \\citep[e.g.,][]{Steidel1999}. Presumably, since stars form in molecular clouds, these star forming galaxies have significant molecular gas mass that has not been discovered through quasar absorption line studies. The obvious conclusion, then, is that these molecular regions have very low covering fraction on the sky and/or contain sufficient columns of dust to obscure any background quasar.  In this paper we address these issues by making use of the studies of the CO distribution in nearby galaxies. \\citet{Zwaan2005b} recently showed that most DLA properties (luminosities, impact parameters between quasars and DLAs, and metal abundances) are consistent  with them arising in galaxies like those in the local universe. Building on this result, we use here information on the molecular content of nearby galaxies, to make predictions of the detectability of \\htwo\\ in DLAs. The \\htwo\\ distribution in nearby galaxies is generally derived through mapping of the CO(1--0) line at a frequency of 115 GHz. The \\htwo\\ column densities derived from these observations are typically higher than $5\\times 10^{20}\\,\\icmsq$, which is very close to those reached by millimeter CO absorption line observations in redshifted DLAs against radio-loud background sources. ",
        "conclusions": "We have used observations of CO in nearby galaxies to make predictions on the detection rate of \\htwo\\ in DLAs. We derived a column density distribution function \\fnhtwo\\ and find that it forms a natural extension of \\fnhi\\ at higher hydrogen column densities. The inferred redshift number density $dN/dz$ of $\\nhtwo>10^{21}\\,\\icmsq$ is $3 \\times 10^{-4}$, which is a factor 150 lower than the corresponding number for DLAs. This implies that of the total number of DLAs currently known, only a handful are expected to show high column density molecular absorption lines. This absorption is expected to arise at impact parameters smaller than 6.5 kpc in 90\\% of the cases. Approximately 97\\% of the cosmic \\htwo\\ mass density resides in this dense, central gas, which also is associated with the bulk of the cosmic star formation rate density \\sfr. The detectability of this molecular gas is further reduced by the facts that  {\\em i)\\/} the high column density \\htwo\\ is likely associated with high dust columns, which obscures background quasars and could take them out of magnitude-limited surveys, and {\\em ii)\\/} a small fraction ($\\la 10$\\%) of the \\htwo\\ is associated with sub-DLA \\lya\\ absorption.  {We apply the Schmidt law for star formation to \\fnhtwo\\ and \\fnhi\\ and find that molecules should be taken into account to derive meaningful values for the cosmic star formation rate density \\sfr. Even though the cross section of high column density \\htwo\\ is low,  this gas is associated with a large fraction of \\sfr. Since the molecular gas is difficult to detect in absorption, the method of using the Schmidt law and \\fn\\ to derive \\sfr\\ remains of limited use.}   The cold atomic gas content of the universe at redshifts $z\\sim 1-5$ has successfully been determined from blind surveys for DLA absorbers. What are the prospects for measuring the molecular mass density at intermediate and high redshifts using similar techniques? With present technology blind molecular absorption line surveys against radio-loud background sources seem difficult given the required redshift interval. The prospects for future instruments seem better. For example, with the Atacama Large Millimeter Array\\footnote{http://www.alma.info/} (ALMA) the CO(3--2) line corresponding to $\\log\\nhtwo>21$ should be detectable at $z=2$ to $3$ in one minute against a 50 mJy background source. A survey of several days should be able to turn up a useful number of high \\nhtwo\\ absorbers. {Such a survey would by simultaneously sensitive to HCO$^+$ (a good \\htwo\\ tracer) at lower redshifts.} The Square Kilometer Array\\footnote{http://www.skatelescope.org/} (SKA) should be able to detect the CO(1--0) and HCO$^+$(1--0) lines at redshifts $z>3.6$ and $>2.6$, respectively. The expected noise levels are much lower than for ALMA, implying that this instrument will be ideal for high-$z$ molecular absorption line surveys, giving good statistics for measuring $\\Omega_{{\\rm H}_{2}}$."
    },
    "0601/astro-ph0601180_arXiv.txt": {
        "abstract": "We report the results of a study of X-ray point sources coincident with the High Velocity System (HVS) projected in front of NGC\\,1275. A very deep X-ray image of the core of the Perseus cluster made with the \\emph{Chandra Observatory} has been used. We find a population of Ultra-Luminous X-ray sources (ULX; 7 sources with L$_{\\rm X}(0.5-7.0 \\rm{\\,keV})>7\\times 10^{39}~ \\rm{erg ~ s^{-1}}$).  As with the ULX populations in the Antennae and Cartwheel galaxies, those in the HVS are associated with a region of very active star formation. Several sources have possible optical counterparts found on \\emph{HST} images, although the X-ray brightest one does not. Absorbed power-law models fit the X-ray spectra, with most having a photon index between 2 and 3. ",
        "introduction": "\\begin{figure*} \\includegraphics[width=0.8\\textwidth]{lobulos.jpg.eps} \\caption{NGC\\,1275 in the 0.3--7.0 keV band. Pixels are 0.49 arcsec in size and the N-S height of the image is 97~arcsec.  North is to the top and east is to the left in this image.  The point sources which we identify as ULX are located in the high-velocity system is seen in absorption to the north of the bright nucleus. None are seen in the south lobe}\\label{fig:whole_picture} \\end{figure*} \\begin{figure*} \\includegraphics[width=0.8\\textwidth]{HVS.jpg.eps} \\caption{The central region of NGC\\,1275 in the 0.3--0.8 keV band. Pixels are 0.49 arcsec in dimension and the entire image is $1.77\\times 0.89 ~ \\rm{arcmin^{2}}$. North is to the top and east is to the left in this image. The high-velocity system is seen in absorption to the north of the bright nucleus which is at RA $3^{\\rm h} 19^{\\rm m} 48^{\\rm s}$, Dec. +$41^o$30'42\"}\\label{fig:HVS} \\end{figure*}  The study of Ultra-Luminous X-ray sources (ULX) has been greatly expanded by the high spatial resolution and spectral grasp of the \\emph{Chandra} and \\emph{XMM-Newton} observatories, respectively. ULX sources (Fabbiano \\& White 2003; Miller \\& Colbert 2004) have 2--10~keV X-ray luminosities exceeding $10^{39}\\ergps$ and are found some distance from the centres of galaxies; they are not active galactic nuclei.  Their luminosity exceeds that for a $10\\Msun$ black hole accreting at the Eddington limit which radiates isotropically and so have created much interest in the possibility that they contain even higher mass black holes, such as InterMediate Black Holes (IMBH) of $\\sim10^3\\Msun$ (Makishima et al 2000; Miller, Fabian \\& Miller 2004). Alternatively they may appear so luminous because of beaming (Reynolds et al 1999; King et al 2001, Zezas \\& Fabbiano 2002) or due to super Eddington accretion (Begelman 2002).  ULX are most common in starburst galaxies and in very active star-forming regions, such as in the Antennae and the Cartwheel galaxy, where populations of tens of them are found (Zezas et al 2002; Gao et al 2003; Wolter \\& Trinchieri 2004). In some cases variability rules out the possibility that they are just clusters of lower-luminosity objects. The origin of IMBH is unclear. They may form as a result of binary interactions in dense stellar environments (Portegies Zwart \\& McMillan 2002).  A comparison of IMBH ULX candidates with a number of well known stellar-mass black holes candidates (BHC; Miller et al 2004) demonstrates that the ULX are more luminous but have cooler thermal disk components than standard stellar-mass BHC.  Therefore, ULX in this sample are clearly different from the sample of stellar-mass BHC and are consistent with being IMBH.  Here we report on the discovery of a population of 8 point X-ray sources to the N of the nucleus of NGC\\,1275, which is the central galaxy in the Perseus cluster. All exceed $10^{39}\\ergps$ in X-ray luminosity, and 7 are formally ULX, if they are at the distance of the cluster. The spatial region where they lie coincides with the High Velocity System of NGC\\,1275. We assume that they are part of that system. We see no other point sources (apart from the nucleus) over the body of NGC\\,1275 (Fig.~1).  NGC\\,1275 is embedded in a complex multiphase environment. Optical imaging and spectroscopy first established the existence of two distinct emission-line system toward NGC\\,1275: a low-velocity component associated with the galaxy itself at 5200 km $\\rm{s^{-1}}$ and a high-velocity component at 8200 km $\\rm{s^{-1}}$ projected nearby on the sky (Minkowski 1955, 1957). This latter component is associated with a small gas-rich galaxy falling into the cluster along our line of sight (Haschick, Crane \\& van der Hulst 1982). A merger scenario has been proposed (Minkowski 1955, 1957). However, interaction of the low and/or high-velocity system with a third gas-rich galaxy or system of galaxies (Holtzman et al. 1992; Conselice, Gallagher \\& Wyse 2001), or influences from the surrounding dense intracluster medium (ICM) (Sarazin 1988; Boroson 1990; Caulet et al. 1992) have been discussed.  Deep \\emph{Chandra} observations have clarified the position of the High Velocity System (HVS).  The depth of the observed X-ray absorption (e.g. Fig.~1) is nor infilled by emission from hot gas projected along the line-of-sight so the HVS must lie well in front of NGC\\,1275. Gillmon, Sanders \\& Fabian (2004) have estimated a lower limit on the distance of the HVS from the nucleus of 57 kpc.  The low- and high-velocity system are therefore not yet directly interacting. The HVS $\\emph{is}$ however strongly interacting with the ICM of the Perseus cluster, which has triggered strong star formation.  In this paper we describe the detailed analysis of the X-ray spatial and spectral properties of the discrete sources in the high velocity system.  The paper is organized as follows: in Sect. 2 and 3 we present reduction and results from the imaging analysis and spectral analysis, respectively; in Sect. 4 we discuss the results; and Sect. 5 summarizes our findings.  Throughout this paper we use a redshift of 0.018 and $H_0=70~\\rm{km~s^{-1}~Mpc^{-1}}$. This gives a luminosity distance to the cluster of 80 Mpc; 1 arcsec corresponds to a physical distance of 370~pc. ",
        "conclusions": "We have described the detailed analysis of the spatial and spectral properties of the discrete X-ray sources detected with a deep \\emph{Chandra} ACIS-S observation around NGC\\,1275. Our results are summarized below:  \\begin{enumerate} \\item We have detected a total of 8 sources to the north of NGC\\,1275 nucleus. \\item The sources are spatially coincident with the High Velocity System and thus probably associated with it. They are therefore ULX. \\item Four of the sources have an optical counterpart in the I and R bands (from \\emph{HST} images); two of which are point-like sources and the other two are associated with star-forming regions. \\item In all the cases a single component power law gives satisfactory fits, with spectral index of $\\rm{\\Gamma=}$[1.78-3.51] and an equivalent column density of $\\rm{N_H=[2.05-4.03]\\times 10^{21}cm^{-2}}$. \\item The minimum luminosity is $\\rm{L_X(0.5-7.0 keV)=3.2\\times 10^{39}erg~ s^{-1}}$ (source N1), which is already above the limit of canonical ULX. \\item No variability was detected in the two brightest sources found. \\end{enumerate}  Our results add to the growing evidence that some episodes of rapid star formation lead to the production of ULX.  Young, massive, star clusters may be involved in some, but not all of the sources."
    },
    "0601/astro-ph0601463_arXiv.txt": {
        "abstract": "Strong constraints on the cosmic star formation history (SFH) have recently been established using ultraviolet and far-infrared measurements, refining the results of numerous measurements over the past decade. Taken together, the most recent and robust data indicate a compellingly consistent picture of the SFH out to redshift $z \\approx 6$, with especially tight constraints for $z \\lesssim 1$. We fit these data with simple analytical forms, and derive conservative bands to indicate possible variations from the best fits. Since the $z \\lesssim 1$ SFH data are quite {\\em precise}, we investigate the sequence of assumptions and corrections that together affect the SFH normalisation, to test their {\\em accuracy}, both in this redshift range and beyond. As {\\em lower} limits on this normalisation, we consider the evolution in stellar mass density, metal mass density, and supernova rate density, finding it unlikely that the SFH normalisation is much lower than indicated by our direct fit. Additionally, predictions from the SFH for supernova type~Ia rate densities tentatively suggests delay times of $\\sim 3$\\,Gyr. As a corresponding {\\em upper} limit on the SFH normalisation, we consider the Super-Kamiokande (SK) limit on the electron antineutrino ($\\overline{\\nu}_e$) flux from past core-collapse supernovae, which applies primarily to $z\\lesssim 1$. We find consistency with the SFH only if the neutrino temperatures from SN events are relatively modest. Constraints on the assumed initial mass function (IMF) also become apparent. The traditional Salpeter IMF, assumed for convenience by many authors, is known to be a poor representation at low stellar masses ($\\lesssim 1\\,M_{\\odot}$), and we show that recently favoured IMFs are also constrained. In particular, somewhat shallow, or top-heavy, IMFs may be preferred, although they cannot be too top-heavy. To resolve the outstanding issues, improved data are called for on the supernova rate density evolution, the ranges of stellar masses leading to core-collapse and type~Ia supernovae, and the antineutrino and neutrino backgrounds from core-collapse supernovae. ",
        "introduction": "\\label{int}  In the past few years measurement of the evolution of galaxy luminosity functions at a broad range of wavelengths has rapidly matured. One consequence of this has been the refinement in our understanding of how the space density of the galaxy star formation rate (SFR) evolves \\citep{Lil:96,Mad:96}. In particular the recent results from the Sloan Digital Sky Survey (SDSS), the Galaxy Evolution Explorer (GALEX), and Classifying Objects by Medium-Band Observations (COMBO17) at ultraviolet (UV) wavelengths, and from the Spitzer Space Telescope at far-infrared (FIR) wavelengths, now allows this cosmic star formation history (SFH) to be quite tightly constrained (to within $\\approx30-50\\%$) up to redshifts of $z\\approx1$. Combined with measurements of the SFH at higher redshifts from FIR, sub-millimeter, Balmer line and UV emission, the SFH is reasonably well determined (within a factor of about 3 at $z\\gtrsim1$) up to $z\\approx6$ \\citep[e.g.,][]{Hop:04}.  Additional results from the Super-Kamiokande (SK) particle detector provide a strong limit on the electron antineutrino ($\\overline{\\nu}_e$) flux, 1.2\\,cm$^{-2}$\\,s$^{-1}$ (for $E_{\\nu}>19.3\\,$MeV), originating from supernova type~II (SNII; here and throughout, we assume the inclusion of all core-collapse supernovae: types II, Ib, and Ic) events associated with the SFH \\citep{Mal:03}. This limit on the diffuse supernova neutrino background (DSNB) acts to constrain the normalisation of the SFH \\citep{FK:03,And:04a,Str:04,Str:05}. An exploration of quantities predicted from the SFH, the stellar and metal mass density evolution, and supernova (SN) rate evolution, provides further insight into the allowable normalisation of the SFH \\citep{Str:05}. This series of interconnected physical properties of galaxies and SNe provides an emerging opportunity for determining the level of the SFH normalisation, and the SFH measurements particularly for $z\\lesssim 1$ now have the precision to allow this exploration of their accuracy. Constraining the normalisation of the SFH will support a range of quantitative analyses of galaxy evolution, including the mass-dependence of the SFH \\citep[e.g.,][]{Pap:06, Jun:05, Hea:04}, and the reasons underlying the decline in the SFH to low redshifts \\citep[e.g.,][]{Bell:05}.  The sequence of assumptions explored here starts from observed luminosity density measurements. (1)~To these, dust corrections (where necessary), SFR calibrations and IMF assumptions are applied, to produce corrected SFH measurements. (2)~From the SFH, assumptions about the high-mass IMF fraction that produce SNII lead to predictions for the SNII rate density. This can be compared directly with measurements of this quantity. (3)~Assumptions about the neutrino emission per SNII then give a prediction of the DSNB, for comparison with the SK limit. Since optical SNII can be hidden from observations by dust obscuration, the present SNII rate density measurements may merely be lower limits. In contrast, since neutrinos are unaffected by dust, the DSNB provides an absolute upper limit on the true SNII rate. The consistency that we find between the SFH data and the SK upper limit on the DSNB indicates that there is little room to increase the overall normalisation of the SFH. Thus, this constrains each of the dust corrections, the assumed IMF, and the neutrino emission per supernova. None can be increased significantly without requiring at least one of the others to be reduced below its expected range. Below, we focus especially on the assumed IMF, noting the more detailed treatment of the neutrino emission per supernova by \\citet{Yuk:05}. A parallel set of assumptions to step (2) regarding the generation of SNIa lead to predictions for the SNIa rate density, and this is explored in some detail with tantalising implications regarding the extent of the SNIa delay time.  In \\S\\,\\ref{data} we update the SFH data compilation of \\citet{Hop:04} and address some of the assumptions that affect the normalisation. We identify the best parametric fit to the most robust subset of this data in \\S\\,\\ref{fits}, consistent with the $\\overline{\\nu}_e$ limits from SK. In \\S\\,\\ref{results} we present the results of this fitting in terms of the stellar and metal mass density evolution and the SNII and SNIa rate evolution. The implications for the assumed IMF and SNIa properties are discussed further in \\S\\,\\ref{disc}.  The 737\\footnote{Thanks to Sandhya Rao \\citep{Rao:06} for this terminology.} cosmology is assumed throughout with $H_0=70\\,$km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_M=0.3$, $\\Omega_{\\Lambda}=0.7$ \\citep[e.g.,][]{Spe:03}. ",
        "conclusions": "\\label{disc}  \\subsection{Stellar mass and metal mass densities}  The predictions from the SFH for both $\\rho_*(z)$ and $\\rho_{Z}(z)$ are difficult to analyse, for different reasons. Most predictions of $\\rho_*(z)$ based on SFH measurements seem to be larger than the observed stellar mass density at high redshift, and numerous simulations imply that the measurements might be underestimating the total $\\rho_*(z)$ \\citep[e.g.,][]{Nag:04,Men:04,Som:01,Gra:00}. Indeed it is suggested by \\citet{Dic:03} that additional obscuration added to their maximally old component used in estimating stellar masses could cause arbitrarily large masses to be derived, and it is perhaps not unreasonable to expect about a factor of two larger stellar mass densities as a result of reasonable obscuration levels. This would bring the measurements more into line with the predictions from the SFH. Similar issues may affect the other high redshift measurements of $\\rho_*(z)$, and other issues that have also been raised include incomplete galaxy population sampling and cosmic variance that may affect surveys probing small fields of view \\citep[see discussion in][]{Nag:04}.  At low redshift, the discrepancy between the measurements of $\\rho_*(z)$ and the SFH prediction is more of a concern. A first attempt at resolving this might be to suggest that the measured SFH is too high at $z=0$, and the technique used (combining $\\dot{\\rho}_*$ from FIR and UV estimates) is not accurate. This cannot be the whole solution as it does not address the problem at $z\\approx 1$, where the SFH is dominated by the FIR contribution, and the $\\rho_*(z)$ values inferred from the SFH are similarly higher than the measurements. Another point against this solution is the equivalent diagram in Figure~3 of \\citet{Hop:05}, where (for a Salpeter IMF) a broader region, encompassing the majority of SFH measurements, is used to predict $\\rho_*(z)$ rather than the confidence regions fitted here. Even in that diagram, the $\\rho_*(z)$ measurement at $z=0$ lies at the extreme lower boundary of the SFH prediction. This suggests that there is something more subtle underlying this discrepancy than simple obscuration correction errors in the SFH. A partial solution might be found through the underlying measurement techniques used for the different quantities. The SFH measurements rely on inferring SFRs from the luminosity generated by massive stars, while the $\\rho_*(z)$ measurements come by inferring total stellar masses based on the luminosity of low-mass stars. These are connected in the prediction of $\\rho_*(z)$ from the SFH through the assumed IMF shape, and it is possible that the discrepancy seen in Figure~\\ref{fig:stars} may be reflecting limitations in our understanding of the relative shapes of the low and high mass ends of our assumed IMFs. Although further exploration of this discrepancy is beyond the scope of the present investigation, we refer the reader to \\citet{Far:06} who provide a detailed analysis of this issue, incorporating limits from the total extragalactic background radiation along with the SFH and stellar mass evolution.  Regarding the evolution of the metal mass density, $\\rho_{Z}(z)$, investigation of predictions from the SFH are complicated by the limited number of estimates for this quantity at $z>0$. This is observationally a difficult measurement to make, particularly as much of the metals may exist in an ionised intergalactic medium component. Measurements of the contribution from various components (stellar, dust, gas) have been explored with varying estimates for how much of the metal mass density budget might be contained in different components at high redshift \\citep[e.g.,][]{Dun:03,BLP:05}. Additional discussion of this issue, emphasising the (minimal) contribution from damped Ly$\\alpha$ absorbers, is presented by \\citet{Hop:05}.  \\subsection{Supernovae type~II} \\label{snii}  Almost identical predictions for $\\dot{\\rho}_{\\rm SNII}$ are obtained for both the SalA and SalB IMFs detailed in \\citet{Bal:03}, a consequence of SalA having a slightly higher conversion factor than SalB for transforming $\\dot{\\rho}_*$ from the traditional Salpeter IMF, and a slightly lower conversion factor to transform between $\\dot{\\rho}_*$ and $\\dot{\\rho}_{\\rm SNII}$. The BG IMF result (Figure~\\ref{fig:snrate1}b) appears to be marginally more consistent with both the $\\dot{\\rho}_{\\rm SNII}$ and $\\dot{\\rho}_{\\rm SNIa}$ measurements than the prediction assuming the SalA IMF (Figure~\\ref{fig:snrate1}a) although both are acceptable. The uncertainties affecting the $\\dot{\\rho}_{\\rm SNII}$ measurements are treated in some detail by \\citet{Dah:04}, and the error bars shown are representative of the combination of both statistical and systematic uncertainties. The $z=0$ measurement may be somewhat low, and recent investigations \\citep{Man:03, And:05a} suggest that this point is in fact likely to be higher by a factor of two or three, consistent with the SFH predictions. \\citet{Dah:04} indicate that the main concern in the measurements is the level of obscuration, and that a change from their assumed $E(B-V)=0.15$ of $\\Delta E(B-V)=\\pm0.06$ would alter their measurements by one standard deviation. It is certainly likely that this is the case for the highest redshift point at $z=0.7$, where obscuration is expected to be higher than at lower redshifts \\citep[e.g.,][see also Figure~\\ref{fig:sfh}]{Tak:05,Hop:04}. This would have the effect of raising the $z=0.7$ measurement to be consistent with the predictions from the SFH.  The $\\dot{\\rho}_{\\rm SNII}$ measurements provide a strong lower bound on the normalisation of the SFH. Particularly given that uncertainty regarding obscuration corrections is more likely to raise than lower the $\\dot{\\rho}_{\\rm SNII}$ measurements, the SFH normalisation cannot realistically be much lower than that obtained from assuming the BG IMF (Figure~\\ref{fig:sfhfit}b). The scaling factors for the SFH fits suggest that this provides further evidence against the $T=8\\,$MeV assumption, for both assumed IMFs. Moreover, the $\\dot{\\rho}_{\\rm SNII}$ measurements are unlikely to be affected by sufficient obscuration to support an SFH normalisation much higher than that obtained with the SalA IMF. The level of obscuration to be inferred in order to make the $\\dot{\\rho}_{\\rm SNII}$ measurements lie above the upper edge of the SalA IMF envelope requires $E(B-V)\\approx0.3$ at $z=0.3$ and $E(B-V)\\approx0.5$ at $z=0.7$. These values are quite extreme, corresponding to correction factors of approximately $3.3$ and $6$ respectively, and values this high are not even inferred from the luminosity dependent obscuration corrections of \\citet{Hop:04} for UV data at similar redshifts. Higher SFH normalisations may still be possible, though, and here we return to the assumptions regarding the mass range over which we assume stars become core collapse SNe. If we allow only a mass range of $10<M<30\\,M_{\\odot}$, the $\\dot{\\rho}_{\\rm SNII}$ predictions from the SalA IMF would be lowered by as much as a factor of $0.6$ ($-0.2\\,$dex). Assuming this mass range would place predictions from even the traditional Salpeter IMF at such a level (i.e., $\\approx -0.2\\,$dex from the confidence region in Figure~\\ref{fig:snrate1}a). So, at the expense of an increased lower limit of integration for core collapse SN production, IMFs providing quite high normalisations for the SFH can still be made consistent with the $\\dot{\\rho}_{\\rm SNII}$ measurements.  Another point to be considered here is the issue of a possible neutrino flux generated from a population of massive stellar objects not observable as SNII \\citep{BBM:01,Heg:03,Str:05,Dai:05}. For example, if stars in the mass range $30<M<50\\,M_{\\odot}$ produce the same kind of burst of neutrinos at the end of their life as SNII, but do not become core collapse SNe, instead progressing directly to a black hole or other exotic end, then the neutrino flux inferred from the SFH via an assumed SNII rate will be biased, and the observed SNII rate density could legitimately be $\\approx 10\\%$ lower than that predicted from the SFH using this technique. If even lower mass progenitors become failed supernovae, then this difference could be even larger. Certainly more robust information regarding the connection between the final stages of stellar evolution and neutrino emission would help to refine this type of analysis, and strengthen the implications regarding the favoured SFH normalisation and corresponding IMF.  Our general conclusions here are (1) that the IMF cannot be much more shallow at the high mass end than the BG IMF without predicting values for $\\dot{\\rho}_{\\rm SNII}$ that are too low \\citep[see also][]{Loe:06}; and (2) in the absence of a lower integration limit as high as $\\approx 10\\,M_{\\odot}$ for core collapse SN production, the IMF cannot produce SFH normalisations much higher than does the SalA IMF without predicting values for $\\dot{\\rho}_{\\rm SNII}$ that begin to require quite extreme obscuration corrections. Some slightly less extreme IMFs, while still producing higher SFH normalisations than the SalA IMF, may be allowed if the core collapse SN lower mass cut-off lay between 8 and 10\\,$M_{\\odot}$. This would still seem to be hard to justify, though, as evidence is growing for SN progenitors at these low masses. The SNII analysed by \\citet{Van:06}, for example, (2003gd), favours a progenitor mass $8-9\\,M_{\\odot}$ \\citep[see also][]{Sma:04}. A second SNII (2005cs) also favors a low progenitor mass, $9^{+3}_{-2}\\,M_{\\odot}$ \\citep{MSD:05}. To confirm and refine these IMF constraints, a larger selection of independent $\\dot{\\rho}_{\\rm SNII}$ measurements, spanning a broad range of redshift, would be invaluable. As noted in \\citet{Str:05}, sufficiently precise SNII rate data could be used to directly predict the DSNB flux, independently of assumptions about the SFH and IMF.  \\subsection{Supernovae type~Ia}  The prediction for $\\dot{\\rho}_{\\rm SNIa}$ from the SFH is also particularly intriguing. The assumption of the fixed $t_{Ia}=3\\,$Gyr has the effect of matching the $z\\gtrsim 3$ turnover in the fitted SFH with the apparent decline in $\\dot{\\rho}_{\\rm SNIa}$ seen in the highest redshift measurement from the GOODS sample of \\citet{Dah:04}. It is possible, indeed probable, that this is simply a coincidence as it is a single $\\dot{\\rho}_{\\rm SNIa}$ measurement, with large uncertainties, that is suggestive of the decline, and the turnover in the SFH is driven almost entirely by the $z\\approx 6$ measurement of \\citet{Bun:04}. It is thus still highly possible that the decline in both the SFH and $\\dot{\\rho}_{\\rm SNIa}$ lie at somewhat higher redshift. In particular, recent spectroscopic results from a complete magnitude limited sample \\citep{LeFev:05} suggest that the SFH inferred at $z=3-4$ is up to two or three times higher than that estimated from colour-selected LBGs \\citep{Ste:99}. This may imply that the shape of the SFH is flatter between $2<z<6$ than our current fits suggest. Even more tantalisingly, the higher SFH estimates from the spectroscopic measurements compared to the colour-selected samples at $3<z<4$ suggest that the photometric dropout selected samples at even higher redshift ($z\\approx 6$) may also be underestimating the total SFH. This effect would be compounded by the recent evidence suggesting non-negligible obscuration at high-$z$ \\citep{Cha:05,And:05}, and appears to provide evidence that the expected high redshift decline in the SFH may still not be well established yet. With this in mind, it is nonetheless interesting to note that the robustness of the GOODS measurements, again, has been explored in some detail by \\citet{Dah:04} who are confident of the reliability of this feature in the $\\dot{\\rho}_{\\rm SNIa}$ data. If the turnover in the SFH also turns out to be reliable, the $\\dot{\\rho}_{\\rm SNIa}$ evolution can be used to constrain the delay time for SNIa \\citep[e.g.,][]{For:06}. The predictions for $\\dot{\\rho}_{\\rm SNIa}$ are quite different for $t_{Ia}=1\\,$Gyr (Figure~\\ref{fig:snrate2}a) compared with $t_{Ia}=3\\,$Gyr, and refining the measurements of $\\dot{\\rho}_{\\rm SNIa}$ between $1<z<3$ would be quite revealing. Alternatively, if the SFH turnover does indeed lie at $z>6$ that would imply either that $t_{Ia}>3\\,$Gyr, a result that appears to lie outside the currently favoured range, or that the true turnover in the $\\dot{\\rho}_{\\rm SNIa}$ measurements is also at somewhat higher redshift than the $z\\approx 1.5$ from \\citet{Dah:04}.  To explore a more physically motivated connection between SNIa generation and the underlying stellar populations, \\citet{SB:05} introduced a two component model, dependent on both SFR and stellar mass densities, to derive the SNIa rate density. This has the effect of allowing a contribution to the SNIa rate from old stellar populations where the current SFR may be low, while maintaining a contribution from the currently star forming systems. Taking a logical next step, \\citet{Nei:06} allow a delay time to be incorporated into such a two component model, and find a characteristic delay time $\\tau=3\\,$Gyr (given a Gaussian distribution of delay times), consistent with our simple result above, although their model also includes a contribution from a component proportional to $\\rho_*$. \\citet{Man:06} find a bimodal distribution of delay times, with a ``prompt\" component having a delay time of $\\approx0.1\\,$Gyr. (\\citeauthor{Pin:06} \\citeyear{Pin:06} find evidence that the joint IMF of binary stars favors ``twins\" of nearly equal mass, and suggest that this may provide a natural explanation for these prompt SNIa.) To illustrate the power of a well constrained SFH, we use our best fitting SFH for the BG IMF (from Table~\\ref{tab1}) to explore the available parameters in a \\citet{SB:05} model and, keeping in mind the cautions regarding the high redshift declines in both $\\dot{\\rho}_{\\rm SNIa}$ and $\\dot{\\rho}_*$, do not try to overinterpret our results. We explored the connection between $\\dot{\\rho}_{\\rm SNIa}$, $\\dot{\\rho}_*$ and $\\rho_*$ using the relation \\begin{equation} \\dot{\\rho}_{\\rm SNIa}(t) = A\\dot{\\rho}_*(t-t_{Ia}) + B\\rho_*(t) \\end{equation} with all quantities in the physical units used in this investigation. Performing a simple $\\chi^2$ minimisation, allowing $A$, $B$ and $t_{Ia}$ to vary, we find the best fitting solution favours $A=1.15\\times10^{-3}\\,M_{\\odot}^{-1}$, $B=0\\,M_{\\odot}^{-1}$yr$^{-1}$, and $t_{Ia}=2.7\\,$Gyr. This result is illustrated in Figure~\\ref{fig:snrate2}b, and is clearly driven fairly strongly (as expected intuitively from Figure~\\ref{fig:snrate1}) by the combination of the high-$z$ $\\dot{\\rho}_{\\rm SNIa}$ measurement and the turnover in the SFH driven by the \\citet{Bun:04} measurement. The value of $A=1.15\\times10^{-3}\\,M_{\\odot}^{-1}$ is consistent with the range of $0.001 \\le f_{Ia} \\le 0.0014 \\,M_{\\odot}^{-1}$ assumed for the limits in Figure~\\ref{fig:snrate1}. The $B=0$ result seems to arise from the strong similarity in shape between the SFH and the $\\dot{\\rho}_{\\rm SNIa}(z)$ evolution, while the $\\rho_*(z)$ evolution has the opposite shape (increasing to lower redshifts, rather than decreasing). It seems clear that allowing negative values of B would allow some non-zero solution, although what physical interpretation this would have is not clear. It certainly would not be in the spirit of the model as intended by \\citet{SB:05}. As concluded above, and since $B=0$, our tentative inference appears to favour SNIa delay times close to $t_{Ia}\\approx 3\\,$Gyr, although there are clearly much more sophisticated models, including those allowing for a distribution in delay times, that we have not explored here, and the locations of the apparent downturns in both $\\dot{\\rho}_{\\rm SNIa}$ and $\\dot{\\rho}_*$ clearly have a very strong influence on this result.  \\subsection{Detecting the DSNB}  With the best fitting SFH models explored here, the predictions for the DSNB appear to lie excitingly close to the measured $\\overline{\\nu}_e$ flux limit (Figure~\\ref{fig:sneas}). It is clear that directly observing the DSNB will allow much greater insight into the properties of star formation. Already the DSNB constraint indicates a preferred IMF range and normalisation for the SFH. It also illustrates that stronger constraints on the SFH have implications for understanding the details of both SNII and SNIa production, and the physical basis of neutrino generation by SNII is intimately associated with all these predictions. Being able to detect the DSNB and its energy spectrum will allow a more sophisticated analysis of the detailed connections between all these aspects of star formation and the cosmic SFH.  Methods for increasing the sensitivity of particle detectors to DSNB antineutrinos and neutrinos have been detailed elsewhere, and we briefly reiterate some of these conclusions, to illustrate the potential for detecting the DSNB. \\citet{BV:04} propose loading SK with dissolved gadolinium trichloride to allow tagging of neutron captures, thus significantly lowering backgrounds, in order to directly detect the DSNB $\\overline{\\nu}_e$ spectrum. This provides perhaps the best immediate possibility for measuring the DSNB spectrum shape, which, given the minimum normalisation of the SFH assuming the BG IMF, seems to be lying just below the current flux limit. \\citet{BS:06} describe how the Sudbury Neutrino Observatory (SNO) could improve the current flux limit for DSNB $\\nu_e$ by about three orders of magnitude by coupling the background analysis from SK with the sensitivity to $\\nu_e$ at SNO. Combining information on both DSNB $\\nu_e$ and $\\overline{\\nu}_e$ populations will allow further exciting insight and constraints on SNII neutrino production \\citep[the connection between the two is explored further by][]{Lun:06}."
    },
    "0601/astro-ph0601149_arXiv.txt": {
        "abstract": "We investigate the dependence of the time delays for the large-separation  gravitationally lensed quasar SDSS J1004+4112  on the inner mass profile of the lensing cluster. Adopting the mass model whose innermost density profile is parameterized as $\\rho\\propto r^{-\\alpha}$, we derive a series of mass models which can fit observational data and then compute the probability distribution functions of time delays. We find that larger $\\alpha$ has longer time delays, longer tails at the higher end of the probability distribution, and larger model uncertainties. The ratios of time delays slightly depend on the slope $\\alpha$. Among others, time delays between images C and A (or B) have little dependence on the inner slope, particularly when the time delays are short. The dependence of time delays on $\\alpha$ is well fitted by a linear form, which reflects well-known degeneracy between the mass profile and time delays. We perform a Monte-Carlo simulation to illustrate how well the inner slope can be constrained from measurements of time delays. We find that measurements of more than one time delays result in reasonably tight constraints on the inner slope ($\\sigma_{\\alpha}\\lesssim0.25$), while only one time delay cannot determine the inner slope very well.  Our result indicates that time delays indeed serve as a powerful tool to determine the mass profile, despite the complexity of the lensing cluster. ",
        "introduction": "Recent high-resolution $N$-body simulations in the Cold Dark Matter (CDM) universe have suggested that dark matter halos are described by a universal mass profile (\\cite{navarro96}, \\yearcite{navarro97}, hereafter NFW). Higher-resolution simulations have revealed that the innermost region may have steeper inner slopes than originally suggested and may not be universal (\\cite{fukushige97}, \\yearcite{fukushige01}, \\yearcite{fukushige03}; \\cite{moore99,ghigna00,jingsuto00,power03,fukushige04,hayashi04, diemand04,tasitsiomi04,reed05}). While the accurate value of the inner slope and the universality remain controversial, the existence of deep potential well at the center of dark matter halos in the CDM model appears quite robust. Here the important fact is that the inner profile of dark halos is extremely sensitive to nature of dark matter: For instance, if one introduces self-interacting cross sections of dark matter particles, the central density profile becomes much shallower than expected in the standard CDM universe \\citep{spergel00}. Therefore it serves as a powerful test of collisionless CDM paradigm. Such observational studies to constrain parameters of the density profile have been conducted. For instance, assuming the hydrostatic equilibrium the mass profile of a dark matter halo can be evaluated from the gas density and the temperature profiles of intra-cluster medium (e.g., \\cite{sato00}); weak lensing reconstruction can reproduce the mass profile of a cluster from ellipticities of galaxies behind the cluster \\citep{kaisq93,schneider05}. In particular, rotation curve observations of dark-matter-dominated dwarf and low surface brightness disk galaxies have indicated that they favor mass profiles with a flat density core, being inconsistent with the NFW profile \\citep{salubu00,deblok02,swaters03,gentile04}. However, this apparent discrepancy might be because of bias in gas rotation speed \\citep{hayashi04} or the destruction of central cusps by the formation of a primordial bar \\citep{weinberg02}. Thus, investigations of the halo density profiles in different systems has been one of the most essential issues.  Among others, one of the most promising methods to test the NFW profile would be strong lensing by clusters of galaxies. Gravitational lensing is unique in the sense that it allows one to detect the mass distribution {\\it directly} (\\cite{refsdal64}; \\cite{bourassa73}, \\yearcite{bourassa75}, \\yearcite{bourassa76}; \\cite{wambsganss94,keetonks97,mao98,chiba02, kawano04,kocda04}). In particular, strong lensing by clusters is a indispensable tool to probe the innermost region of dark matter halos, since most baryons in clusters remain hot and diffuse and therefore the density profile of clusters can be well approximated by that of dark halos seen in $N$-body simulations. For instance, giant arcs of background galaxies with different redshifts (Einstein radii) can strongly constrain the mass profiles of clusters \\citep{broadhurst05}. This is because the mass inside the Einstein ring with radius $r_{\\rm E}$ at which arcs build up is proportional to $r_{\\rm E}^2$ and thus multiple arcs with different redshifts determine the masses within the different radii, equivalently the mass slope. Multiply imaged quasars due to clusters of galaxies offer another direct probe of the mass profiles of lensing clusters. Indeed, they have several advantages over giant arcs. First, for lensed quasar systems it is easier to correct selection bias of lensing clusters (i.e., difference between lensed and unlensed cluster populations; see e.g., \\cite{hennawi05}) because of well-known source population of quasars. Second, we can measure time delays for lensed quasars, which serves as additional strong constraints on the mass distribution.  Recently, the first example of such quasar-cluster lens system, SDSS J1004+4112, has been discovered \\citep{inada03,oguri04} in the Sloan Digital Sky Survey (SDSS; \\cite{york00}). It has an unusual separation of $\\sim 15$ arcsec, which is more than twice larger than the second largest lens Q 0957+561 \\citep{walsh79}. The quasar and lensing cluster have redshifts of $1.74$ and $0.68$, respectively. The central region of the lens cluster is dominated by the brightest cluster galaxy. The fifth image of the lensed quasar, which constrains inner mass distribution of the the brightest cluster galaxy strongly, was discovered by \\citet{inada05}.  How did the lens system SDSS J1004+4112 constrain the inner mass distribution of the dark matter halo? \\citet{oguri04} modeled the lens with singular isothermal ellipsoid (SIED) plus NFW, and found that the image configurations and fluxes are reproduced well. More generally, \\citet{willisaha04} adopted a free-form mass reconstruction technique \\citep{sahawi97} and pointed out that the quadruple images alone do not give useful constraints on the inner slope of the dark matter halo. Therefore, we need additional information: One such information comes from arcs of background galaxies which have already observed by \\citet{sharon05}. Arcs at several radii (redshifts) can constrain the inner slope since we can estimate the mass inside a arc from the the arc position and shape. Another information, which is unique for lensed quasar systems, is time delays between each images. It has been shown that the Hubble constant and the slope of the density profile are both sensitive to time delays \\citep{wambsganss94,oguri02,wucknitz02,kochanek02,ogukawa03,kochanek05}, therefore by assuming the value of the global Hubble constant, which is now determined with better than 10\\% accuracy \\citep{freedman01,spergel03}, we will be able to obtain useful information on the mass profile. The image configuration of the system is very similar to that of PG 1115+080 \\citep{weymann80}, however the separations are scaled by a factor of $\\sim 8$ and hence the time delays are longer by a factor of $\\sim 8^2$, which means that the shortest and longest time delays would be $\\sim 10$ days and $\\sim 1000$ days, respectively. However, the complexity of the cluster mass distribution makes the possible range of time delays quite large \\citep{oguri04,willisaha04}. Therefore it is not obvious whether time delays can really give meaningful constraints on the inner profile of the dark matter halo. In this paper, we investigate in detail the capability of measuring time delays and constraining the inner mass slope of the cluster.  This paper is organized as follows. In section \\ref{sec:model} we describe our model to predict time delays for SDSS J1004+4112. Section \\ref{sec:predtij} gives the predictions of time delays and the probability distributions of them. Then, a likelihood analysis is performed in section \\ref{sec:const} to see how time delays constrain the radial slope of the cluster mass. We finally discuss the results and give conclusions in section \\ref{sec:discon}. ",
        "conclusions": "} We have presented predictions of time delays for the giant quadruple lensed quasar SDSS J1004+4112, and investigated the relationship between the time delays and the inner slope of the lensing cluster. We adopt a two-component model in which the brightest cluster galaxy and the cluster are described by the singular isothermal ellipsoid and the generalized NFW profile. We parameterize the inner slope of the cluster component by $\\alpha$ such that $\\rho\\propto r^{-\\alpha}$ in the innermost region. The values of the slope $\\alpha$ we have considered are 0.5, 0.75, 1.0, 1.25, and 1.5, while fixing the scale length to be $40\\arcsec$. We have derived a group of mass models that fit the observables, and calculated the range of predicted time delays for each $\\alpha$. For observables, we have used the data of image positions, fluxes, and the position angle of the galaxy to predict the time delays. We have obtained a set of 100 acceptable models for each $\\alpha$, and constructed the PDFs of the predicted time delays and the ratios between them.  We have found that predicted time delays indeed depend on the inner slope, such that the steeper inner profiles (larger $\\alpha$) predict longer time delays. All the two of them are approximately proportional to each other, but the ratios depend slightly on the inner slope, which suggests that different $\\alpha$ leads to a different structure of the PDFs of multiple time delays. The larger $\\alpha$ models predict longer minimum and maximum time delays, and have longer tails at the higher end of the cumulative probabilities. The time delays $\\Delta t_{\\rm CB}$ and $\\Delta t_{\\rm CA}$ give slightly different feature in short time delay --- the models with different $\\alpha$ show almost same distributions. It is interesting that the model uncertainties resemble the inner slope uncertainty in prediction of time delays, and it would be a clue to investigate the model uncertainties.  To illustrate how well we can constrain the inner slope by adding measurements of time delays, we have calculated the likelihood function for $\\alpha$. Figure \\ref{allike} has shown that the slope $\\alpha$  is constrained weakly with the measurement of one time delay. However, we also found that multiple time delays will results in reasonably strong constraints ($\\sigma_{\\alpha}\\lesssim0.25$) on the inner slope. We note that we assumed observational errors of time delays to be 10\\%, which may be too conservative. For instance, Q 0957+561 has a time delay of $\\sim 420$ days and the error is measured to be 1\\% \\citep{kundic97,oscoz01}. If this level of errors can be achieved for SDSS J1004+4112, we will be able to determine the inner slope more tightly.  Our model predictions offer useful guidance for photometric monitoring of this lens system to determine the time delays. We found that $95\\%$ of models of $\\alpha=0.5 - 1.5$ have the predicted time delays of $\\Delta t_{\\rm BA}\\lesssim 28$, $\\Delta t_{\\rm CB}\\lesssim 1400$, and $\\Delta t_{\\rm CD}\\lesssim 3700$ in units of $h_{70}^{-1}$ days. The first time delays is similar to those in any other lensed quasars, and therefore it can be measured easily. Indeed, the preliminary detection of the time delay between B and A has been made to be $\\sim 25$ days (C. S. Kochanek, private communication). The second shortest time delay $\\Delta t_{\\rm CB}$ is $\\sim 30-40$ times larger than  $\\Delta t_{\\rm BA}$, thus assuming the measured $\\Delta t_{\\rm BA}$ we predict $\\Delta t_{\\rm CB}=750-1000$ days. This is somewhat longer than observed time delays in any other systems, but is not impossible to measure. The longest time delays need monitoring more than $\\sim 10$ years, making the measurement quite challenging.  The mass profile of the lensing cluster can be constrained much better if we combine the time delay measurements with other observations. For instance we may add constraints from multiply imaged galaxies behind the cluster \\citep{sharon05}. The multiple arcs that have different redshifts (and therefore different Einstein ring radii) are an independent tool to constrain the cluster mass profile, especially the radial mass slope $\\alpha$. Other such observations includes X-ray measurements of intra-cluster medium and the measurement of the velocity dispersion of the brightest cluster galaxy (eq. [\\ref{eq:bv}]). By combining these complementary information, we will be able to reveal the detailed distribution of the dark matter of the lensing cluster in a robust manner.  \\bigskip We are grateful to Takahiko Matsubara, Naohisa Inada, and Chris Kochanek for discussions."
    },
    "0601/astro-ph0601713_arXiv.txt": {
        "abstract": "We present CCD photometry, light curve and time series analysis of the classical nova V382 Vel (N Vel 1999). The source was observed for 2 nights in 2000, 21 nights in 2001 and 7 nights in 2002 using clear filters. We report the detection of a distinct period in the light curve of the nova P=0.146126(18) d (3.5 h). The period is evident in all data sets, and we interpret it as the binary period of the system. We also measured an increase in the amplitude modulation of the optical light (in magnitude) by more than 55$\\%$ from 2000 to 2001 and about 64$\\%$ from 2001 to 2002. The pulse profiles in 2001 show deviations from a pure sinusoidal shape which progressively become more  sinusoidal by 2002. The main cause of the variations in 2001 and 2002 can be explained with the occultation of the accretion disk by the secondary star. We interpret the observed deviations from a pure sinusoidal shape as additional flux resulting from the aspect variations of the irradiated face of the secondary star. ",
        "introduction": "A classical nova outburst arises from the explosive ignition of accreted matter (i.e., thermo-nuclear runaway, TNR) in a cataclysmic binary system where a Roche-lobe filling secondary transfers hydrogen rich material typically via an accretion disk onto the white dwarf (WD) primary. During the outburst, the envelope of the WD expands to $\\sim$ 100 R$_{\\odot}$ and 10$^{-7}$ to 10$^{-3}$\\ M$_{\\odot}$ of material that exceeds the escape velocity is expelled from the system (Shara 1989; Warner 1995; Starrfield 2002). The outburst stage usually lasts from a few months to several years and finally, the system returns to a quiescent state after the outburst.  Nova Vel 1999 (V382 Vel) was discovered on 1999, May 20.6 UT (Lee at al. 1999). It is believed to be an outburst on an ONeMg white dwarf because of the strong Ne line emission detected in the optical wavelengths (Woodward et al. 1999). Spectroscopic observations from 5 to 498 days (after the eruption) indicate that the nova belongs to a broad Fe II spectroscopic class with an absolute magnitude at the optical maximum M$_v$$\\sim$-8.9, t$_2$=6 d and t$_3$=10 d (Della Valle et al. 2002). The early evolution of the nova shows an iron curtain phase and P-Cygni profiles on all the important resonance lines with expansion velocities between 2000-4000 km s$^{-1}$ (see Shore et al. 2003 and Della Valle et al. 2002 for the optical and Burwitz et al. 2002 for the X-ray wavelengths). Fragmentation in the ejecta is apparent and the ejected mass is calculated to be 4-5$\\times$10$^{-4}$ M$_{\\odot}$. The nova is found to be  enhanced in Ne/Ne$_{\\odot}$=17($\\pm$3), N/N$_{\\odot}$=17($\\pm$4), C/C$_{\\odot}$=0.6($\\pm$0.3), Al/Al$_{\\odot}$=21($\\pm$2), and Mg/Mg$_{\\odot}$=2.6($\\pm$0.1) (Shore et al. 2003). The distance to the nova is estimated as 1.7-2.5 kpc (Della Valle et al. 2002; Shore et al. 2003). In addition, Nova Vel 1999 appeared as a Super Soft X-ray Source (Orio et al. 2002; Ness et al. 2005) and the hard X-ray emission was detected originating from the shock heated material in the shell (Orio et al. 2001; Mukai $\\&$ Ishida 2001; Burwitz et al. 2002; Ness et al. 2005)  Bos et al. (2002) discovered  a 3.5-h periodicity; 0.14615(1) d together with its second harmonic in the optical light curve of the nova. In this paper, we will elaborate that report, discuss possible physical mechanisms for the variation and try to understand the characteristics of the long term light curve and the detected period. ",
        "conclusions": "We have detected  a consistent and coherent periodicity P=0.146126(18) in the optical light curve of the classical nova V382 Vel using the data obtained in the years 2000 (two nights), 2001 (21 nights) and 2002 (seven nights). The observations were conducted at the Wharemaru Observatory (2000; 0.25m telescope) and Molehill Observatory (2000-2001; 0.3m telescope) in New Zealand. We interpret this variation as the binary period of the underlying system. We conclude that the cause of variations is the occultation of an accretion disk by the secondary star. We detect increasing modulation amplitudes as the nova itself cools off in time  in the optical wavelengths, and the contribution of light from the occulted disk into the total light curve increases enhancing the amount of variation in the light curve. The increase in the variation amplitude from 2000 ($\\Delta$m$<$0.007) to 2001 ($\\Delta$m=.014$\\pm$0.003) is larger than 55$\\%$ and from 2001 to 2002 ($\\Delta$m=0.023$\\pm$0.008), it is 64$\\%$. We also observe variation of the shape and depth of the mean light curves particularly in 2001. We favor a scenario where variations due to the aspect changes of the irradiated secondary star intervenes with the optical light curve at phases 0.10-0.15, 0.5-0.6 and 0.8-0.9.  The soft X-ray radiation from the WDs in the outburst stage of classical novae could trigger irradiation induced high mass transfer resulting in re-establishment of the accretion disk. Irradiation induced mass transfer cycles occur in compact binaries if the donor star has a shallow convective envelope (fast thermal time scale), or the system has mass transfer driven on a long dynamical timescale, or the photosphere scale height is small (Buning $\\&$ Ritter 2004; Ritter, Zhang, $\\&$ Kolb 2000). Once the nova is a soft X-ray source it has the potential to irradiate its companion initiating mass transfer (stable/limit cycle) if the basic conditions described above are also met depending on other important parameters like the masses of the two stars, the binary separation, etc. The systems which have orbital periods above 3 hrs tend to form mass transfer limit cycles whereas the systems with periods below 2.5 hrs show stable mass transfer unless the deriving rate is no larger than a few times the gravitational braking rate (Buning $\\&$ Ritter 2004). The H burning phase of classical novae outbursts will be interesting laboratories to study irradiation induced mass transfer phenomenon for CVs. In addition, the hibernation scenario for novae suggests that classical novae remain bright for a few centuries after the eruption because of irradiation induced mass transfer (Shara et al. 1986; Prialnik $\\&$ Shara 1986). Nova systems are known to show irradiation effects and/or superhumps in their light curves after the optical decline (see Retter $\\&$ Naylor 2000).  The three years of monitoring observations  of the classical nova V382 Vel in the optical wavelengths have revealed several important facts on the evolution of classical novae during the outburst stage. It shows that accretion is established and a disk is present early in the evolution (as early as 645 days after the outburst). It also shows evidence that the irradiation of the secondary stars by the hot stellar remnants (WDs) could be a common phenomenon that is apparent in the outburst light curves of the first few years till after the stellar remnants cool off of the X-ray wavelengths. Only few novae have been detected during the SSS phase in the X-rays that lasts a few weeks up to several years after the outburst (eg. less than 9-10 years: Orio 2004) and such observations are difficult to obtain due to the tight schedules of the X-ray satellites. The evolution of the light variations as a result of the aspect changes of the irradiated secondary stars can yield additional indirect proof of the evolution of the hot white dwarfs after the outburst and complement the X-ray observations if systematic long term ground-based observations can be maintained in the optical wavelengths."
    },
    "0601/astro-ph0601239_arXiv.txt": {
        "abstract": "{The star-to-star scatter in lithium abundances observed among otherwise similar stars in the solar-age open cluster M~67 is one of the most puzzling results in the context of the so called ``lithium problem\". Among other explanations, the hypothesis has been proposed that the dispersion in Li is due to star-to-star differences in Fe or other element abundances which are predicted to affect Li depletion. } {The primary goal of this study is the determination of the metallicity ([Fe/H]), $\\alpha$- and Fe-peak abundances in a sample of Li-poor and Li-rich stars belonging to M~67, in order to test this hypothesis. By comparison with previous studies, the present investigation also allows us to check for intrinsic differences in the abundances of evolved and unevolved cluster stars and to draw more secure conclusions on the abundance pattern of this cluster.} {We have carried out an analysis of high resolution UVES/VLT spectra of eight unevolved and two slightly evolved cluster members using MOOG and measured equivalent widths. For all the stars we have determined [Fe/H] and element abundances for O, Na, Mg, Al, Si, Ca, Ti, Cr and Ni.} {We find an average metallicity [Fe/H]~$=0.03\\pm0.01$, in very good agreement with previous determinations. All the [X/Fe] abundance ratios are very close to solar. The star-to-star scatter in [Fe/H] and [X/Fe] ratios for all elements, including oxygen, is lower than 0.05~dex, implying that the large dispersion in lithium among cluster stars is not due to differences in these element abundances. We also find that, when using a homogeneous scale, the abundance pattern of unevolved stars in our sample is very similar to that of evolved stars, suggesting that, at least in this cluster, RGB and clump stars have not undergone any chemical processing. Finally, our results show that M~67 has a chemical composition that is representative of the solar neighborhood.} {} ",
        "introduction": "``Classical\" or ``standard\" models of stellar evolution include convection only as internal mixing process and neglect more complex physical processes like diffusion, magnetic fields, rotation; these models predict that solar-type stars should not deplete lithium (Li)  while on the main sequence (MS), since the base of the convective zone does not reach the layers in the stellar interior where the temperature is high enough for Li reactions. Furthermore, according to standard models, stars with the same age, mass, and chemical composition should undergo the same amount of Li depletion. In sharp contrast with these predictions, not only is now well established on observational grounds that solar-type stars, including the Sun, do deplete a significant amount of Li during the MS phases, but Li depletion can be different for similar stars (e.g., Randich \\cite{r_cast} and references therein). The factor of about 10 star-to-star scatter in Li abundances observed among otherwise identical F- and G-type members of the 4.5~Gyr old cluster M~67 indeed represents one of the most puzzling results in the context of the so-called ``Lithium problem\" (Spite et al. \\cite{spi87}; Garcia L\\'opez et al. \\cite{gar88}; Pasquini et al. \\cite{pas97}; Jones et al. \\cite {jon99}).  A large dispersion in Li abundances is present among old field stars (e.g., Pallavicini et al. \\cite{pal87}; Pasquini et al. \\cite{pas94}). On the other hand, whereas no dispersion is seen in the 600~Myr old Hyades, in the three  2~Gyr old clusters IC~4651, NGC~3680 and NGC~752 (Randich et al. \\cite {r00}; Sestito et al. \\cite{ses_752}), nor in 6~Gyr old NGC~188 (Randich et al \\cite{n188}), preliminary results of Li measurements among large samples of stars in the 2~Gyr, metal rich NGC~6253 and in the very old Collinder~261 suggest that they could be characterized by some amount of scatter (Pallavicini et al. \\cite{pall_cast}; Randich \\cite{r_cast}). In other words, the appearance of a dispersion seems to depend more on the characteristics of the cluster than on age.  It has been shown that non membership and/or binarity are not the reasons for the scatter in M~67 (Pasquini et al. \\cite{pas97}); also, the effects of chromospheric activity on the formation of the Li~{\\sc i} line, which are proposed as a possible explanation for the dispersion in young clusters (Jeffries \\cite{jef_06} and references therein), are unlikely to play a role, given that at the old age of M~67 the level of chromospheric activity should be low enough not to affect the Li~{\\sc i} line. Hence, the scatter is most likely intrinsic and, under the very reasonable assumption that cluster stars were all born with the same Li content, it must reflect different amounts of Li depletion.  Different possibilities were proposed to explain the existence of the spread (Randich \\cite{r_cast}); among them, it has been suggested that M~67 members do not have an identical chemical composition and that the scatter in Li may reflect a scatter in iron content or, more in general, in heavy element composition (Garcia L\\'opez et al. \\cite{gar88}; Piau et al. \\cite{piau03}).  As well known, chemical composition affects stellar opacities and thus internal structure, mixing processes (both standard and non standard ones) and, in principle, Li depletion. The effect of variations of the chemical composition on pre-main sequence Li depletion has been theoretically investigated in different studies (e.g, Swenson et al. \\cite{sf_92}; Piau \\& Turck-Chieze \\cite{ptc}; Sestito et al. \\cite{ses_06}) and all of them agree in that even relatively small changes of the mass fraction of elements critical for the opacity can have significant effects on the amount of Li depletion. Similarly, Piau et al. (\\cite{piau03}) have shown that variations in CNO abundances can change the amount of Li depletion during the MS phases and suggested that star-to-star differences in these element abundances or, more in general in $\\alpha$ and Fe-peak abundances, could explain the observed scatter in M~67. More quantitatively, they found that a difference of 0.05~dex in [CNO/Fe] would result in a difference in $\\log$~n(Li) of $\\sim 0.5$~dex (i.e, smaller than the observed spread), implying that larger differences in CNO abundances are needed to explain the whole spread.  Garcia L\\'opez et al. (\\cite{gar88}) did not find any difference in the overall metallicity of Li-poor and Li-rich stars and ruled out that different [Fe/H] values within M~67 could be the reason for the scatter in Li. Here we extend their work to elements other than iron and to a larger sample of stars, to investigate {\\it a)} whether a large ($> 0.05$~dex) star-to-star scatter in heavy elements is present among unevolved cluster members; and, in case such a spread is detected, {\\it b)} whether it is related to the dispersion in Li.  In a more general context, the determination of the metallicity and chemical composition of open clusters covering a large interval of ages, metallicities, and Galactocentric distances is a critical tool to investigate the formation and evolution of the Galactic disk (Friel \\cite{friel} and references therein). M~67 is one of the closest and best studied old open clusters; nevertheless, relatively few studies focused on the determination of its chemical composition (Tautvai\\v sien\\. e et al.  \\cite{tau00} and references therein), and most of them are based on few and/or evolved stars that may have undergone chemical processing. The present work represents the first abundance study of M~67 based on a rather large sample of unevolved (or slightly evolved) cluster members, allowing us to check whether intrinsic differences exist between abundances of unevolved and evolved stars and to draw more secure conclusions on the abundance pattern of this cluster. The paper is structured as follows: in Sect.~2 the sample and the observations are described, while the abundance analysis is presented in Sect.~3. The results, together with a discussion of internal and systematic errors, and a comparison with findings from previous studies are presented in Sect.~4. Finally, a discussion of the results and conclusions are given in Sect.~5. ",
        "conclusions": "\\subsection{Star-to-star scatter and Li} \\label{scatter} As mentioned in Sect.~1, the investigation of the presence (or lack thereof) of a significant dispersion in $\\alpha$ and Fe-peak element abundances among cluster stars, that could possibly explain the large star-to-star scatter in Li abundances, was the main motivation for the present study.  Figures~\\ref{figure2} and \\ref{figure3} together with Tables~\\ref{sample}, \\ref{ratios}, \\ref{oxygen} and \\ref{errors} clearly indicate that the standard deviations $\\sigma$ of [Fe/H] and [X/Fe] ratios for the whole sample are comparable or even smaller than measurement uncertainties of individual stars. This on the one hand implies that our estimate of internal errors may be somewhat conservative; on the other hand, and most important, $\\sigma$ values represent upper limits to the dispersion in [Fe/H] and [X/Fe] ratios, suggesting that we can exclude the presence of a scatter in [Fe/H] and [X/Fe] ratios at a level larger than $\\sim 0.05$~dex. \\setcounter{table}{7} \\begin{table*} \\caption{Reanalysis of Tautvai\\v sien\\. e et al. (\\cite{tau00}) data. The subscript ``our\" indicates the values obtained through our analysis, while the subscript ``TETI\" indicates their original values.}\\label{taut_n} \\begin{tabular}{cccccccc}\\hline Ion & \\multicolumn{2}{c}{F84} & \\multicolumn{2}{c}{F141} & \\multicolumn{2}{c}{F151} & average$_{\\rm our}$\\\\ & [X/Fe]$_{\\rm TETI}$ & [X/Fe]$_{\\rm our}$ & [X/Fe]$_{\\rm TETI}$ & [X/Fe]$_{\\rm our}$ &[X/Fe]$_{\\rm TETI}$ & [X/Fe]$_{\\rm our}$ &\\\\ & & & & & & &\\\\ Na~{\\sc i} & 0.21 & 0.11 & 0.19 & 0.12 & 0.21 & 0.12 & 0.12\\\\ Al~{\\sc i} & 0.08 & $-0.04$ & 0.12 & $-$0.03 & 0.09 & $-0.01$ & $-$0.03 \\\\ Si~{\\sc i} & 0.09 & $-$0.03 & 0.06 & $-$0.03 & 0.08 & $-0.04$ & $-$0.03 \\\\ \\hline \\end{tabular} \\end{table*} In order to investigate more in detail the relationship between heavy element abundances and Li dispersion, sample stars shown in Fig.~\\ref{figure2} are denoted with different symbols according to whether they are Li-rich or Li-poor. Analogously, we show in Fig.~\\ref{figure4} [X/Fe] vs. [Fe/H] for our sample stars with Li-rich and Li-poor objects indicated by different symbols. We define as Li-rich/poor stars those lying on the upper/lower envelope of the $\\log$~n(Li) vs. \\teff~distribution of M~67 (see Pasquini et al. \\cite{pas97}; Jones et al. \\cite{jon99}). Excluding stars S1034 and S1239 that are already on the subgiant branch and have undergone a certain amount of post-MS Li dilution (Balachandran \\cite{bala95}), three and five of the remaining sample stars fall in the Li-poor/rich category (open and filled circles, respectively). The figure shows that no systematic difference between the abundance pattern of Li-rich and Li-poor stars is present, in the sense that Li-poor stars are not systematically more metal-rich/poor and/or have larger/lower [X/Fe] ratios than the other stars. In particular the two stars S1252 and S1256 (similar temperatures, but a factor of almost 10 difference in Li) have very similar [Fe/H] and similar (almost identical in some cases) abundances of $\\alpha$-elements, including oxygen, as well as of Fe-peak elements. As for the other three stars with similar temperatures, but different Li abundances (S2205 and S988 both Li-poor and S994 Li-rich), they have similar [X/Fe] ratios for most elements ---again including O--- with the exception of Si, Ca, and Cr. [Si/Fe], [Ca/Fe] and [Cr/Fe] are somewhat higher in S988 than in the other two stars. However, not only the difference is within errors, but the largest difference in heavy element abundances is seen between the two Li-poor stars themselves, with S988 being on average more metal rich than S2205. S994 has [X/Fe] values intermediate between those of S2205 and S988.  We conclude that, based on our sample of M~67 members, it seems very unlikely that dispersion of a factor about 10 in Li abundances measured among otherwise similar cluster stars is due to differences larger than $\\sim 0.05$~dex in heavy element abundances. We note in particular that no dispersion in O -one of the most critical elements for stellar opacities and thus Li destruction- is present among our sample stars.  Our result on the lack of a dispersion in heavy element abundances leaves early depletion due to angular momentum loss and transport as the most likely explanation for the dispersion (Jones et al. \\cite{jon99}). In this scenario, stars with different initial rotational velocities would undergo different amounts of Li depletion, resulting in a dispersion in Li at the age of M~67. This scenario however encounters two main difficulties: first, models including mixing due to angular momentum transport predict a correlated Li and Be depletion, which is instead not seen in M~67 (Randich et al. \\cite{R02}). Second, since observations of young cluster show that they all have similar rotation distributions, one would expect that all old clusters should be characterized by a spread in Li, which is instead not the case. Further studies are most obviously warranted. \\subsection{M~67 in the disk} In Fig.~\\ref{disk} we compare the average [X/Fe] vs. [Fe/H] pattern of M~67 with the sample of field stars thin and thick disks from Bensby et al. (\\cite{bfl03}) and Bensby et al. (\\cite{bensby04}) for O. With the caveat that the comparison of M~67 with field stars can be made only on qualitative grounds, since abundances are on different scales and systematic offsets might be present, the figure indicates that the average [X/Fe] ratios for M~67 very well fit into the general trend of field stars. The $\\alpha$-elements have solar ratios (or slightly below solar in the case of Ti) and none of them is significant over- or under-abundant with respect to field stars. [O/Fe] is also solar and very well fits within the distribution of field stars. Na, which we find to be somewhat enhanced with respect to the Sun, lies on the upper envelope of the distribution of field stars, but is still consistent with it. Similarly, Al is slightly below the solar ratio and the lower envelope of field stars, but, given the uncertainties, consistent with it. In other words, our measurements confirm that M~67 has a ``normal\" abundance pattern for the solar neighborhood and it is very much similar to the Sun. As a final remark, we wish to stress that, given its age and global abundance pattern, M~67 represents one of the most promising clusters where to look for true solar-analogs."
    },
    "0601/astro-ph0601098.txt": {
        "abstract": "{We study the polarization of the SiO maser emission in a representative sample of evolved stars in order to derive an estimate of the strength of the magnetic field, and thus determine the influence of this magnetic field on evolved stars. We made simultaneous spectroscopic measurements of the 4 Stokes parameters, from which we derived the circular and linear polarization levels. The observations were made with the IF polarimeter installed at the IRAM 30m telescope. A discussion of the existing SiO maser models is developed in the light of our observations. Under the Zeeman splitting hypothesis, we derive an estimate of the strength of the magnetic field. The averaged magnetic field varies between 0 and 20 Gauss, with a mean value of 3.5 Gauss, and follows a $1/r$ law throughout the circumstellar envelope. As a consequence, the magnetic field may play the role of a shaping, or perhaps collimating agent of the circumstellar envelopes in evolved objects.} ",
        "introduction": "The prodigious mass loss observed in numerous and widespread evolved stars make these objects the main recycling agents of the interstellar medium, and thus one of the most important objects in the Universe. Even though our knowledge of evolved stars has considerably improved over recent years, some of their main characteristics remain insufficiently understood (see the review by Herwig 2003):  which mechanisms are responsible for their drastic change of geometry when evolving to the Planetary Nebula (hereafter PN) stage ? What is powering so efficiently the mass loss and could the magnetic field play a major role ?  Important information about the physics and chemistry prevailing in the circumstellar envelope  (hereafter {\\em CSE}) of evolved stars can be retrieved from radiowave line emission of molecules, specially from maser emission (see the review by Bujarrabal 2003). These envelopes can be probed at different depths through the study of three masing molecules, OH, \\water and SiO. Our current knowledge indicates that: \\begin{itemize} \\item OH radiation traces the outer part of the envelope, at 1000-10000 AU from the central star; \\item \\water molecules are located at intermediate distances, i.e. a few 100 AU; \\item SiO maser emission comes from the inner regions of the envelope, between 5 to 10 AU (a few stellar radii R$_{\\star}$). \\end{itemize}  The SiO maser emission is produced in small gas cells, and  is known to be polarized. The polarization (circular or linear) and angle of the emission can be measured and thus improve our knowledge of these objects. In addition, studying the maser polarization can shed light on the maser theory itself. As explained further in this paper, several uncertainties in the theory make data interpretation often difficult, and new observational data are helpful. One of the most interesting quantities that can be derived from polarization measurements is the stellar magnetic field. According to theory (e.g. Elitzur 1996 or 2002), measurement of the maser radiation polarization can lead to an estimation of the magnetic field strength B and can reveal its spatial structure (via interferometric observations). In single dish observations, all of the maser components get smeared within the beam and only the mean value of B along the line of sight ($B_{//}$) can be derived. Only SiO masers are capable of tracing the magnetic field as close as $\\sim$ 5 AU from the central star. But if SiO masers are to be used as a B-field tracer, we first need to give evidence that SiO masers are reliable B-field tracers. This requires more detailed theories than available today. Nevertheless, we tentatively derive in this work the field strength in the CSE inner layers of several evolved stars.  Research on astronomical masers polarization is very active but is made difficult both by the lack of specific instrumental facilities and by the excitation and propagation of the masers themselves. Until now, numerous polarimetric observations of OH masers have been done, several of \\water masers, but few of SiO maser emission. Few SiO polarimetric observations have been done with VLBI giving the very first images of the magnetic field  in some objects (e.g. Kemball \\& Diamond 1997 in TX Cam). Most of the early SiO studies were done in linear polarization. The first complete SiO polarimetric observations were performed by Johnson \\& Clark (1975), then by Troland \\etal (1979); emission was found to be typically 15-30 \\% linearly polarized and to exhibit no circular polarization. Barvainis, McIntosh \\& Predmore (1987), and McIntosh \\etal (1989) measured circular polarization of $1-9$ \\% in several stars. Circular (0-4 \\%) and linear (3.7-9.7 \\%) polarizations were measured in VY CMa by McIntosh, Predmore \\& Patel (1994). Later, Kemball \\& Diamond (1997) made the first image of the magnetic field in the atmosphere of TX Cam, measuring a circular polarization level of 5 \\% with some features showing polarization up to 30-40 \\%.  It must be stressed that SiO, as H$_2$O, is a non-paramagnetic species. Zeeman splitting exists but the sublevels overlap; the effect is thus undetectable and hence only net polarization can be used to trace the magnetic field. The current status of our knowledge on the magnetic field strength can be summarized as follows: \\begin{itemize} \\item between 1000-10000 AU, $B_{//}\\sim 5-20$ mG (OH masers, e.g. Kemball \\& Diamond 1997, Szymczak \\& Cohen 1997); \\item at a few 100 AU from the star, $B_{//} \\sim$ a few 100 mG (\\water masers, e.g. Vlemmings, Diamond \\& van Langevelde 2001, Vlemmings, van Langevelde \\& Diamond 2005); \\item at 5-10 AU, $B_{//} \\sim 5-10$ G (SiO masers; Kemball \\& Diamond 1997, in TX Cam). \\end{itemize}  The main purpose of this work is to measure and analyze the SiO maser polarization in terms of magnetic field strength in a representative sample of evolved stars. Our observations are presented in Section 2; they include simultaneous spectroscopic measurements of the 4 Stokes parameters. The results for individual stars are discussed in Section 3. In Section 4, we compare our data with predictions from existing SiO maser models and initiate a discussion on the validity of these models. Within the limitations of one of these models we derive the magnetic field strength and try to determine the role of the magnetic field. More broadly, a summary of the magnetic field topic in evolved stars is also given in Section 4. In Section 5 we give some concluding remarks. ",
        "conclusions": "We have made a study of the SiO maser polarization in a representative sample of evolved stars, simultaneously measuring, for the first time, the 4 Stokes parameters. From our measurements we derive the circular and linear polarization levels and shows that, due to the beam averaging of our polarization measurements, we cannot firmly discriminate between the two dominant theories of SiO maser emission. In particular, VLBI observations of our source sample are absolutely necessary to distinguish between Zeeman or non-Zeeman theories. Nevertheless, the magnetic field strength was derived assuming Elitzur's model. $B_{\\\\}$ varies between 0 and 20 Gauss, with a mean value of 3.5 G. As a consequence, we suggest that the magnetic field plays a significant role in the evolution of these objects. Within the frame of the Zeeman theory the magnetic field could shape or even collimate the gas layers surrounding the AGB objects. Emission from Mira-type objects clearly tends to have a higher linear  ( ${<p_L>}_{Mira} \\simeq 30$\\%, ${<p_L>}_{SR} \\simeq 11$\\%) and circular polarization  (${<p_C>}_{Mira} \\simeq 9$\\%, ${<p_C>}_{SR} \\simeq 5$\\%). Basically, if there is a real correlation between $p_C$ and the strength of the magnetic field, this trend may indicate that the magnetic field may be  stronger in Mira objects than in Semi-Regular variables (at least in the inner layers of the circumstellar envelope).  To better understand the mechanisms at work with the magnetic field, complementary studies have to be conducted and in particular the presence of a companion has to be investigated in a large sample of objects. Of course VLBI maps of the magnetic field in these stars are essential. Another important objective is to investigate the evolution of the magnetic field and its influence during the transition from the AGB star phase to the PN stage."
    },
    "0601/astro-ph0601707_arXiv.txt": {
        "abstract": "A toy model for gamma-ray burst supernovae (GRB-SNe) is discussed by considering the effects of screw instability of magnetic field in black hole (BH) magnetosphere. The screw instability in the Blandford-Znajek (BZ) process (henceforth SIBZ) can coexist with the screw instability in the magnetic coupling (MC) process (henceforth SIMC). It turns out that both SIBZ and SIMC occur inevitably, provided that the following parameters are greater than some critical values, i.e., (i) the BH spin, (ii) the power-law index describing the magnetic field at the disk, and (iii) the vertical height of the astrophysical load above the equatorial plane of the rotating BH. The features of several GRBs are well fitted. In our model the durations of the long GRBs depend on the evolve time of the half-opening angle. A small fraction of energy is extracted from the BH via the BZ process to power a GRB, while a large fraction of energy is extracted from the BH via the MC process to power an associated supernova. In addition, the variability time scales of tens of msec in the light curves of the GRBs are fitted by two successive flares due to SIBZ. ",
        "introduction": "Recently, observations and theoretical considerations have linked long-duration GRBs with ultra-bright Type Ib/c supernovae (SNe; Galama et al. 1998, 2000; Bloom et al. 1999). The first candidate was provided by SN 1998bw and GRB 980425, and the recent HETE-II burst GRB 030329 has greatly enhanced the confidence in this association (Stanek et al. 2003; Hjorth et al. 2003). Extremely high energy released in very short time scale suggests that GRBs involve the formation of a black hole (BH) via a catastrophic stellar collapse event or possibly a neutron star merger, implying that an inner engine could be built on an accreting BH.  Among a variety of mechanisms of powering GRBs the Blandford-Znajek (BZ) process (Blandford {\\&} Znajek, 1977) has its unique advantage in providing ``clean'' (free of baryonic contamination) energy by extracting rotating energy from a BH and transferring it in the form of Poynting flow in the outgoing energy flux (Lee et al. 2000, hereafter Lee00; Li 2000c, hereafter Li00).  Not long ago Brown et al. (2000, hereafter B00) worked out a specific scenario for a GRB-SN connection. They argued that the GRB is powered by the BZ process, and the SN is powered by the magnetic coupling (MC) process, which is regarded as one of the variants of the BZ process (Blandford 1999; van Putten 1999; Li 2000b, 2002; Wang, Xiao {\\&} Lei 2002, hereafter W02). It is shown in B00 that about $10^{53}ergs$ are available to power both a GRB and a SN. However, they failed to distinguish the fractions of the energy for these two objects.  More recently, van Putten and his collaborators (van Putten 2001; van Putten {\\&} Levinson 2003, hereafter P03) worked out a poloidal topology for the open and closed magnetic field lines, in which the separatrix on the horizon is defined by a finite half-opening angle. The duration of a GRB is set by the lifetime of the rapid spin of the BH. It is found that GRBs and SNe are powered by a small fraction of the BH spin energy. This result is consistent with observations, i.e., duration of GRBs of tens of seconds, true GRB energies distributed around $5\\times 10^{50}ergs$ (Frail et al. 2001), and aspherical SNe kinetic energies of $2\\times 10^{51}ergs$ (Hoflich et al. 1999).  Very recently, Lei et al. (2005, hereafter Lei05) proposed a scenario for GRBs in type Ib/c SNe invoking the coexistence of the BZ and MC processes. In Lei05 the GRB is powered by the BZ process, and the associated SN is powered by the MC process. The overall time scale of the GRB is fitted by the duration of the open magnetic flux on the horizon.  Besides the features of high energy released in very short durations most GRBs are highly variable, showing very rapid variations in flux on a time scale much shorter than the overall duration of the burst. Variability on a time scale of milliseconds has been observed in some long bursts (Norris et al. 1996; McBreen \\textit{et al.}, 2001; Nakar and Piran, 2002). Unfortunately, the origin of the variations in the fluxes of GRBs remains unclear. In this paper we intend to discuss the mechanism for producing the variations in the fluxes of GRBs by virtue of the screw instability in BH magnetosphere.  It is well known that the magnetic field configurations with both poloidal and toroidal components can be screw-unstable. According to the Kruskal-Shafranov criterion the screw instability will occur, if the toroidal magnetic field becomes so strong that the magnetic field line turns around itself once or more (Kadomtsev 1966; Bateman 1978).  Some authors have addressed the screw instabilities in BH magnetosphere. Gruzinov (1999) argued that the magnetic field with a bunch of closed field lines connecting a Kerr BH with a disk can be screw-unstable, resulting in the release of magnetic energy with the flares at the disk. Li (2000a) discussed the screw instability of the magnetic field in the BZ process, leading to a stringent upper bound to the BZ power. Wang et al. (2004, hereafter W04) studied the screw instability in the MC process. They concluded that this instability could occur at some place away from the inner edge of the disk, provided that the BH spin $a_ * $ and the power-law index $n$ for the variation of the magnetic field on the disk are greater than some critical values.  In this paper we attempt to combine the screw instability of the magnetic field with the coexistence of the BZ and MC processes. To facilitate the description henceforth we refer to the screw instability of the magnetic field occurring in the BZ and MC processes as SIBZ and SIMC, respectively. It is shown that both SIBZ and SIMC can occur, provided that the following parameters are greater than some critical values: (\\ref{eq1}) the BH spin, (\\ref{eq2}) the power-law index describing the variation of the magnetic field at the disk, and (\\ref{eq3}) the vertical height of the astrophysical load above the equatorial plane of the Kerr BH.  The features of several GRB-SNe are well fitted in our model. (\\ref{eq1}) The overall duration of the GRBs is fitted by the evolution of the half-opening angles. (\\ref{eq2}) The true energies of several GRBs are fitted by the energy extracted in the BZ process, and the energies of associated SNe are fitted by the energy transferred in the MC process. (\\ref{eq3}) The variability time scales of tens of msec in the light curves of several GRBs are fitted by two successive flares due to SIBZ.  This paper is organized as follows. In $\\S$ 2 we derived a criterion of SIBZ based on the Kruskal-Shafranov criterion and some simplified assumptions on the remote load. In $\\S$ 3 we discuss the time scale and energy extraction from a Kerr BH in the context of the suspended accretion state. In $\\S$ 4 we propose a scenario for the origin of the variation in the light curves of GRBs based on the flares arising from SIBZ. Finally, in $\\S$ 5, we summarize the main results and discuss some issues related to our model. Throughout this paper the geometric units $G = c = 1$ are used. ",
        "conclusions": "\\subsection{Mechanism for the recurrent occurrence of SIBZ}  In this paper we discuss the possibility of the modulations of SIBZ on the light curves of GRBs. One of the puzzles is what mechanism leads to the recurrent occurrence of SIBZ, and prevents the magnetic field from settling to a screw-stable configuration.  As argued in MT82 the BH magnetosphere consists of a series of magnetic surfaces connecting the horizon with the loads, and the total electric current $I$ flowing downward through an \\textbf{\\textit{m}}-loop is proportional to the toroidal magnetic field $B^T$ by Ampere's law. It is argued that these magnetic surfaces can be regarded as an equivalent circuit, in which each loop consists of two adjacent magnetic surfaces (W02). The total electric current flowing downward through an \\textbf{\\textit{m}}-loop is exactly equal to the algebraic sum of the poloidal currents flowing in the loops (W03).  As shown in Figure 1 the critical magnetic surface (henceforth CMS) for SIBZ is represented by the critical line $M{M}'$. According to the criterion (\\ref{eq7}) only the toroidal magnetic field outside CMS is depressed by SIBZ, while the toroidal and poloidal components of the magnetic field within CMS are little affected. On the other hand, the poloidal magnetic field outside CMS still exists in spite of the occurrence of SIBZ, and the toroidal magnetic field outside CMS will be recovered because of the twist of the poloidal magnetic field arising from the rotation of the BH. So, the rotation of the BH is the main mechanism for the recurrent occurrence of SIBZ in the BH magnetosphere.   \\subsection{Rotation period of BH and time scales of recovery of toroidal magnetic fields}  If we take BH mass as 10$M_ \\odot $, the rotation period of the BH is only $\\sim $1msec for the BH spin required by the criterion of SIBZ. This result implies that toroidal magnetic fields can not be recovered in one period of a rotating BH. How to explain the discrepancy between BH rotation and the time scales required for recovery of toroidal fields ?  In spite of lack of detailed knowledge of the screw instability, it is helpful to imagine the magnetic field line as an elastic string. The rotating BH always twists the field line, while the field line tries to untwist itself. Once the toroidal component of the magnetic field is strong enough to satisfy the criterion, the screw instability will occurs, just as a twisted elastic string releases its energy under appropriate conditions. In our model we simulate the process for twisting the field line by a transient process for accumulating magnetic energy in the inductor $\\Delta L$ in the equivalent $R-L$ circuit, which corresponds to the increase of the toroidal magnetic field, and the variability time scales of tens of msec in the light curves of several GRBs are fitted as the time interval between two successive flares due to SIBZ. Thus the threshold of toroidal magnetic field (magnetic energy) cannot be recovered in only one rotation of BH, just as a threshold of twisting more than one turn is required for an elastic string to release its energy.  \\ \\subsection{An explanation for GRBs with XRFs and XRRs }  Recently, much attention has been paid to the issue of X-ray flashes (XRFs), X-ray-rich gamma-ray bursts (XRRs) and GRBs, since HETE-2 provided strong evidence that the properties of these three kinds of bursts form a continuum, and therefore these objects are probably the same phenomenon (Lamb et al. 2004a, 2004b, 2005). The observations from HETE-2 motivate some authors to seek a unified model of these bursts. The most competitive unified models of these bursts are off-axis jet model (Yamazaki, Ioka, {\\&} Nakamura 2002; Lamb et al. 2005) and two-component jet model (Berger et al. 2003; Huang et al. 2004), in which XRFs, XRRs and GRBs arise from the differences in the viewing angles. Unfortunately, a detailed discussion on producing different viewing angles for XRFs, XRRs and GRBs has not been given in the above works.  Our argument on SIBZ and SIMC may be helpful to understanding this issue. It is believed that a disk is probably surrounded by a high-temperature corona analogous to the solar corona (Liang {\\&} Price 1977; Haardt 1991; Zhang et al 2000). Very recently, some authors argued that the coronal heating in some stars including the Sun is probably related to dissipation of currents, and very strong X-ray emissions arise from variation of magnetic fields (Galsgaard {\\&} Parnell 2004; Peter et al. 2004). Analogously, if the corona exists above the disk in our model, we expect that it might be heated by the induced currents due to SIMC and SIBZ. Therefore a very strong X-ray emission would be produced to form XRFs or XRRs.  Although our model may be too simplified and idealized with respect to the real situation, it provides a possible scenario for the occurrence of the screw instability in BH magnetosphere, and it may be helpful to understanding some astrophysical observations. We hope to improve our model by combining more observations in the future."
    },
    "0601/astro-ph0601477_arXiv.txt": {
        "abstract": "The New Luyten Two-Tenths catalog contains a large number of high-proper motion white dwarf candidates that remain to be spectroscopically confirmed. We present new spectroscopic observations as well as SDSS archival spectra of 49 white dwarf candidates which have been selected from the revised NLTT catalog of \\citet{sal2003}. Out of these, 34 are cool DA white dwarfs with temperatures ranging from approximately 5000 K up to 11690 K, and 11 are DC white dwarfs with temperatures ranging from 4300 K (NLTT~18555) up to 11000 K. Three of the DA white dwarfs also display abundances of heavy elements (NLTT 3915, NLTT 44986 and NLTT 43806) and one is a cool magnetic white dwarf (NLTT 44447) with an estimated magnetic field strength of 1.3 MG. We also present a new cool DQ white dwarf (NLTT 31347) with an estimated temperature of 6250 K. We supplement our sample with SDSS {\\it ugriz} photometry for a fraction of the newly identified white dwarfs. A kinematical study of this sample of white dwarfs, characterized by proper motions ranging from 0.136 to 0.611$\\arcsec$yr$^{-1}$ suggest that they belong to the thin disk population. ",
        "introduction": "The current census of white dwarfs in the solar neighborhood is believed to be complete only to about 13 pc, and is measurably incomplete within 20 pc \\citep{hol2002,sch2004}. The local sample of white dwarfs appears to be an old population consisting of mostly cool white dwarfs. The average absolute magnitude for the local census of white dwarfs is $M_V = 13.7$ which corresponds to a temperature of $\\sim 7000$ K. A significant fraction of white dwarfs within 13 pc \\citep[$20\\pm8\\%$]{kaw2003} show the presence of a magnetic field. \\citet{hol2002} determined that over a quarter of the white dwarfs within 20 pc are in binary systems. Based on recent literature, we updated the spectral classifications of the sample of white dwarfs within 20 pc \\citep{hol2002} and found that about 70\\% of the stars are hydrogen-rich. The remaining 30\\% are helium-rich, and out of these helium rich white dwarfs about half are DC white dwarfs, about a third are cool DQ white dwarfs, while the remainder are DZ white dwarfs.  The New Luyten Two-Tenths (NLTT) catalog has been used to search for nearby cool dwarfs and subdwarfs \\citep{giz1997,rei2003,yon2003,rei2005} and for white dwarfs \\citep{lie1977,hin1986,ven2003,kaw2004,kaw2005}. However, a large number of objects still require a formal spectroscopic confirmation. Recently, \\citet{sal2003} have revised the coordinates and proper motions of most stars in the NLTT catalog by cross-correlating the catalog with the Two Micron All Sky Survey (2MASS) and the USNO-A catalogs. One problem with the original catalog was that there was no obvious separation between different populations because the R and B photographic bands do not provide a broad enough baseline to distinguish between main-sequence, subdwarf and white dwarf populations. \\citet{sal2002} showed that using an optical/infrared ($V-J$) reduced proper motion diagram they could distinguish between these groups of stars and from this diagram they listed 23 white dwarf candidates within 20 pc from the Sun. \\citet{ven2003} and \\citet{kaw2004} found that only a third of these nearby candidates are white dwarfs, with the remainder being cool red dwarfs and subdwarfs.  Recently, several projects have been initiated to search for high-proper motion stars, such as the SuperCOSMOS-RECONS survey \\citep{ham2004} and LSPM Catalog \\citep{lep2005} which utilizes the SuperCOSMOS Sky Survey \\citep{ham2001} and SUPERBLINK \\citep{lep2002}, respectively. These surveys have detected many high-proper motion stars which are either nearby or have been labeled as halo candidates. One such object is the cool white dwarf WD 0346$+$246 \\citep{ham1997}, which has definite halo kinematics and is very cool ($T\\sim3800$ K) and hence old enough to belong to the halo. A total of 299 new stars with $0.4\\arcsec$ yr$^{-1}$ have been discovered using the SuperCosmos-RECONS survey \\citep{sub2005}, out of which 148 have proper motions greater than $0.5\\arcsec$yr$^{-1}$ and 43 are within 25 pc from the Sun. \\citet{lep2002,lep2003} reported the discovery of 198 new stars with proper motions greater than $0.5\\arcsec$yr$^{-1}$ using SUPERBLINK. Their survey also included the nearby halo white dwarf PM J13420$-$3415 \\citep{lep2005b}.  The NLTT catalog appears to be complete at Galactic latitudes $|b| > 15^\\circ$, however at low Galactic latitudes the catalog appears to be significantly incomplete. On the other hand, \\citet{sal2003} notes that the completeness for white dwarfs is very high, i.e., the coverage of white dwarfs appears to be completely uniform. The reason for this may be that Luyten concentrated on blue objects, which are no more common in the plane than elsewhere, and hence Luyten was able to detect most high-proper motion white dwarfs above $\\delta > -32.5\\degr$. Another survey to search for cool white dwarfs with high-proper motion using the Guide Star Catalog II was initiated by \\citet{car2005} with the aim of placing constraints on the halo white dwarf space density. Their search resulted in the discovery of 24 new white dwarfs. Many earlier proper motion surveys, in particular the Luyten proper motion surveys relied on plate pairs taken a decade or more apart, and very high-proper motion stars were most likely lost in the background. A search for ultra-high proper motion stars was performed by \\citet{tee2003} using the SkyMorph database of the Near Earth Asteroid Tracking (NEAT) project \\citep{pra1999} and found a main-sequence star with spectral type M6.5 and with a proper motion of $5.05\\arcsec\\pm0.03\\arcsec {\\rm yr}^{-1}$ at a distance of $2.4^{+0.7}_{-0.4}$ pc which they obtained from the trigonometric parallax.  Since most of the stars in the NLTT catalog are brighter than 19th magnitude, we do not expect to find many halo white dwarf candidates. Since the halo star formation history is assumed to be a burst occurring 12 Gyrs ago, then most of the halo white dwarfs would be required to have cooling ages greater than 10 Gyrs (and hence $M_V \\ga 16$). Given that the percentage of halo white dwarfs in the local population is $\\sim 2\\%$ \\citep{pau2005}, then most halo candidates are likely to be at large distances and hence too faint to be included in the NLTT catalog.  In this paper, we pursue our study of white dwarf candidates \\citep{kaw2004} from the revised NLTT (rNLTT) catalog of \\citet{sal2003}. Some preliminary results were presented in \\citet{kaw2005}. We present the list of white dwarfs observed in \\S 2 and their spectroscopic observations in \\S 2.1 with complementary optical and infrared photometric data presented in \\S 2.2. The model atmospheres and spectra we used to analyze our data are described in \\S 3 and in \\S 4 we present our analysis of the new white dwarf stars. We discuss our findings in \\S 5 and summarize in \\S 6. ",
        "conclusions": "\\begin{figure} \\plotone{f8.eps} \\caption{{\\it Top:} Effective temperature and surface gravity distribution for the DA white dwarfs with the mass-radius relations ($0.4- 1.2 M_\\odot$ (in step of $0.2 M_\\odot$)) of \\citet{ben1999} for carbon interiors with a hydrogen envelope of $M_H/M_* = 10^{-4}$ and a metallicity of Z=0. {\\it Bottom:} Masses and ages for the DA white dwarfs ({\\it filled circles}). White dwarfs for which we have assumed $\\log{g}=8.0$ (i.e., DC, DQ, DZ) are shown as open circles. Typical errorbars for the masses and ages are shown at the top left corner of the figure. \\label{fig_age_mass}} \\end{figure}  The sample of rNLTT white dwarfs are old white dwarfs with ages of the order of $10^9$ to $7\\times10^9$ years. Figure~\\ref{fig_age_mass} shows the effective temperature versus surface gravity of the DA white dwarfs compared to the mass-radius relations of \\citet{ben1999}, as well as the masses and ages for all the white dwarfs. No independent gravity measurements are available for non-DA white dwarfs and for the ultra-cool DA white dwarf NLTT~8581, and we simply assumed $\\log{g}=8$. Three objects (NLTT~14307, NLTT~43827, and NLTT~44986) have masses in excess of $1.2 M_{\\odot}$ and belong to a sequence of high-mass white dwarfs also identified in EUV surveys \\citep{ven1997} and in common-proper motion surveys \\citep{sil2001}. Note also that three cool white dwarfs ($T_{eff} < 5200$K) also seem to be characterized by a low surface gravity ($\\log{g} \\sim 7$); the low surface gravity was diagnosed by their relatively strong H$\\alpha$ line strengths for their assigned effective temperatures. For the remaining 37 DA white dwarfs, the mass average and dispersion are $<M>= 0.69 M_{\\odot}$ and $\\sigma_M = 0.17 M_{\\odot}$ which is slightly higher than the mass average ($<M>= 0.61 M_{\\odot}$) of the H-rich subsample of \\citet{ber2001}. The distribution among various spectral types follows established trends. First, 45 objects (74\\%) from the present selection of 61 white dwarfs are hydrogen-rich, three are DAZ white dwarfs (NLTT~3915, NLTT~43806 and NLTT~44986). For at least three of the new DA white dwarfs the classification is primarily secured by the presence of H$\\alpha$, which underlines the importance of obtaining red spectroscopy for a proper classification of cool white dwarfs. We suspect that a few DC white dwarfs in the \\citet{mcc1999} catalog of spectroscopically identified white dwarfs may turn out to be DA white dwarfs. Some 14 objects are classified as DC white dwarfs (23\\%) and the two remaining objects consist of a new DQ white dwarf (NLTT~31347) and the DZ white dwarf NLTT~40607 \\citep{kaw2004}. For comparison, we determined the composition of 200 well-studied rNLTT white dwarfs using the \\cite{mcc1999} catalog: some 134 objects (67\\%) received a primary label ``DA'', 14 received the label ``DB'' (7\\%), 15 are DQ (8\\%), 11 are ``DZ'' (5\\%), and in last resort 26 received the label ``DC'' (13\\%). Combining the two data sets ($N=261$), 69\\% of rNLTT white dwarfs are spectral type DA, and 31\\% non-DA. The DA:non-DA number ratio varies as a function of temperatures. By dividing the survey in the temperature bins $<6$, $6-8$, $8-10$, and $10-12\\times 10^3$K, we observe a varying ratio 1.5:1, 3.6:1, 3.5:1, and 5:1, respectively. The ratio increases sharply with effective temperatures. The result of this proper motion survey are to be contrasted with the result of the parallax survey of \\citet{blr2001} who estimated DA:non-DA ratios varying from 1.2:1 in the $<6$ bin, to 2:1 in a sparsely populated $8-10$ bin. The number of DC white dwarfs are markedly lower in the present proper motion survey. The DA:non-DA ratio in $<6$ bin is only a lower limit, since for many white dwarfs in this bin it is very difficult to distinguish between H-or He-rich atmospheres without infrared photometry. For many of the objects accurate infrared photometry is required to ascertain the classification. Note that in this proper motion survey we selected the brighter candidates for spectroscopic observations, and therefore there may exist a slight bias against DC white dwarfs which tend to be fainter than DA white dwarfs.  Over 150 white dwarf candidates from our original selection \\citep{kaw2004} remain to be spectroscopically confirmed. Eight white dwarfs (NLTT 529, NLTT 3915, NLTT 8435, NLTT 14307, NLTT 18555, NLTT 43806, NLTT 53447 and NLTT 56805) in the present compilation are found at distances closer than 20 pc from the Sun and therefore contribute to the local census. Other nearby white dwarfs are potentially among these remaining $\\sim 150$ white dwarf candidates. We evaluate that $\\sim 7\\%$ of these are likely to lie within 20pc of the Sun. We assumed that all of these candidates are white dwarfs for which we calculated their absolute magnitude using $VJH$ photometry, and then determined their distance.  \\begin{figure} \\epsscale{.80} \\plotone{f9.eps} \\caption{$U$ vs. $V$ diagram showing the white dwarfs listed in Table~\\ref{tbl2}. The $2\\sigma$ velocity ellipse of the thin disk population is shown ({\\it full line}), along with the $2\\sigma$ ellipse of the thick disk ({\\it dashed line}), and the $1\\sigma$ ellipse of the halo ({\\it dotted line}) populations. \\label{fig_kinematic}} \\end{figure}  Most white dwarfs in the neighborhood of the Sun belong to the thin disk population, but a significant number of thick disk white dwarfs are expected. We have calculated the velocity components $U$,$V$ and $W$ using \\citet{joh1987} and assuming $v_{\\rm rad} = 0$. We have calculated absolute magnitudes using the temperatures and surface gravities as described in the previous sections and listed in Table~\\ref{tbl2}. We obtained the distances using the apparent magnitudes and the absolute magnitudes listed in Table~\\ref{tbl1} and Table~\\ref{tbl3}. Figure~\\ref{fig_kinematic} compares measured $U$ and $V$ to the $2\\sigma$ velocity ellipse of the thin disk population, the $2\\sigma$ ellipse of the thick disk population and the $1\\sigma$ ellipse of the halo population \\citep{chi2000}. The following velocities are all expressed in km~s$^{-1}$. All three distributions, $UVW$, were fitted with Gaussian functions, where $\\sigma_U = 34$, $\\sigma_V = 22$, $\\sigma_W = 19$ and $<U> = -2$, $<V> = -13$, $<W> = 1$. Our velocity dispersions are in good agreement with those of a thin disk population \\citep[$\\sigma_U = 34$, $\\sigma_V = 21$, $\\sigma_W = 18$][]{bin1998}, however our $<V> = -13$ is only slightly more negative then their quoted value of $<V> = -6$ for normal stars in the thin disk. Hence, the measured velocity distributions imply that the majority of the objects in this study belong to the thin disk population.  The local sample of white dwarfs should consist of a significant number of white dwarfs that belong to the thick disk population. Several estimates of the fraction of thick disk white dwarfs in the Solar neighborhood have been made, from $\\sim5\\%$ \\citep{han2003} up to 25\\% \\citep{rei2005b}. The main reason for the difference between the two values is that the first assumes a single power-law for the initial mass function, which produces more low-mass main-sequence stars and hence a small white dwarf mass-fraction in the disk ($N_{WD}/N_{MS}$). \\citet{rei2005b} used a two-component power law, which produces fewer low-mass main sequence stars and hence the white dwarf mass fraction for the thick disk will be higher. Therefore, the thick disk could contribute as much as 25\\% of the local white dwarfs.  It is difficult to disentangle thin  versus thick disk populations solely based on the kinematics. A study of the distribution of white dwarfs in the $UV$ plane led \\citet{kaw2004} to conclude that 5\\% of their sample of $N=417$ white dwarf candidates extracted from the rNLTT catalog of \\citet{sal2003} lay in  a high-velocity tail and, therefore, may belong to the thick disk population. \\citet{kaw2004} used the $V$-velocity distribution which is most sensitive to the presence of thick disk white dwarfs which are characterized by a velocity dispersion twice the dispersion of thin disk white dwarfs ($\\sigma_{\\rm thick} =50$ versus $\\sigma_{\\rm thin} = 21$). Therefore, the fraction of thick disk white dwarfs estimated by \\citet{kaw2004} is a lower limit, considering that several objects assimilated with the thin disk population may in fact be part of the bulk of the thick disk population only betrayed by its high-velocity tail. Similarly in this study of 61 objects, four objects in Figure~\\ref{fig_kinematic} are clearly lying outside the $2\\sigma$ velocity ellipse of the thin-disk population. If these four objects were assumed to belong to the thick-disk, then the percentage of thick-disk white dwarfs in our sample would be $6.5\\%$ which is consistent with \\citet{han2003}. Note that this is a lower limit."
    },
    "0601/astro-ph0601194_arXiv.txt": {
        "abstract": "{} {We present a detailed analysis of a very high resolution (R~$\\approx 112,000$) spectrum of  the quasar HE {0515$-$4414} obtained using the High Accuracy Radial velocity Planet Searcher (HARPS) mounted on the ESO 3.6~m telescope at the La Silla observatory. The main aim is to use HARPS spectrum of very high wavelength calibration accuracy (better than 1~m\\AA), to constrain the variation of $\\alpha=e^2/\\hbar c$ and investigate any possible systematic inaccuracies in the wavelength calibration of the UV Echelle Spectrograph (UVES) mounted on the ESO Very Large Telescope (VLT).} {A cross-correlation analysis between the Th-Ar lamp spectra obtained with HARPS and UVES is carried  out to detect any possible shift between the two spectra. Absolute wavelength calibration accuracies, and how that translate to the uncertainties in \\dela are computed using Gaussian fits for both lamp spectra. The value of \\dela  at \\zabs~=~1.1508 is obtained using Many Multiplet method, and simultaneous Voigt profile fits of HARPS and UVES spectra.} {We find the shift between the HARPS and UVES spectra has mean around zero with a dispersion of $\\sigma\\simeq 1$ m\\AA. This is shown to be well within the wavelength calibration accuracy of UVES (i.e $\\sigma\\simeq 4$ m\\AA). We show that the uncertainties in the wavelength calibration induce an error of about, \\dela $~\\le 10^{-6}$, in the determination of the variation of the fine-structure constant. Thus, the results of non-evolving \\dela reported in the literature based on UVES/VLT data should not be heavily influenced  by problems related to wavelength calibration uncertainties. Our higher resolution spectrum of the \\zabs~=~1.1508 Damped Lyman-$\\alpha$ system toward HE {0515$-$4414} reveals more components compared to the UVES spectrum. Using only \\feii lines of  \\zabs~=~1.1508 system, we obtain \\dela~=~${ (0.05\\pm0.24)\\times10^{-5}}$. This result is consistent with the earlier measurement for this system using the UVES spectrum alone.} {} ",
        "introduction": "\\label{sect:Int} Some of the modern theories of fundamental physics, such as SUSY, GUT and Super-string theory, allow possible space and time variations of the fundamental constants, thus motivating an experimental search for such a variation (Uzan 2003 and 2004 for a detail review on the subject).  Murphy et al. (2003), applying the Many Multiplet method (MM method) to 143 complex metal line systems, claimed a non-zero variation of the fine-structure constant, $\\alpha=e^2/\\hbar c$: $\\langle\\Delta\\alpha/\\alpha\\rangle = (-0.57\\pm0.11)\\times10^{-5}$ for $0.2\\le z \\le 3.5$, where $\\Delta\\alpha/\\alpha=(\\alpha_{z}-\\alpha_{0})/\\alpha_{0}$, with $\\alpha_{0}$ being the present value and $\\alpha_{z}$ its value at redshift $z$. This result, if true, would have very important implications to our understanding of fundamental physics and has therefore motivated new activities in the field. Search for the possible time-variation of $\\alpha$ using alkali doublets has started long ago (Bahcall et al. 1967). The alkali-doublet method is a clean method for constraining the variation in $\\alpha$ using spectral lines because it uses transitions from the same species (Wolfe et al. 1976; Levshakov 1994; Potekhin et al. 1994; Cowie \\& Songaila, 1995; Varshalovich et al. 1996; Varshalovich et al. 2000; Murphy et al. 2001a; Martinez et al. 2003; Chand et al. 2005). The tightest constraint obtained using this method till date is \\dela = (0.15$\\pm$0.44) $\\times10^{-5}$ at $z \\sim 2$ (Chand et al. 2005). \\par Studies based on heavy element molecular absorption lines seen in the radio/mm wavelength range are more sensitive than that based on optical/UV absorption lines. They usually provide constraints on the variation of a combination of the fine-structure constant, the proton g-factor ($G_{p}$) and the electron-to-proton mass ratio ($\\mu$). Murphy et al. (2001b) have obtained \\dela = $(-0.10\\pm0.22)\\times10^{-5}$ at z = 0.2467 and \\dela = $(-0.08\\pm0.27)\\times 10^{-5}$ at z = 0.6847, assuming a constant proton g-factor ($G_{p}$). It has been pointed out that OH lines are very useful in simultaneously constraining various fundamental constants (Chengalur \\& Kanekar 2003; Kanekar et al. 2004; Darling 2003, 2004). These studies have provided \\dela = $(0.6\\pm1.0)\\times 10^{-5}$ for an absorption system at \\zabs = 0.247 toward PKS~1413+135. Such studies have not been performed yet at higher redshift (i.e $z$~$\\ge1$) due to the lack of molecular absorption systems. \\par Constraints on the variations of $\\alpha$ are also obtained from terrestrial measurements. The most stringent constrain has been obtained from the analysis of the Oklo phenomenon. Fujii et al. (2000)  find that $\\Delta\\alpha/\\alpha = (-0.8\\pm1.0)\\times10^{-8}$ over a period of about 2 billion years (or $z\\simeq0.45$). Laboratory experiments also give very stringent constraints on the local variation of $\\alpha$. Marion et al. (2003) have obtained, $\\Delta\\alpha/\\alpha\\Delta t = (-0.4\\pm16)\\times10^{-16}\\,{\\rm yr}^{-1}$, by comparing the hyperfine transition in $^{87}$Rb and  $^{133}$Cs over a period of 4 years assuming no variation in the magnetic moments. Fischer et al. (2004) have obtained, $\\Delta\\alpha/\\alpha\\Delta t = (-0.9\\pm2.9)\\times10^{-16}\\,{\\rm yr}^{-1}$, by comparing the absolute $1S-2S$ transition of atomic hydrogen to the ground state of Cesium. A linear extrapolation gives a constraint of $-1.3\\times10^{-6}\\le$ \\dela$\\le1.9\\times10^{-6}$ at $z = 1$ for the most favored cosmology ($\\Omega_m = 0.27$, $\\Omega_{\\Lambda} =0.73$ and $ h=0.71$). \\par Clearly all the experimental results summarized above are consistent with no variation of $\\alpha$. However, these results do not directly conflict with the positive detection by Murphy et al. (2003) either because of the insufficient sensitivity of the method (as in the case of alkali doublets) or because of  the different redshift coverage (as in the case of radio and terrestrial measurements). However,  recent attempts using the MM method (or its modified version) applied to very high quality UVES spectra have resulted in null detections. The analysis of Fe~{\\sc ii} multiplets and Mg~{\\sc ii} doublets in a homogeneous sample of 23 systems has yielded a stringent constraint, \\dela = $(-0.06\\pm0.06)\\times10^{-5}$ (Chand et al. 2004; Srianand et al. 2004). Modified MM method analysis of \\zabs = 1.1508 toward HE 0515$-$4414 that avoids possible complications due to isotopic abundances has resulted in \\dela = $(0.01 \\pm 0.17)\\times 10^{-5}$ (Quast et al. 2004). Levshakov et al. (2005b) have re-analysis this system using the single ion differential alpha measurement method as described in Levshakov et al. (2005a), and obtained \\dela = $(-0.007\\pm0.084)\\times10^{-5}$. Clearly all studies based on VLT-UVES data are in contradiction with the conclusions of Murphy et al. (2003). \\par A first possible concern about these studies is the real accuracy and robustness of the various calibration procedures. A second possible source of uncertainty comes from the multi-component Voigt-profile decomposition. It is very important to check how sensitive the derived constraints are to the profile decomposition. This can be done by performing the analysis on data of higher resolution than typical UVES (or HIRES) spectra. The best way to investigate all this is to compare data taken by UVES (or HIRES) with data on the same object taken with another completely independent, well controlled, and higher spectral resolution instrument. The advent of HARPS mounted on the ESO 3.6~m telescope makes this possible. Unfortunately this is only possible on the brightest quasar in the southern sky, HE~0515$-$4414. \\par This forms the basic motivations of this work. We report the analysis of the \\zabs = 1.15 DLA system toward QSO HE 0515$-$4414 (De la Varga et al. 2000, Quast et al. 2004, 2005) using very high resolution (R$\\sim112,000$) spectra obtained with HARPS mounted on the ESO 3.6~m telescope. The organization of the  paper is  as follows. The HARPS observations of HE 0515$-$4414 are described in  Section 2. Calibration accuracy and comparison with the UVES observations are discussed in Section 3. In Section 4 we present the joint analysis of the HARPS and UVES spectra. Results are summarized and discussed  in Section 5. ",
        "conclusions": "\\label{sect:result} In this paper, we present a very high resolution (R~=~112,000) spectrum of QSO HE 0515$-$4414 obtained using HARPS. We have used the high wavelength calibration accuracy and high spectral resolution capabilities of HARPS to address the following issues.\\par We compare the lamp spectra obtained with UVES and HARPS. Using cross-correlation analysis we show that any possible relative shift between the two spectra are within 2~m\\AA. Using Gaussian fits to unblended lamp emission lines, we find that the absolute wavelength calibration of HARPS is very robust with rms deviation of 0.87~m\\AA~with respect to the wavelengths tabulated in Cuyper et al. (1998). This is about a factor of 4 better than that of UVES ($\\sigma=4.08$ m\\AA,~see Fig.~\\ref{dlam.fig}). Thus the small shifts noted between the HARPS and UVES lamp spectra are well within the typical wavelength calibration accuracy of UVES. We have derived the error on \\dela measurements due to the calibration accuracy alone. For UVES and HARPS spectra this is found to be  respectively $\\sigma=0.96\\times10^{-6}$ and  $\\sigma=0.19\\times10^{-6}$ for a typical system with three well detached components. The value obtained for the UVES spectrum is also consistent with that of HIRES (Murphy et al. 2003). \\par This shows that HARPS is the ideal instrument for this kind of measurement. Unfortunately it is mounted on the 3.6~m telescope at La~Silla and only HE~0515$-$4414 is bright enough to be observed in a reasonable amount of time. This shows as well that the UVES spectra reduced (or calibrated) with the UVES pipeline and used in the literature to constrain \\dela (Srianand et al. 2004 and Chand et al. 2004, Quast et al. 2004, Chand et al. 2005) do not suffer from major systematic error in the wavelength calibration. \\par We have obtained the accurate multi-component structure using the higher resolution data (R~$\\approx 112,000$ for HARPS compared to $\\approx 55,000$ for UVES). The best fit to the profiles obtained by fitting simultaneously the HARPS data (of higher resolution) and the UVES data (of better S/N ratio) require additional components as compared to the fit using the UVES data alone (Quast et al. 2004). Using this new sub-component decomposition and both HARPS and UVES data, we find \\dela~$=(0.05\\pm0.24)\\times10^{-5}$. This is consistent with the results derived by Quast et al. (2004) from the UVES data alone. Indeed, we have in addition re-analyzed the UVES data which was used in Quast et al. (2004) (without using the component structure from HARPS data), to estimate the effect of different independent algorithms used to obtain error spectra, to combine the data, to fit the continuum and to fit the absorption lines. We find that the best-fit parameters as well as the \\dela measurement (\\dela~=~${[0.10\\pm0.22]\\times10^{-5}}$), obtained by our independent analysis, are consistent with  that of Quast et al. (2004) (\\dela~=~${[0.01\\pm0.17]\\times10^{-5}}$).\\par We note that the precision on the \\dela measurement obtained using the HARPS spectrum, which is of high resolution and low S/N ratio, is similar to that obtained from the UVES spectrum, which is of lower resolution and higher S/N ratio. Therefore, the improvement in the wavelength calibration accuracy by an order of magnitude using  HARPS will be effective to improve the constrain on \\dela  only if high S/N ratio can also be obtained. This could be possible if an instrument such as HARPS can be mounted on bigger telescopes."
    },
    "0601/astro-ph0601688_arXiv.txt": {
        "abstract": "A high quality Keck spectrum of the \\ha\\ line in NGC~4395 reveals symmetric exponential wings, $f_v\\propto e^{-v/\\sigma}$, with $\\sigma\\simeq 500$~km~s$^{-1}$. The wings extend out to $\\gtsim 2500$~km~s$^{-1}$ from the line core, and down to a flux density of $\\ltsim 10^{-3}$ of the peak flux density. Numerical and analytic calculations indicate that exponential wings are expected for optically thin, isotropic, thermal electron scattering. Such scattering produces exponential wings with $\\sigma\\simeq 1.1\\sigma_e(\\ln\\tau_e^{-1})^{0.45}$, where $\\sigma_e$ is the electron velocity dispersion, and $\\tau_e$ is the electron scattering optical depth. The \\ha\\ wings in NGC~4395 are well fit by an electron scattering model with $\\tau_e=0.34$, and an electron temperature $T_e=1.1\\times 10^4$K. Such conditions are produced in photoionized gas with an ionization parameter $U\\simeq 0.3$, as expected in the broad line region (BLR). Similar analysis of the [O~III]~$\\lambda 5007$ line yields $\\tau_e<0.01$, consistent with the lower ionization in the narrow line region. If the electron scattering interpretation is correct, there should be a tight correlation between $\\tau_e$ and the ionizing flux on time scales shorter than the BLR dynamical time, or $\\sim 1$~week for NGC~4395. In contrast, the value of $\\sigma$ should remain nearly constant on these time scales. Such wings may be discernible in other objects with unusually narrow Balmer lines, and they can provide a useful direct probe of $T_e$ and $\\tau_e$ in the BLR. ",
        "introduction": "Early studies of the broad line profiles in Type~1 Active Galactic Nuclei (AGN) suggested that the large observed line widths ($\\gtsim 3000$~km~s$^{-1}$) are produced by electron scattering within the photoionized gas (Weymann 1970; Mathis 1970). With the development of photoionization models for the Broad Line Region (BLR), it was however realized that the expected electron scattering optical depth, $\\tau_e$, within the BLR gas is $\\ll 1$, which is too low to explain the observed line profiles (Davidson \\& Netzer 1979). Later studies of the size of the BLR, based on reverberation mappings, indicated that the BLR is more compact than previously thought (e.g. Rees, Netzer, \\& Ferland 1989), and thus $\\tau_e$ is expected to be higher than previously thought, though it is not expected to reach $\\tau_e> 1$, required to explain the line profiles purely by electron scattering. There is also good evidence now that the line broadening in the BLR is mostly due to virialized bulk motion (e.g. Peterson \\& Wandel 2000). However, electron scattering by a significantly hotter gas outside the BLR gas ($T\\sim 10^{6-7}K$ vs. $T\\sim 10^4K$) remains a viable option (Shields \\& McKee 1981; Kallman \\& Krolik 1986; Emmering et al. 1992; Bottorff et al. 1997). Evidence for electron scattering of continuum photons in AGN is provided by wavelength independent polarization in some objects (e.g. Antonucci et al. 1993), but there is yet no conclusive detection of electron scattering effects on the broad line profiles.  A possible exception may occur in NGC~1068, where spatially resolved spectropolarimetry reveals a broadened \\ha\\ line, which may result from electron scattering in warm ($T\\sim 10^5K$) gas (Miller, Goodrich \\& Mathews 1991). A spectropolarimetric search for polarized broad lines in low luminosity AGN (Barth et al. 1999) revealed a polarized BLR component in the lowest luminosity type~I AGN, NGC~4395. Based on that detection, Barth et al. suggested that higher S/N spectropolarimetry may reveal a broadened \\ha\\ component, if the polarization is induced by electron scattering. Here we report on a high quality Keck observation of the \\ha\\ line in NGC~4395, which indeed reveals extended symmetric exponential wings (though their polarization is unknown). As shown below, these wings are well fit by optically thin electron scattering by electrons at $T_e=1.14\\times10^4K$, with $\\tau_e=0.34$, which are plausible conditions for the BLR gas.  The paper is structured as follows. Section 2 describes the observations which revealed the exponential wings in the \\ha\\ profile of NGC~4395. Section 3 describes the numerical calculation of the electron scattering line profiles, which yield exponential wings, $f_v=f_0e^{-v/\\sigma}$. A fitting formula for $\\sigma$ as a function of $T_e$ and $\\tau_e$ is provided, together with a simple analytic derivation. In section 4 the observed \\ha\\ profile is fit with the electron scattering model, yielding $T_e$ and $\\tau_e$, the origin of the scattering gas is discussed, and some predictions are made. The main conclusions are summarized in section 5. ",
        "conclusions": "High quality Keck observations of the \\ha\\ line in NGC~4395 reveal symmetric exponential wings of the form $f_v\\propto e^{-v/\\sigma}$, with $\\sigma\\simeq 500$~km~s$^{-1}$. The wings are fit well by an isothermal and isotropic electron scattering model with $\\tau_e=0.34$, and $T_e=1.14\\times 10^4$K, plausibly generated in photoionized gas with $U\\simeq 0.3$. The lack of electron scattering wings for the [O~III]~$\\lambda 5007$ line, and the value of $U$ indicate that the scattering occurs within the BLR gas where \\ha\\ is formed.  The electron scattering interpretation can be tested by looking for a strong correlation between $\\tau_e$ and the ionizing flux. High quality spectra of other AGN with very narrow Balmer lines will provide a further test of the electron scattering interpretation. If the electron scattering origin is verified, then the extended \\ha\\ wings can provide us with a new direct probe of $T_e$ and $\\tau_e$ within the BLR gas, and their time variability."
    },
    "0601/astro-ph0601127_arXiv.txt": {
        "abstract": "We attempt to detect optical GRB afterglows on images taken by the Canada France Hawaii Telescope for the Very Wide survey, component of the Legacy Survey. To do so, a Real Time Analysis System called \"Optically Selected GRB Afterglows\" has been installed on a dedicated computer in Hawaii. This pipeline automatically and quickly analyzes Megacam images and extracts from them a list of variable objects which is displayed on a web page for validation by a member of the collaboration. The Very Wide survey covers 1200 square degrees down to i'=23.5. This paper briefly explain the RTAS process. \\\\ \\\\ PACS 95.75.Mn - Image processing (including source extraction)  PACS 95.75.Wx - Time series analysis, time variability  PACS 98.70.Rz - $\\gamma$-ray sources; $\\gamma$-ray bursts  PACS 01.30.Cc - Conference proceedings ",
        "introduction": "\\subsection{CHFT at glance} The CFHT is a 3.6 m telescope located on the Mauna Kea in Hawaii. Built in the late 70's, it has been recently equiped with a new instrument, Megacam. The specifications of this CCD imager are the following: \\begin{itemize} \\item[-] 36 2048x4612 pixel CCDs with 0.185 arcsec/pixel resolution \\item[-] 1 square degree field of view \\item[-] 5 filters set (u*, g', r', i', z') \\end{itemize}  \\begin{figure}[h] \\begin{center} \\includegraphics[height=.075\\textheight]{malacrino_f1.eps} \\end{center} \\caption{This diagram shows the observing strategy for one field (each vertical line is for one exposure, exposure time depends on the filter, but is typically of the order of 100 second) } \\end{figure}  These characteristics, combined with the excellent climatic conditions at the site, provide images with a very good quality. Moreover, the images are pre-processed in quasi-real time with a pipeline called Elixir\\footnote{See http://www.cfht.hawaii.edu/Instruments/Elixir/}, which flattens and defringes each CCD and computes gross astrometry and photometry.  \\subsection{The Legacy Survey} This 5 years' project has started in march 2003, and is based on 3 different observing strategies: \\begin{itemize} \\item[-] The Wide Synoptic Survey covers 170 square degrees with all the 5 Megacam filters (u* g' r' i' z') down to approximatively i'=25.5. The main goal of this survey is to study large scale structures and matter distribution in the universe. \\item[-] The Deep Synoptic Survey covers only 4 square degrees, but down to i'=28.4 and through the whole filter set. Aimed mainly at the detection of 2000 type I Sne and the study of galaxy distribution, this survey will allow an accurate determination of cosmological parameters. \\item[-] The Very Wide Survey covers 1200 square degrees down to i'=23.5, but only through 3 filters (g' r' i'). Initially conceived to discover and follow 2000 Kuiper Belt Objects, we'll show that this observing strategy can be used to detect optical afterglows. \\end{itemize}  \\subsection{The Very Wide Survey} Each field is observed several times, according to the strategy explained in Figure 1. This recurrence can be used to compare images between them in order to detect variable or transient objects, such as GRB afterglows ",
        "conclusions": ""
    },
    "0601/astro-ph0601311_arXiv.txt": {
        "abstract": "A dispersion relation for electromagnetic wave propagation in a strongly magnetized cold plasma is deduced, taking photon--photon scattering into account. It is shown that the combined plasma and quantum electrodynamic effect is important for understanding the mode-structures in magnetar and pulsar atmospheres. The implications of our results are discussed. ",
        "introduction": "The quantum electrodynamical (QED) phenomenon of elastic photon--photon scattering, due to the interaction of photons with virtual electron--positron pairs, has recently received increased attention \\cite{Ding1992,Moulin1999,BMS2001,Eriksson2004,Shen-etal,Shen-Yu,Bulanov-etal,JPP,% Marklund-Brodin-Stenflo,MSSBS2005,RMP}. Several papers are motivated by the desire to detect photon--photon scattering in laboratories \\cite {Ding1992,Moulin1999,BMS2001,Eriksson2004}, whereas others \\cite {Shen-etal,Shen-Yu,Bulanov-etal} concern phenomena that might be relevant when the laser power is further increased to produce electric fields strengths close to the Schwinger field $\\sim 10^{18}\\, \\mathrm{V\\,m^{-1}}$ \\cite{Bulanov-etal}% . Up to now, however, observable effects of photon--photon scattering are likely to occur only for astrophysical systems \\cite {Baring-Harding,Shaviv-etal,Marklund-Brodin-Stenflo,magnetar,Bialynicki-Birula1970,% Adler,MSSBS2005}, where the large magnetic field strengths in pulsar and magnetar environments \\cite{magnetar,Duncan-Thompson,Palmer-etal} open up for QED processes to play an important role, leading to phenomena such as frequency down-shifting \\cite{Adler,Bialynicki-Birula1970} and lensing \\cite{Shaviv-etal}. \\ The frequency down-shifting \\ is a result of so called photon splitting \\cite {Adler,Bialynicki-Birula1970}, which is one of the consequences of elastic photon-photon scattering, and the process may even be responsible for the radio silence of magnetars \\cite{Baring-Harding}. Another QED-process of interest in pulsar and magnetar environments is pair-production \\cite{Beskin-book} due to the strong field interactions, which lead to the presence of an electron-positron pair plasma in the pulsar and magnetar atmospheres. However, with a few exceptions (e.g.\\ \\cite{MSSBS2005,JPP}), we note that most papers considering photon--photon scattering have omitted plasma effects when considering electromagnetic wave propagation under these conditions.  In the present paper we will consider electromagnetic wave propagation at an arbitrary angle to a strong external magnetic field $\\mathbf{B}_{0}$, and include the QED-effects associated with that field, as well as the influence of an electron-positron pair plasma. The former effect is described within the framework of the Heisenberg--Euler Lagrangian, which constitutes an effective theory of photon--photon scattering \\cite {Heisenberg-Euler,Schwinger}, and the latter contribution follows from elementary plasma theory. A comparatively general dispersion relation will be derived. It reduces to previous results in a number of limiting cases \\cite{Chen-book,Bialynicki-Birula1970,MSSBS2005}. In order to determine the contribution from the pair-plasma on the propagation properties in pulsar and magnetar atmospheres, we adopt the Goldreich-Julian expression for the plasma density \\cite{Beskin-book}, and evaluate the dispersion relation for field strengths in the pulsar and magnetar range, $B_{0}\\sim 10^{8}-10^{10}% \\,\\mathrm{T}$. In the radio-wave regime it then turns out that for one of the EM-wave polarizations, the plasma effects are typically negligible as compared to the QED-effects, whereas for the other polarization, the opposite is true in most cases. Noting that important processes in pulsar and magnetar environments, e.g.\\ photon splitting, typically involve both EM-wave polarizations, we will conclude that QED and plasma effects should be simultaneously included when studying radio wave propagation in such environments. ",
        "conclusions": "QED-effects associated with the external magnetic field are likely to be of importance in environments with extreme magnetic fields, in particular in the vicinity of astrophysical objects like pulsars and magnetars. For example, the radio silence of magnetars is assumed to be connected with QED-effects associated with the magnetar fields \\cite{Baring-Harding}, which could reach $10^{10} - 10^{11}\\,\\mathrm{T}$ $\\ $\\ close to the surface. However, in the same environments, we also expect the presence of an electron-positron plasma \\cite{Beskin-book}. Thus we evaluate (\\ref{Full-DR}) with $% \\sum_{s}=\\sum_{e,p}$, where $e$ and $p$ denotes electrons and positrons, respectively. Considering propagation at an arbitrary angle to the magnetic field in an electron-positron plasma, letting $\\omega _{\\mathrm{p}e,\\mathrm{p}p}\\!\\sim\\! \\omega\\! \\ll \\! \\left| \\omega _{\\mathrm{c}e,\\mathrm{c}p}\\right| $ , using $\\xi\\! \\ll\\! 1$ and noting that the factor $[\\sum_{e,p}\\omega _{\\mathrm{p}s}^{2}\\omega _{\\mathrm{c}s}/\\omega (\\omega ^{2}-\\omega _{\\mathrm{c}s}^{2})]^{2}$ then becomes negligibly small (due to the approximate cancellation of the electron and positron contributions), we find from (% \\ref{Full-DR}) that the dispersion relation separates in two modes that can be approximated by \\begin{equation} 1-n^{2}+8\\xi n_{\\bot }^{2}+\\frac{\\omega _{\\mathrm{p}}^{2}}{\\omega _{\\mathrm{c}}^{2}(1-4\\xi )}% \\approx 0  \\label{Magnetosonic-DR} \\end{equation} and \\begin{equation} (1-n^{2})(1-4\\xi ) - \\left( -14\\xi +\\frac{\\omega _{\\mathrm{p}}^{2}}{\\omega ^{2}}% \\right) (1-n_{\\Vert }^{2})  \\approx 0 , \\label{O-mode-DR} \\end{equation} where $\\omega _{\\mathrm{p}}=(\\omega _{\\mathrm{p}e}^{2}+\\omega _{\\mathrm{p}p}^{2})^{1/2}$ is the total plasma frequency, and $\\omega _{\\mathrm{c}}$ $=eB_{0}/m$ is the magnitude of the electron (or positron) cyclotron frequency. In the vicinity of pulsars or magnetars where $\\omega _{\\mathrm{c}}\\sim 10^{19}\\! -\\! 10^{21}\\,\\mathrm{rad\\,s^{-1},}$ the last term of (\\ref{Magnetosonic-DR}) is negligible unless the plasma density is extremely high. Omitting that term, the dispersion relation (\\ref {Magnetosonic-DR}) is then the same as that used in \\cite {Bialynicki-Birula1970} for the high phase velocity mode when considering photon-splitting. Similarly the ordinary mode described by (\\ref {O-mode-DR}), reduces to the mode with the lower phase velocity of reference \\cite {Bialynicki-Birula1970} when the plasma is removed. However, for the latter dispersion relation we note that a relatively modest plasma density is enough to significantly affect the propagation properties in the radio wave regime. To make a concrete estimate, we adopt the Goldreich--Julian density \\begin{equation} n_{\\mathrm{GJ}}=7\\times 10^{15}\\left(\\frac{0.1}{\\tau }\\right)\\left(\\frac{B_{\\mathrm{pulsar}}}{10^{8}}\\right)\\,\\mathrm{m}^{-3} \\label{Julian-Goldreich} \\end{equation} where $\\tau $ is the pulsar period time (in seconds) and $B_{\\mathrm{pulsar}}$ the pulsar magnetic field (in tesla). The pair plasma density is expected to satisfy $% n_{e}=n_{p}=Mn_{\\mathrm{GJ}}$, where $M$ is the multiplicity \\cite {Beskin-book,Luo-etal}. Moderate estimates then give $M=10$ \\cite{Luo-etal}. Choosing this value and letting $\\tau =1\\,\\mathrm{s}$, we note that for magnetar field strengths, $B_{\\mathrm{pulsar}}=10^{10}\\,\\mathrm{T}$, the term due to the plasma $\\propto \\omega _{\\mathrm{p}}^{2}/\\omega ^{2}$ in (\\ref{O-mode-DR}) dominates over the term due to QED $\\propto 14\\xi $ for frequencies up to $% \\omega \\sim 10^{14} - 10^{15}\\,\\mathrm{rad\\,s^{-1},}$ i.e.\\ in the infrared regime and below. Furthermore, we note that photon splitting \\cite {Adler,Bialynicki-Birula1970} as described by standard QED (i.e.\\ with zero plasma density) requires that the phase velocity of the dispersion relation in (% \\ref{Magnetosonic-DR}) is higher than that of (\\ref{O-mode-DR}). While this is always true in the absence of a plasma, we note that for wave frequencies in the infra-red regime and below, the Goldreich--Julian density given by (\\ref{Julian-Goldreich}) is enough to increase the phase velocity of the mode in (\\ref{O-mode-DR}) above that of (\\ref{Magnetosonic-DR}), unless we choose the period time $\\tau $ extremely low. \\ Thus we conclude that photon--photon splitting as described by vacuum theories is not likely to apply to magnetar atmospheres, unless the pair-production \\cite {Beskin-book} responsible for the Goldreich-Julian expression is effectively suppressed. Wave cascade processes as a mechanism to explain the radio silence of magnetars \\cite{Baring-Harding} could still be possible, but for densities of the order of (\\ref{Julian-Goldreich}), plasma nonlinearities are likely to dominate over the pure QED effects."
    },
    "0601/astro-ph0601061_arXiv.txt": {
        "abstract": "We present the first MIR spectrum of the $z=2.2856$ ultraluminous, infrared galaxy FSC 10214+4724, obtained with the Infrared Spectrograph onboard the {\\it Spitzer Space Telescope}.  The spectrum spans a rest wavelength range of $2.3-11.5\\mu$m, covering a number of key diagnostic emission and absorption features.  The most prominent feature in the IRS spectrum is the silicate emission at rest-frame $\\sim 10$\\ \\mic.  We also detect an unresolved emission line at a rest wavelength of $7.65\\mu$m which we identify with [NeVI], and a slightly resolved feature at 5.6 \\mic\\ identified as a blend of [Mg VII] and [Mg V].  There are no strong PAH emission features in the FSC 10214+4724 spectrum.  We place a limit of $0.1\\mu$m on the equivalent width of $6.2\\mu$m PAH emission but see no evidence of a corresponding 7.7 \\mic\\ feature.  Semi-empirical fits to the spectral energy distribution suggest $\\sim 45$\\% of the bolometric luminosity arises from cold ($\\sim 50$\\ K) dust, half arises from warm (190 K) dust, and the remainder, $\\sim 5\\%$ originates from hot ($\\sim 640$\\ K) dust.  The hot dust is required to fit the blue end of the steep MIR spectrum.  The combination of a red continuum, strong silicate emission, little or no PAH emission, and no silicate absorption, makes FSC 10214+4724 unlike most other ULIRGs or AGN observed thus far with IRS.  These apparently contradictory properties may be explained by an AGN which is highly magnified by the lens, masking a (dominant) overlying starburst with unusually weak PAH emission. ",
        "introduction": "FSC 10214+4724, at a redshift of $z=2.2856$\\ \\citep{Rowan-Robinson 1991}, was initially thought to be the most luminous object in the Universe, but was later revealed to be gravitationally lensed by a foreground galaxy \\citep{Broadhurst 1995, Graham 1995, Serjeant 1995, Eisenhardt 1996}.  The lensing model of \\cite{Eisenhardt 1996}\\ suggests that the central (optical/UV) source is magnified by a factor of $\\sim 100$, and that the lensing arc is an image of the central $\\sim 0^{\\prime \\prime}.005$\\ (40 pc) of the source at B-band rest wavelength.  They further conclude that the bolometric luminosity of the source is produced over a larger region (240 pc), implying a bolometric magnification factor of 30 and an intrinsic luminosity of $\\sim 2 \\times 10^{13}$\\ \\Lsun\\ , placing it in the class of ultraluminous infrared galaxies (ULIRGs).  The large magnification makes FSC 10214+4724 a unique subject for the study of high redshift ULIRGs, offering greater effective sensitivity and spatial resolution than possible in observations of unlensed sources.  A central question in understanding FSC 10214+4724 is the relative contribution of starburst and AGN components to its luminosity.  The rest-frame UV-optical spectrum appears to be similar to that of Seyfert 2 galaxies \\citep{Elston 1994}.  Strong UV polarization shows that much of the UV continuum results from scattered light, and broad lines in the polarization spectrum identify the presence of an AGN \\citep{Lawrence 1993, Goodrich 1996}. The bulk of the luminosity, however, is emitted in the infrared \\citep[as much as 99\\%,][]{Rowan-Robinson 1991}.  CO observations point to a large reservoir of molecular gas \\citep{Solomon 1992, Scoville 1995}. Sub-mm and millimeter detections \\citep{Downes 1992, Rowan-Robinson 1993,Benford 1999} suggest substantial emission from cold dust, usually associated with extended star formation.  Recent {\\it Chandra}\\ observations show the object to have weak X-ray emission, consistent with vigorous star formation or a Compton-thick AGN \\citep{Alexander 2005}.  Lens models suggest differential amplification, so that the central (AGN-dominated) region is more highly magnified than the surrounding starburst \\citep[e.g.,][]{Eisenhardt 1996, Lacy 1998}.  The sensitivity and large wavelength coverage of the {\\it Spitzer}\\ Infrared Spectrograph \\citep[IRS][]{Houck 2004}\\ makes it possible to explore the dust emission and absorption features in the rest-frame mid-infrared spectra of dusty galaxies at low and high redshift, and thus identify the power sources which may be hidden in the UV and optical.  In this paper, we present the first MIR spectrum of FSC 10214+4724.  We describe the observations and data reduction in Section 2, present the spectrum and a dust emission model fit to the spectral energy distribution in Section 3, and discuss the implications in Section 4.  Throughout, we assume a $\\Lambda$-dominated flat universe, with $H_0=70$\\ km s$^{-1}$\\ Mpc$^{-1}$, $\\Omega_{\\Lambda}=0.7, \\Omega_{m}=0.3$. ",
        "conclusions": "The shape of the SED of FSC 10214+4724 is dominated by dust at a variety of temperatures.  Substantial cold dust must be present, given the strong emission at long wavelengths \\citep[$> 40$\\ \\mic, rest-frame; e.g., ][]{Rowan-Robinson 1993}. At the same time, warmer dust will be required to explain the 5-10 \\mic\\ continuum.  Dust warmer than 100 K can produce the silicate emission \\citep[e.g.,][]{Li 2001}, but a hot dust component (several hundred K) is required to explain the $\\sim 5$\\ \\mic\\ continuum.  The rest-frame near-infrared (NIR) will have a contribution from both star light and hot dust.   We have fit the SED with a multi-component model which includes three graphite and silicate dust grain components and a 3500K blackbody stellar component \\citep{Marshall 2006}. The three dust components in the model are not meant to represent three distinct physical structures, but rather they are indicative of the average temperature ranges within the source.  The model was fit to the SED from observed frame NIR to millimeter wavelengths (see Figure \\ref{fig: SED fit}.  NIR data included the photometry of \\cite{Soifer 1991}\\ and the IRAC data described in Section 2.  Longer wavelength data included IRAS photometry at 60 and 100 \\mic, the 350 \\mic\\ detection of \\cite{Benford 1999}, and the sub-mm and mm data of \\cite{Rowan-Robinson 1993}\\ and \\cite{Downes 1992}.   A minimum of three dust components are required to fit the FSC 10214+4724 SED, one at $\\sim 50$\\ K (cold), one at $\\sim 190$\\ K (warm), and one at $\\sim 640$\\ K (hot). Their relative contributions to the bolometric luminosity are given in Table \\ref{tab: SED fit} . Note that hot dust contribution is an upper limit, because it is dominated by the IRAC data and no correction has been made for contamination by the foreground lensing galaxy. The stellar emission is a negligible contribution to the bolometric luminosity, but is needed to fit the shortest wavelength data. component.  The cold dust component ($51 \\pm 6$\\ K) is in good agreement with the estimate of \\cite{Benford 1999}, 55 K.  Each component contains a distribution of grains with different equilibrium temperatures depending on their size and composition.  Additionally, each component contains emission from dust at different radial distances, and therefore temperatures, from the illuminating source. We define the characteristic temperature of a component to be the temperature of the most luminous grain size at the distance from the source contributing the majority of the luminosity.  This luminosity dominating distance corresponds to a $\\tau(UV) \\sim 0.5$, where approximately half of the UV-source photons have been absorbed.  Each component therefore contains dust above and below the characteristic temperature.  With this definition, the characteristic temperature roughly corresponds to the expected peak in the dust modified (grey-body) Planck function.  The observed SED of FSC 10214+4724 appears consistent with low redshift, AGN-dominated ULIRGs.  Such sources can have cold dust components which account for up to 40\\%\\ of the bolometric luminosity, due to both AGN-heated cold dust far from the nucleus and the presence of a small underlying starburst \\citep{Armus 2004, Armus 2006}. \\cite{Rowan-Robinson 2000} also estimates that the AGN contributes most of the luminosity of FSC 10214+4724.  In addition, FSC 10214+4724 shows silicate emission, similar to other AGN observed with the IRS, and a hot (graphite) dust continuum.  However, the hot dust contribution is quite small compared to many AGN-dominated ULIRGs or QSOs.  The only obvious fine-structure lines are [Ne VI], [Mg VII] and [Mg V], high-ionization species not observed in starburst galaxies. There is little if any PAH emission.  The upper limit to the $6.2\\mu$m EW is approximately a factor of $5-10$ lower than most pure starburst galaxies \\citep{Brandl 2005}\\ and ULIRGs dominated by star formation \\citep{Armus 2004, Armus 2006}.  However, differential magnification is likely to be enhancing the central AGN, reducing the PAH EW and making the silicate emission more obvious in the IRS spectrum.  \\cite{Eisenhardt 1996}\\ estimate a magnification factor of $\\sim 100$\\ for the central 40 pc, but only a factor of $\\sim 30$\\ out to at least 240 pc.  The intrinsic contribution of the starburst to the bolometric luminosity could be larger than that inferred by our current model.  We assume that the warm and hot dust heated by the AGN are magnified by an additional factor of 3.3, but that both the cold dust heated by the AGN and any dust heated by starburst are not.  The ratio of cold to warm dust in local AGN- and starburst-dominated ULIRGs is approximately 2:3 and 3:1, respectively (Armus et al. 2006, in preparation).  Taking these assumptions together, we estimate the intrinsic AGN contribution to the bolometric luminosity of FSC 10214+4724 to be $\\sim 35$\\% after correction for the differential magnification.  If the differential magnification factor is correct, this AGN contribution is probably an upper limit because AGN-dominated ULIRGs often have measurable starbursts.   If the starburst contributes 65\\%\\ of the intrinsic luminosity, it is surprising that little or no PAH emission is seen.  Taking our limit of 0.1 \\mic\\ $>EW_{\\mbox{rest}}$\\ and a differential magnification factor of 3.5 would place the limit of the EW for PAH emission within the range of other star-forming ULIRGs\\citep[Armus et al.  2005, in prep;][]{Brandl 2005}.  However, the uncertainty in our PAH measurement leaves this possibility unconfirmed; a lower signal to noise ratio would have made our limit higher.  Furthermore, the lack of evidence for comparably strong PAH emission at 7.7 \\mic\\ makes it likely that the EW at 6.2 \\mic\\ is overestimated.  Nonetheless, the limit demonstrates that substantial star formation can be present in FSC 10214+4724 and not visible in the (differentially magnified) IRS spectrum.  In figure \\ref{fig: SED compare}, we compare the rest frame MIR spectrum of FSC 10214+4724 to the IRS spectra of three other sources.  The intrinsic spectrum of FSC 10214+4724 is redder than the observed spectrum, given the differential magnification, so the differences seen in the figure would be greater if corrected for lensing.  The comparison sources are an AGN-dominated ULIRG (FSC 15307+3252), a silicate-emitting QSO (PG 1351+640) and a pure starburst galaxy (NGC 7714).  PG 1351+640 has similar silicate emission, but has a much stronger hot dust continuum at $\\sim 5$\\ \\mic\\ \\citep{Hao 2005}.  In fact, the cold dust emission in FSC 10214+4724 gives it a much steeper mid-to-far infrared slope than most AGN with silicate emission.  The 30 \\mic\\ to 6 \\mic\\ flux density ratio of 46 in FSC 10214+4724, is higher than the reddest of the quasars in \\cite{Hao 2005}\\ or \\cite{Siebenmorgen 2005}, with a ratio of 11, and the reddest AGN in \\cite{Weedman 2005}, NGC 1275, which has a ratio of 30.  The ULIRG FSC 15307+325 has a similar MIR slope to FSC 10214+4724 in the 2-10 \\mic\\ region, but has silicate absorption rather than emission, and a much weaker cold dust continuum at longer wavelengths.  The continuum of NGC7714 is very red with a 30 \\mic\\ to 6 \\mic\\ flux density ratio of 120, and a strong stellar contribution at the short wavelengths.  In many AGN-dominated ULIRGs, the MIR spectrum is even bluer, with a strong continuum in the $\\sim 5$\\ \\mic\\ region \\citep[e.g., ][]{Laurent 2000}.  ULIRGs with steeper mid-infrared continua tend also to have silicate absorption, not emission.  We see no evidence in the spectrum of FSC 10214+4724 for an underlying silicate absorption feature at 9.7 \\mic. Furthermore, the profile of the silicate emission is consistent with that seen in most other AGN observed with the IRS \\citep{Hao 2005, Siebenmorgen 2005, Weedman 2005}; although see \\cite{Sturm 2005}\\ for a counter example and discussion of the factors influencing the profile shape.  While some absorption may be filled in with emission, the absorption and emission profiles are not necessarily identical.  A coincidence of opacity, temperature, and grain mixture in the emitting and absorbing regions would be required.  The AGN contribution to the bolometric luminosity in FSC 10214+4724 is apparently substantial but does not dominate.  Nonetheless, the PAH emission expected for a starburst is weak or absent.  An unusual dust geometry would be required for an AGN alone to explain the SED, given the amount of cold dust.  The other available evidence points to a standard AGN configuration.  The shape of the spectrum between $2-12\\mu$m is qualitatively similar to dusty torus models \\citep{Pier and Krolik 1992}\\ with inner radius to height ratios of $a/h\\sim 0.3$, and opening angles of $\\sim 50$\\deg -- consistent with a high-luminosity equivalent of Seyfert 2 galaxy, like NGC 1068. The similarlity to NGC 1068 hass previously been noted\\citep[e.g.,][]{Barvainis 1995}.  The presence of silicate emission in the IRS spectrum rules out models where the torus is seen edge-on, or where the AGN is completely obscured by foreground (cold) dust."
    },
    "0601/astro-ph0601582_arXiv.txt": {
        "abstract": "We present 350 $\\mu$m photometry of all 17 galaxy candidates in the Lockman Hole detected in a 1.1 mm Bolocam survey.  Several of the galaxies were previously detected at 850 $\\mu$m, at 1.2 mm, in the infrared by {\\it Spitzer}, and in the radio. Nine of the Bolocam galaxy candidates were detected at 350 $\\mu$m and two new candidates were serendipitously detected at 350 $\\mu$m (bringing the total in the literature detected in this way to three).  Five of the galaxies have published spectroscopic redshifts, enabling investigation of the implied temperature ranges and a comparison of photometric redshift techniques.  $\\lambda$ = 350 $\\mu$m lies near the spectral energy distribution peak for $z \\approx 2.5$ thermally emitting galaxies.  Thus, luminosities can be measured without extrapolating to the peak from detection wavelengths of $\\lambda \\ge$ 850 $\\mu$m. Characteristically, the galaxy luminosities lie in the range $1.0-1.2\\times10^{13}$ L$_\\odot$, with dust temperatures in the range of 40 K to 70 K, depending on the choice of spectral index and wavelength of unit optical depth. The implied dust masses are 3 - 5 $\\times$ $10^8$ M$_\\odot$. We find that the far-infrared to radio relation for star-forming ULIRGs systematically overpredicts the radio luminosities and overestimates redshifts on the order of $\\Delta z \\approx 1$, whereas redshifts based on either on submillimeter data alone or the 1.6 $\\mu$m stellar bump and PAH features are more accurate. ",
        "introduction": "Surveys at submillimeter and millimeter wavelengths have detected hundreds of galaxy candidates by their thermal dust emission.  The galaxies (hereafter referred to as submillimeter galaxies) characteristically have redshifts $z > 1$ and inferred luminosities of $L \\sim 10^{13}$ L$_\\odot$ and star formation rates of $10^3$ M$_\\odot$ per year (assuming dust heating by young stars).  Such enormous luminosities and star formation rates, or accretion rates in the case of super-massive black hole growth, make submillimeter galaxies strong candidates for the progenitors of massive galaxies at the current epoch.  Clearly, it is crucial to characterize the spectral energy distributions (SEDs) where their emission peaks ($\\lambda$ = 350 $\\mu$m for 40 K dust at a redshift of $z = 2.5$), measure their redshifts and luminosity functions, determine their power sources, and integrate them into theories of galaxy formation.  Although submillimeter galaxy SEDs peak at a few hundred microns for all but the highest redshifts, $\\lambda \\ge$ 850 $\\mu$m surveys have been most successful at detecting galaxies because of the lower atmospheric noise and greater transmission, and less stringent telescope surface requirements.  Most of the detections have been low signal-to-noise ratio (just over thresholds of 3-4 $\\sigma$), necessitating multiwavelength confirmation.  Furthermore, the SEDs have been extrapolated shortward from 850-1200 $\\mu$m over the peak, or between 850-1200 $\\mu$m and the far-infrared, to estimate dust temperatures, luminosities, and star formation rates. Clearly, 350 $\\mu$m photometry can confirm galaxy candidates and sample the SEDs near their peaks for more precise inferences of physical parameters. Similarly, because of the difficulty in obtaining spectroscopic redshifts of large numbers of highly obscured galaxies, various photometric redshift estimation techniques have emerged, notably based on the far-infrared to radio luminosity relation in ULIRGs \\citep{carilli99,yun02} and the stellar continuum bump in the infrared for {\\it Spitzer}-detected galaxies \\citep{egami04,sawicki02}.  However, despite the difficulty, candidate spectroscopic redshifts have been obtained for $\\sim$ 73 galaxies \\citep{chapman05}.  Thus, with well-determined dust-emission SEDs, including 350 $\\mu$m, photometric techniques can be compared to spectroscopic redshifts.  In this paper, we present 350 $\\mu$m photometry of all 17 submillimeter galaxy candidates from the Bolocam Lockman Hole survey \\citep{laurent05}, some with previous 850 $\\mu$m and 1200 $\\mu$m detections.  Two new galaxy candidates were serendipitously detected at 350 $\\mu$m, bringing the total number of 350 $\\mu$m-discovered galaxies to three \\citep{khan05}.  We combine these data with infrared and radio data to derive improved luminosities, explore the range of implied dust temperatures and spectral indices, and compare photometric redshift techniques.  This comparison is timely with the imminent launch of the {\\it Herschel Space Observatory} (scheduled for August 2007), which will detect thousands of galaxies at far-infrared and submillimeter wavelengths, but for which spectroscopic redshifts will be attainable for only a small fraction.  Throughout the paper, a cosmology of H$_0 = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_M = 0.3$, and $\\Omega_\\Lambda$ = 0.7 is assumed. ",
        "conclusions": "We have obtained 350 $\\mu$m SHARC II observations toward galaxy candidates from the Bolocam Lockman Hole survey.  The Lockman Hole has rich, deep, multiwavelength observations enabling detailed studies of galaxies.  The 350 $\\mu$m photometry is near the peaks of the SEDs of galaxies with characteristic temperatures of $\\sim$ 50 K and redshifts of $z$ $\\sim$ 2 to 3.  They therefore enable measurements of luminosities and estimates of temperatures and photometric redshifts without interpolating over the peak of the FIR thermal SEDs.  Seven galaxies detected at 1.1 mm with Bolocam were detected at 350 $\\mu$m, two of which have two 350 $\\mu$m counterparts; these were combined with two 350 $\\mu$m detections from the survey of \\cite{kovacs05}, bringing the total number of Bolocam galaxies detected with SHARC II to nine.  Two additional galaxies not associated with the Bolocam sources were also detected.  The SHARC II detections range in significance from 3.0 $\\sigma$ to 6.8 $\\sigma$, with flux densities ranging from 14 mJy to 64 mJy.  We combined our observations with 850 $\\mu$m and 1.2 mm photometry from the literature to fit the submillimeter/millimeter-wave spectra to thermal dust models.  We found that two models with significantly different dust temperatures (40 K and 68 K) and spectral indices $\\beta$ (1.6 and 1.8, respectively) yielded similar quality fits owing to the degeneracy in T and $\\beta$, rendering them indistinguishable without better SED sampling.  However, there is little consequence of the degeneracy to the derived luminosities, photometric redshifts, and dust masses within the statistical uncertainties.  Five of the galaxies have spectroscopic redshifts in the literature, with four ranging from $z$ = 2.1 to 3.0 and one at $z$ = 0.689.  The four high-z galaxies have luminosities of (1.0 - 1.2) $\\times$ 10$^{13}$ L$_\\odot$, while the $z$ = 0.689 galaxy is best fit by a 20 K, $\\beta$ = 1.0, spectrum with a much lower luminosity:  1.3 $\\times$ 10$^{11}$ L$_\\odot$.  (Given the source confusion in the optical and radio, along with consistent photometric redshifts, we suggest that the $z$ = 0.689 spectroscopic redshift of Bolocam source 17 may be a misidentification.)  The characteristic dust masses for the four high-$z$ spectroscopic galaxies are 4 $\\times$ 10$^8$ M$_\\odot$, implying gas masses of 4 $\\times$ 10$^{10}$ M$_\\odot$.  The dominant uncertainties in this estimation are the dust opacity and the gas-to-dust conversion factor, which make the estimation uncertain to a factor of a few.  Assuming a Salpeter IMF and that the submillimeter emission derives completely from star formation yields stellar masses of 10$^{10}$ to a few times 10$^{11}$ M$_\\odot$, broadly consistent with the stellar content of modern-day elliptical galaxies.  The photometric redshifts of the full sample of seven galaxies span the range of $z$ = 2.0 to $z$ = 4.3, with statistical uncertainties of $\\Delta z$ = 0.3 to 0.6 (1 $\\sigma$).  Photometric redshifts utilizing composite radio/FIR spectra representative of local star-forming ULIRGs yields systematically higher redshifts, on the order of $\\Delta z$ = 1.  For the four galaxies with optical spectroscopic redshifts the anomolously high redshifts arise from systematically low 1.4 GHz observed flux densities.  The discrepancy could arise from small number statistics, inverse-Compton losses of high energy cosmic rays off the CMB, heating by radio-quiet AGN, or suppressed synchrotron emission from supernova remnants in the unusually luminous galaxies.  For comparison, photometric redshifts derived using only Spitzer MIPS and IRAC data points yielded slightly more precise and accurate redshifts than the submillimeter/millimeter-wave data alone, with discriminatory power between heating by AGN and star formation (albeit with limited bolometric luminosity constraints)."
    },
    "0601/gr-qc0601026_arXiv.txt": {
        "abstract": "We present results from fully nonlinear simulations of unequal mass binary black holes plunging from close separations well inside the innermost stable circular orbit with mass ratios $q \\equiv M_1/M_2 = \\lbrace 1,0.85,0.78,0.55,0.32 \\rbrace$, or equivalently, with reduced mass parameters $\\eta \\equiv M_1M_2/(M_1+M_2)^2 = \\lbrace 0.25, 0.248, 0.246, 0.229, 0.183 \\rbrace$. For each case, the initial binary orbital parameters are chosen from the Cook-Baumgarte equal-mass ISCO configuration. We show waveforms of the dominant $\\ell=2,3$ modes and compute estimates of energy and angular momentum radiated. For the plunges from the close separations considered, we measure kick velocities from gravitational radiation recoil in the range \\kickrange. Due to the initial close separations our kick velocity estimates should be understood as a lower bound. The close configurations considered are also likely to contain significant eccentricities influencing the recoil velocity. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601257_arXiv.txt": {
        "abstract": "We use optical data on 10 Kuiper Belt objects (KBOs) to investigate their rotational properties. Of the 10, three (30\\%) exhibit light variations with amplitude $\\Delta m \\ge 0.15\\,$mag, and 1 out of 10 (10\\%) has $\\Delta m \\ge 0.40\\,$mag, which is in good agreement with previous surveys.  These data, in combination with the existing database, are used to discuss the rotational periods, shapes, and densities of Kuiper Belt objects. We find that, in the sampled size range, Kuiper Belt objects have a higher fraction of low amplitude lightcurves and rotate slower than main belt asteroids. The data also show that the rotational properties and the shapes of KBOs depend on size. If we split the database of KBO rotational properties into two size ranges with diameter {\\em larger} and {\\em smaller} than $400\\,$km, we find that: (1) the mean lightcurve amplitudes of the two groups are different with 98.5\\% confidence, (2) the corresponding power-law shape distributions seem to be different, although the existing data are too sparse to render this difference significant, and (3) the two groups occupy different regions on a {\\em spin period} vs. {\\em lightcurve amplitude} diagram. These differences are interpreted in the context of KBO collisional evolution. ",
        "introduction": "The Kuiper Belt (KB) is an assembly of mostly small icy objects, orbiting the Sun beyond Neptune. Kuiper Belt objects (KBOs) are likely to be remnants of outer solar system planetesimals \\citep{1993Natur.362..730J}.  Their physical, chemical, and dynamical properties should therefore provide valuable information regarding both the environment and the physical processes responsible for planet formation.  At the time of writing, roughly 1000 KBOs are known, half of which have been followed for more than one opposition. A total of $\\approx 10^{5}$ objects larger than 50~km are thought to orbit the Sun beyond Neptune \\citep{2000prpl.conf.1201J}.  Studies of KB orbits have revealed an intricate dynamical structure, with signatures of interactions with Neptune (Malhotra 1995). The size distribution follows a differential power-law of index $q=4$ for bodies $\\gtrsim 50\\,$km \\citep{2001AJ....122..457T}, becoming slightly shallower at smaller sizes \\citep{2004AJ....128.1364B}.   KBO colours show a large diversity, from slightly blue to very red \\citep{1996AJ....112.2310L,2000Natur.407..979T,2001AJ....122.2099J}, and seem to correlate with inclination and/or perihelion distance \\citep[e.g.,][]{2001AJ....122.2099J,2002AJ....124.2279D, 2002ApJ...566L.125T}. The few low-resolution optical and near-IR KBO spectra are mostly featureless, with the exception of a weak 2$\\,\\mu$m water ice absorption line present in some of them \\citep{1999ApJ...519L.101B,2001AJ....122.2099J}, and strong methane absorption on 2003$\\,$UB$_{313}$ \\citep{2003UB313}.  About 4\\% of known KBOs are binaries with separations larger than 0\\farcs15 \\citep{2002AJ....124.3424N}. All the observed binaries have primary-to-secondary mass ratios $\\approx$ 1. Two binary creation models have been proposed. \\citet{2002Icar..160..212W} favours the idea that binaries form in three-body encounters. This model requires a 100 times denser Kuiper Belt at the epoch of binary formation, and predicts a higher abundance of large separation binaries. An alternative scenario \\citep{2002Natur.420..643G}, in which the energy needed to bind the orbits of two approaching bodies is drawn from the surrounding swarm of smaller objects, also requires a much higher density of KBOs than the present, but it predicts a larger fraction of close binaries. Recently, \\citet{2004AJ....127.3023S} have shown evidence that $2001\\,\\rm{QG}_{298}$ could be a close or contact binary KBO, and estimated the fraction of similar objects in the Belt to be $\\sim10\\%$--$20\\%$.  Other physical properties of KBOs, such as their shapes, densities, and albedos, are still poorly constrained.  This is mainly because KBOs are extremely faint, with mean apparent red magnitude $m_R\\sim$23 \\citep{2001AJ....122.2740T}.  The study of KBO rotational properties through time-series broadband optical photometry has proved to be the most successful technique to date to investigate some of these physical properties.  Light variations of KBOs are believed to be caused mainly by their aspherical shape: as KBOs rotate in space, their projected cross-sections change, resulting in periodic brightness variations.  One of the best examples to date of a KBO lightcurve -- and what can be learned from it -- is that of (20000)$\\,$Varuna \\citep{2002AJ....123.2110J}. The authors explain the lightcurve of (20000)$\\,$Varuna as a consequence of its elongated shape (axes ratio, $a/b\\sim 1.5$). They further argue that the object is centripetally deformed by rotation because of its low density, ``rubble pile'' structure. The term ``rubble pile'' is generally used to refer to gravitationally bound aggregates of smaller fragments. The existence of rubble piles is thought to be due to continuing mutual collisions throughout the age of the solar system, which gradually fracture the interiors of objects. Rotating rubble piles can adjust their shapes to balance centripetal acceleration and self-gravity. The resulting equilibrium shapes have been studied in the extreme case of fluid bodies, and depend on the body's density and spin rate \\citep{1969efe..book.....C}.  \\citeauthor{2003Icar..161..174L} (\\citeyear{2003Icar..161..174L}, hereafter \\citetalias{2003Icar..161..174L}) showed that under reasonable assumptions the fraction of KBOs with detectable lightcurves can be used to constrain the shape distribution of these objects. A follow-up (\\citealt{2003EM&P...92..221L}, hereafter \\citetalias{2003EM&P...92..221L}) on this work, using a database of lightcurve properties of 33 KBOs \\citep[]{2002AJ....124.1757S,2003EM&P...92..207S}, shows that although most Kuiper Belt objects ($\\sim85\\%$) have shapes that are close to spherical ($a/b\\leq1.5$) there is a significant fraction ($\\sim12\\%$) with highly aspherical shapes ($a/b\\geq1.7$).  In this paper we use optical data on 10 KBOs to investigate the amplitudes and periods of their lightcurves. These data are used in combination with the existing database to investigate the distributions of KBO spin periods and shapes. We discuss their implications for the inner structure and collisional evolution of objects in the Kuiper Belt. ",
        "conclusions": ""
    },
    "0601/astro-ph0601492_arXiv.txt": {
        "abstract": "{Laser guide stars with adaptive optics allow astronomical image correction in the absence of a natural guide star. Single guide star systems with a star created in the earth's sodium layer can be used to correct the wavefront in the near infrared spectral regime for 8-m class telescopes. For possible future telescopes of larger sizes, or for correction at shorter wavelengths, the use of a single guide star is ultimately limited by focal anisoplanatism that arises from the finite height of the guide star. To overcome this limitation we propose to overlap coherently pulsed laser beams that are expanded over the full aperture of the telescope, traveling upwards along the same path which light from the astronomical object travels downwards. Imaging the scattered light from the resultant interference pattern with a camera gated to a certain height above the telescope, and using phase shifting interferometry we have found a method to retrieve the local wavefront gradients. By sensing the backscattered light from two different heights, one can fully remove the cone effect, which can otherwise be a serious handicap to the use of laser guide stars at shorter wavelengths or on larger telescopes. Using two laser beams multiconjugate correction is possible, resulting in larger corrected fields. With a proper choice of laser, wavefront correction could be expanded to the visible regime and, due to the lack of a cone effect, the method is applicable to any size of telescope. Finally the position of the laser spot could be imaged from the side of the main telescope against a bright background star to retrieve tip-tilt information, which would greatly improve the sky coverage of the system. ",
        "introduction": "Adaptive optics correction of wavefront distortions is becoming a common tool at large ground based telescopes. Diffraction limited imaging through the turbulent atmosphere can be regarded as routine at near infrared wavelengths and shows great success in many scientific areas (see for example the proceedings of the `Science with adaptive optics' meeting, eds. Brandner \\& Kasper \\cite{Brandner05}). By using laser guide stars this technique can be applied at almost any position in the sky where a suitable star for tip-tilt correction can be found. Nevertheless, the quality of correction when using a single guide star is limited by focal anisoplanatism, or the `cone effect'. Due to the limited height where the artificial guide star is created, some part of the wavefront distortion remains unsampled, and thus uncorrected. For possible future extremely large telescopes, or for correction at short wavelengths, this limit will ultimately prevent diffraction limited performance. To overcome this, tomographic solutions have been proposed (e.g. by Angel \\& Lloyd-Hart \\cite{Angel00a}), utilizing multiple Rayleigh guide stars or multiple stars in the earth's sodium layer. For a large telescope or for correction in the visible a large number of guide stars and wavefront sensors are needed.  Another approach is to use the upward path of a projected laser beam to sample the atmospheric turbulence, either in an inverse Shack-Hartmann geometry by distributing many laser beams in parallel over the telescope, or with a single laser beam expanded over the whole aperture of the telescope. Buscher et al. \\cite{Buscher2002} have proposed a method to sample the local curvature of the atmospheric wavefront distortions from the intensity changes that a collimated laser beam shows when it has traveled to a certain height. In the same publication a method is proposed to compare the full aperture expanded beam with lots of other laser beams in a point-diffraction scheme. Also using a full aperture expanded laser, Bonaccini et al. \\cite{Bonaccini2004} have proposed to measure the wavefront on the downward path when imaging back the distributed beacon with a shearing interferometer. Feasibility studies of all the above mentioned methods are currently being carried out. Limitations for the inverse Shack-Hartmann and the point-diffraction scheme may arise due to both the size of the individual spots that can be reached in the presence of diffraction and also beam quality aspects. For the shearing measurement, the coherence length of the scattered light might be a serious issue.  The method proposed here uses the coherent superposition of tilted laser wavefronts over the whole aperture of the telescope together with the method of laser phase shifting interferometry (LPSI) in order to retrieve a local phase difference that can be used in an adaptive optics system in a similar way to the gradients retrieved from a standard Hartmann sensor. To adapt to different conditions, the sensing resolution can be chosen freely by adding an appropriate tilt to the beams. Detecting the light scattered from low altitudes will result in a ground layer correction. By extending the system to multiple detections at different heights (or two beam systems) with more deformable mirrors, a complete sampling of the turbulent layers in the atmosphere is possible. Therefore high Strehl ratios can be reached and multiconjugate adaptive optics correction over an extended field of view is possible. Additional means can help to expand the sky coverage further, since in a differential measurement far off-axis guide stars can serve as the tip-tilt reference.   \\begin{figure} \\centering \\includegraphics[width=8cm]{4143Fig1.eps} \\caption{Basic principle of measuring atmospheric turbulence with phase shifting interferometry. A pulsed laser beam is split into two beams which are tilted with respect to each other. To one of the beams a known phase shift $\\alpha$ can be applied. The scattered fringe pattern which is formed in the atmosphere at height H, can be viewed with a camera over the full telescope aperture, which is gated according to the time of flight. In a series of laser shots a known phase difference is added to the beam in each pulse. Depending on the algorithm used for the phase retrieval, a set of three or more laser shots will allow the local phase difference to be measured. The frames shown here of a single laser shot are taken from a simulation with multiple turbulence layers of the fringe formation and imaging at an 8-m telescope with $r_0=10$\\,cm.}\\label{fig:basic_principle} \\end{figure} ",
        "conclusions": "We have proposed a method to measure and correct atmospheric distortions for astronomical telescopes using a laser to probe the turbulence. The use of laser wavefronts over the full aperture and the technique of phase shifting interferometry allows the distortions to be measured and corrected without suffering from the cone effect. The system outlined is therefore applicable to any size of telescope including extremely large apertures. While the number of actuators on the deformable mirror and the number of detection pixels on the wavefront sensor have to be adopted accordingly, the laser system itself would not change with telescope size. We have shown that algorithms for the interferometric phase retrieval exist and allow for a calibration of the measurement. The basic scheme outlined, involving a single detection of the scattered light, gives easy access to ground layer adaptive optics correction. By using of two deformable mirrors and two laser beams, a multiconjugate correction can be achieved. The necessary lasers are commercially available, and the power demands can be reduced with a dynamical refocusing device. The total system can be built relatively compactly, and the laser plus adaptive optics would plug into the telescope as one instrument. This removes the need for separate laser projectors, simplifying the implementation. The two examples given here of how to arrange the collection of the interferograms and the correction schemes, can easily be modified to fit the needs of an individual project. While the basic feasibility has been shown in this paper, further study is required in order to address detailed questions arising from the implementation or extension of such a system. Examples are field-of-view considerations in the multiconjugate case, studies for stray light suppression onto the science camera, and multi-color sensing for phase range extension."
    },
    "0601/astro-ph0601347_arXiv.txt": {
        "abstract": "%  There are obvious and direct ways in which the presence of binary companions may affect activity in OB stars through tidal interactions. In this review, however, I consider a more fundamental role for multiplicity and explore claims that the Be phenomenon may be intimately linked to binarity. I describe the binary channel for the formation of Be stars and the ongoing discussion about the relative contribution of this channel to the population of Be stars. I also present evidence suggesting that some environments are more favourable for the appearance of Be stars and explore whether this may be connected to initial conditions, such as the chemical composition or the distribution of rotational velocities on the ZAMS. ",
        "introduction": "%  As the presence of binary (or multiple) companions in a close orbit can affect the evolution of a star, it is only natural to assume that it can also have an effect on the occurrence of different kinds of stellar activity. In some cases, the physical mechanisms connecting the presence of a companion with activity are more or less transparent. For example, the tidal force exerted by a companion can excite different kinds of stellar oscillations, as discussed by \\citet{wil06}. Such induced oscillations are believed to occur in a few binaries containing OB stars \\citep[e.g.,][]{quaint}.  Similarly, the tidal force exerted by a companion on the disk of a Be star can result in photometric variability. For example, regularly spaced, small brightenings of the B6\\,IVe star HR 2968 have been interpreted as due to the presence of a companion \\citep{car99}. However, dedicated searches for periodic photometric variability locked with the orbital period in Be/X-ray binaries have rarely resulted in detections \\citep[cf.][]{roc93, clark99}, even though the presence of a neutron star in a relatively close orbit suggests that the tidal effect on the decretion disk would be rather strong in these systems \\citep[but see][for examples of positive detections in SMC Be/X-ray binaries]{ce04}.  In some Be stars, (quasi-)periodic variations in the shape of emission lines (V/R variability) occur with the orbital period of a companion (see, e.g., Stefl et al.\\ these proceedings). The presence of a companion is one obvious way of providing a clock when a clear periodicity exists.     There are, however, deeper ways in which membership in a stellar system may be connected to activity. Considering the focus of this meeting on the use of active stars as laboratories to study the effects of fast rotation, I will centre on Be stars and explore claims that binarity provides the mechanism to spin up Be stars to high rotational speeds. Afterwards, I will move to a higher level of multiplicity and discuss the likelihood that membership in an open cluster affects the chances of becoming a Be star. Finally, I will address the role that variations in metallicity may have in the occurrence of the Be phenomenon. ",
        "conclusions": "Though the study of large populations of early type stars in open clusters seems to be the only way to gain an understanding of the causes and incidence of the Be phenomenon, data so far has provided a very blurred and, in many aspects, contradictory picture  We know for certain that some fraction of Be stars must have undergone mass transfer in a binary system and we have reasons to suspect that their Be nature may be in some way due to this process. However, different authors have provided completely divergent estimates of the contribution of this binary channel to the formation of Be stars. In my opinion, the lack of evidence for any companion in many Be stars, together with the statistics of the Be fraction and star counts in clusters, make it unlikely that most Be stars form through this channel. On the other hand, I believe that present evidence suggests that a non-negligible fraction of Be stars form in this way.  Several works have shown that B-type stars in (at least massive) open clusters rotate faster than those in the field. This fast rotation is likely to be primordial and related to the initial conditions of star formation, as seems to be confirmed by the results of \\citet{str05}. It is, however, interesting to note that in the case of h \\& $\\chi$ Persei, this higher average $\\Omega$ has not {\\it yet} resulted in a high Be fraction.  There is abundant evidence demonstrating that Be stars in open clusters are more numerous around the MSTO, suggesting some sort of evolutionary effect, but claims of a preferred age with a higher fraction of Be stars are likely based on biased samples. There is also some evidence favouring an inverse dependence of the Be fraction on metallicity. However, at this point we do not really understand which mechanism(s) may cause this dependence. The fact that the fraction of Be stars varies enormously from cluster to cluster, even at a given $Z$, strongly suggests that the conditions leading to the development of the Be phenomenon must be primordial and have to be found at the cluster level."
    },
    "0601/astro-ph0601037_arXiv.txt": {
        "abstract": "We present the results of {\\it Chandra} ACIS-S snapshot observations of six radio-loud quasars (RLQs) at \\hbox{$z \\approx$ 3.5--4.7}. These observations sample luminous RLQs with moderate-to-high radio-loudness \\hbox{($R \\approx$ 200--9600)} and aim to connect the X-ray properties of radio-quiet quasars \\hbox{($R <$ 10)} and highly radio-loud blazars \\hbox{($R >$ 1000)} at high redshift. This work extends a study by Bassett et al. (2004) which used similar methods to examine \\hbox{$z >$ 4} RLQs with moderate radio-loudness \\hbox{($R \\approx$ 40--400)}. All of our targets are clearly detected. A search for extended X-ray emission associated with kpc-scale radio jets revealed only limited evidence for X-ray extension in our sample: three sources showed no evidence of X-ray extension, and the other three had 3--30\\% of their total X-ray fluxes extended $>$1\\arcsec\\ away from their X-ray cores. Additionally, we do not observe any systematic flattening of the optical-to-X-ray spectral index (\\hbox{$\\alpha_{\\rm ox}$}) compared to low-redshift quasars. These results suggest that kpc-scale X-ray jet emission is not dominated by inverse-Compton scattering of CMB-seed photons off jet electrons. We measured X-ray continuum shapes and performed individual and joint spectral fits of our data combined with eight archival RLQs. A single power-law model acceptably fit the data, with best-fit photon indices of 1.72$^{+0.11}_{-0.12}$ for moderate-$R$ sources and 1.47$^{+0.13}_{-0.12}$ for high-$R$ sources. We added an intrinsic absorption component to our model, and neither the moderate-$R$ nor the high-$R$ fits set a lower bound on $N_H$. The upper limits were \\hbox{$N_H$ $<$ 3.0 $\\times$ 10$^{22}$ cm$^{-2}$} and \\hbox{$N_H$ $<$ 2.8 $\\times$ 10$^{22}$ cm$^{-2}$}, respectively. Our spectral results suggest that intrinsic absorption does not strongly depend on radio-loudness, and high-$R$ sources have flatter power laws than moderate-$R$ sources. Overall, our high-redshift RLQs have basic X-ray properties consistent with similar RLQs in the local universe. ",
        "introduction": "Multiwavelength studies of the high-redshift ($z > 4$) universe have become increasingly widespread over the last decade. With the high sensitivities of X-ray observatories like {\\it Chandra} and {\\it XMM-Newton}, scientists have analyzed the physical conditions in the immediate vicinities of the first massive black holes to form in the Universe. Prior to 2000, there were only six published X-ray detections of quasars at $z > 4$ (e.g., Kaspi, Brandt, \\& Schneider 2000).  In the last five years, the number has risen to over one hundred largely via {\\it Chandra} and {\\it XMM-Newton} observations.  Wide-area radio surveys, such as the Parkes-MIT-NRAO survey (PMN; Griffith et al. 1995 and references therein), are vital to constrain the properties of high-redshift quasar radio cores, jets, and lobes. Radio-selected radio-loud quasars (RLQs), although sampling a minority of the total quasar population, are less prone to obscuration bias than optically selected targets, as radio emission is not affected by absorption due to dust.  The radio-loudness parameter, as defined by Kellermann et al.\\ (1989), is given by \\hbox{$R$=\\(f_{\\rm 5~GHz}/f_{\\rm 4400~\\mbox{\\scriptsize\\AA}}\\)} (rest frame), where quasars with $R\\ge10$ are RLQs and those with $R<10$ are radio-quiet quasars (RQQs).  Some of the first X-ray studies of $z>4$ quasars were of a small group of highly radio-loud blazars (\\hbox{$R\\approx$~800--27000}) in which the X-ray radiation is probably dominated by a jet-linked component (e.g., Mathur \\& Elvis 1995; Fabian et al. 1997; Moran \\& Helfand 1997; Zickgraf et al. 1997). These blazars are not suitable for representative statistical studies, however, because they represent a small minority of even the RLQ population. Recent X-ray observations with {\\it Chandra's} Advanced CCD Imaging Spectrometer (ACIS; e.g., Garmire et al. 2003) have generally focused on $z >$ 4 RQQs (\\hbox{$R \\lesssim 2$--10}) which display different \\hbox{X-ray} properties than RLQs and thus do not provide much information on RLQ \\hbox{X-ray} emission mechanisms. Consequently, understanding of high-redshift RLQ \\hbox{X-ray} emission is still limited. Prior to the current paper, there have been ten published detections of $z>4$ RLQs with $R \\approx 10$--$1000$, corresponding to typical RLQs in the local universe (excluding the blazar PMN 0525$-$3343 which has $R \\approx 800$).  In comparison to RQQs, RLQs show enhanced \\hbox{X-ray} and radio emission that can be attributed to relativistic jet-linked components. For core-dominated, flat-spectrum RLQs, the majority of the enhanced \\hbox{X-ray} emission is presumably from the small-scale (sub-parsec) jet close to the core (e.g., Wilkes \\& Elvis 1987; Worrall et al. 1987). In addition, observations with {\\it Chandra} have revealed extended \\hbox{X-ray} emission from the kpc-scale jets of RLQs (e.g., Sambruna et al. 2004; Marshall et al. 2005). Several theories have been developed to explain this unexpectedly strong extended \\hbox{X-ray} emission. Low optical luminosities often preclude pure synchrotron emission from a single electron population, leading some to propose models with more complex electron energy distributions (e.g., Stawarz et al. 2004) or models based on inverse-Compton (IC) scattering. A leading model for extended X-ray jets invokes bulk relativistic motions (\\hbox{$\\Gamma \\sim$ 1--20}) over kpc scales, thereby allowing electrons in a jet to IC scatter photons from the Cosmic Microwave Background (CMB) into the X-ray band (e.g., Tavecchio et al. 2000; Celotti, Ghisellini, \\& Chiaberge 2001). A natural consequence of the IC/CMB model is that jets should remain X-ray bright at high redshift: the \\hbox{$(1+z)^4$} dependence of CMB energy density compensates for the \\hbox{$(1+z)^{-4}$} decrease in surface brightness (e.g., Schwartz 2002). The IC/CMB model predicts that \\hbox{$z > 4$} X-ray jets should often outshine their cores (which dim with luminosity distance squared). As angular size increases with redshift beyond \\hbox{$z \\approx 1.6$} in currently favored cosmological models, the sub-arcsecond angular resolution of {\\it Chandra} is sufficient to resolve jets with projected length scales $ > $ 10 kpc at all redshifts. Evidence of an X-ray jet has been discovered in the \\hbox{$z$ = 4.30} radio-loud (\\hbox{$R$ $\\approx$ 1200}) quasar \\hbox{GB~1508$+$5714} (Siemiginowska et al.\\ 2003; Yuan et al.\\ 2003). However, the jet-to-core X-ray flux ratio is only $\\approx$ 3\\%, substantially less than the $\\approx$ 100\\% predicted by Schwartz (2002).  In a recent paper (Bassett et al. 2004, B04 hereafter), six new $z>4$ RLQ detections were reported. This study aimed to fill the gap between the many X-ray observations of $z>4$ RQQs and the extreme blazars. The B04 targets had moderate $R$ parameters of \\hbox{$\\approx$ 40--400} and were representative of the flat-spectrum RLQ population (Ivezi\\'{c} et al. 2002), with several sources near the mode of the $R$-value distribution for RLQs (see Fig.~2 of B04). B04 found enhanced X-ray emission relative to RQQs of the same redshift and optical luminosity by a factor of $\\approx$ 2. Joint spectral analysis revealed the spectra could be fitted with a power-law model with Galactic absorption and possibly an intrinsic absorption component. The evidence tentatively suggested increased absorption toward these high-redshift RLQs, generally consistent with previous RLQ studies (e.g., Cappi et al. 1997; Fiore et al. 1998; Reeves \\& Turner 2000). No extended X-ray emission associated with jets from the RLQs was detected, raising possible concern about the IC/CMB explanation of X-ray jets.  Here we report new {\\it Chandra} ACIS-S ``snapshot'' detections of six \\hbox{$z\\approx$~3.5--4.7} radio-selected RLQs to continue building the observational X-ray bridge between RQQs and extreme blazars at high redshift. We selected our targets from the large survey of Hook et al. (2002) that searched 7265 deg$^{2}$ for flat-spectrum, radio-loud quasars at high redshift based on a parent sample from the PMN survey. The PMN data cover a significant fraction of the southern sky with \\hbox{$\\delta$ $<$ 10$^{\\circ}$} to a flux-density limit of 20--72 mJy at 5 GHz. Hook et al. found a significant number of new high-redshift RLQs that are bright in the optical and radio wavebands. We chose our targets based on their bright optical and radio fluxes and their high redshifts. Four of our sources were selected to represent the majority population of flat-spectrum RLQs, with moderate-$R$ values ranging 200--1300. Two of our targets are not as representative: PMN~J1230$-$1139 and PMN~J2219$-$2719 are highly radio-loud with \\hbox{$R \\approx$ 7400--9500}. However, these targets are important as they enlarge the radio-loud ``blazar'' population with X-ray observations (only 5 blazars have been detected in X-rays previously). All the sources have published detections at both 1.4 GHz and 5 GHz and have flat radio spectra with $\\alpha_{\\rm r} \\approx -0.4$~to $+$0.3 ($f_{\\nu} \\propto \\nu^{\\alpha}$; Hook et al. 2002). We note that our targets have not been mapped at high-resolution in the radio band over multiple epochs to check for, e.g., superluminal motion. All the sources are detected by {\\it Chandra} with a minimum of $\\approx$~20 counts, as we describe in more detail below. Combined with the results of B04, we have now observed ten moderate-$R$ RLQs at high redshift with {\\it Chandra} that are representative of the overall RLQ population.  This paper brings the total sample of $z>4$ \\xray\\ detected RLQs to 21, more than tripling the number of moderately radio-loud quasars observed before B04 and filling the gap (as shown in Figure~1) between the many \\xray\\ observations of $z>4$ RQQs and the extreme blazars.\\footnote{ A regularly updated list of $z>4$ quasars with X-ray detections can be found at http://www.astro.psu.edu/users/niel/papers/highz-xray-detected.dat} In combination with the results of B04, this set of quasars is large enough to serve as a representative sample for reliable statistical evaluations of high-redshift flat-spectrum RLQ \\xray\\ properties.  Throughout this paper we adopt $H_{0}$=70 km s$^{-1}$ Mpc$^{-1}$ in a $\\Lambda$-cosmology with $\\Omega_{\\rm M}$=0.3 and $\\Omega_{\\Lambda}$=0.7 (e.g., Spergel et al.\\ 2003). Additionally, flux densities are implicitly taken to have units of erg cm$^{-2}$ s$^{-1}$ Hz$^{-1}$. ",
        "conclusions": "Our primary results are the following:  1. We have presented new {\\it Chandra} observations of six RLQs at $z \\approx$ 3.5--4.7, significantly increasing the number of X-ray detected, high-redshift RLQs. Our targets have moderate-to-high radio-loudness parameters, with $R$ $\\approx$ 200--9600.  2. We do not detect significant extended X-ray emission associated with kpc-scale X-ray jets. Nine RLQs from our sample and that of B04 showed no X-ray extension, and three RLQs displayed slight extension with only 3--30\\% of their total X-ray fluxes extending $>$1\\arcsec\\ away from their X-ray cores. This result is on the order of the modest X-ray jet discovered in the $z$ = 4.3 blazar GB 1508$+$5714  (where $\\approx$ 3\\% of the total flux was extended from the core). We conclude any spatially resolvable jets must be at 3--30 times fainter than their X-ray cores. Additionally, we do not observe any systematic flattening of optical-to-X-ray spectral index compared to low-redshift quasars. In combination with the results of B04, we have a large enough sample to reduce sensitivity to jet-orientation effects. Our results contrast with the predictions of the IC/CMB model, and we conclude that kpc-scale X-ray production is not dominated by a straightforward IC/CMB process in flat-spectrum RLQs.  3. Individual and joint spectral fitting revealed 14 RLQs can be fit acceptably using a power-law model with a Galactic absorption component. The average best-fit photon indices of \\hbox{$\\Gamma$ = 1.72$^{+0.11}_{-0.12}$} for the moderate-$R$ objects and \\hbox{$\\Gamma$=1.47$^{+0.13}_{-0.12}$} for the high-$R$ objects are consistent with similar objects at low redshift. To search for evidence of intrinsic X-ray absorption, we added a redshifted neutral-absorption component to the above model. Neither the moderate-$R$ nor the high-$R$ fits set a lower bound on $N_H$: the upper limits were \\hbox{$N_H$ $<$ 3.0 $\\times$ 10$^{22}$ cm$^{-2}$} and \\hbox{$N_H$ $<$ 2.8 $\\times$ 10$^{22}$ cm$^{-2}$}, respectively. Our results are consistent (within the 90 \\% confidence limits) with previous quasar and blazar studies. Our results indicate intrinsic absorption $N_H$ does not have a strong dependence on radio-loudness.  Longer X-ray exposures with {\\it Chandra} and {\\it XMM-Newton} of our sample are necessary to investigate further high-redshift RLQ X-ray extension, X-ray absorption, and X-ray variability. The greater photon statistics of deeper X-ray observations will enable direct detection of X-ray jets as well as provide tighter constraints on intrinsic X-ray absorption. Additionally, longer X-ray exposures will yield better timescales to explore RLQ \\hbox{X-ray} variability and its dependence on redshift. High-resolution radio imaging of high-redshift RLQs can reveal the direction and magnitude of extended radio emission, and these results can be compared to extended X-ray features to test for possible coincidence and to verify putative jets."
    },
    "0601/astro-ph0601201_arXiv.txt": {
        "abstract": "In the first paper of this series we presented EBAS, a new fully automated algorithm to analyse the lightcurves of eclipsing binaries, based on the EBOP code. Here we apply the new algorithm to the whole sample of 2580 binaries found in the OGLE LMC photometric survey and derive the orbital elements for 1931 systems. To obtain the statistical properties of the short-period binaries of the LMC we construct a well defined subsample of 938 eclipsing binaries with main-sequence B-type primaries. Correcting for observational selection effects, we derive the distributions of the fractional radii of the two components and their sum, the brightness ratios and the periods of the short-period binaries. Somewhat surprisingly, the results are consistent with a flat distribution in log P between 2 and 10 days. We also estimate the total number of binaries in the LMC with the same characteristics, and not only the eclipsing binaries, to be about 5000. This figure leads us to suggest that $(0.7\\pm 0.4)$\\% of the main-sequence B-type stars in the LMC are found in binaries with periods shorter than 10 days. This frequency is substantially smaller than the fraction of binaries found by small Galactic radial-velocity surveys of B stars. On the other hand, the binary frequency found by HST photometric searches within the late main-sequence stars of 47 Tuc is only slightly higher and still consistent with the frequency we deduced for the B stars in the LMC. ",
        "introduction": "A number of large photometric surveys have recently produced unprecedentedly large sets of high S/N stellar lightcurves \\citep[e.g.,][]{alcocketal97}. One of these projects is the OGLE study of the SMC \\citep{udalski1998} and the LMC \\citep{udalski2000}, which has already yielded \\citep{lukas2003} a few thousand lightcurves of eclipsing binaries. These two sets of lightcurves enable a statistical analysis of a large sample of short-period extragalactic binaries, for which the absolute magnitudes of all systems are known to a few percent. No such sample is known in our Galaxy.  One such statistical study was the analysis of \\citet[][]{NZ03}, who derived the orbital elements and stellar parameters of 153 eclipsing binaries discovered by OGLE in the SMC.  North \\& Zahn examined the eccentricities of the eclipsing binaries as a function of the ratio between the stellar radii and the binary separation, and compared their result with a similar dependence derived from the elements of the binaries discovered by \\citet{alcocketal97} in the LMC. They then considered the implication of their findings for the theory of tidal circularization of early-type primaries in close binaries \\citep[e.g.,][]{zahn75, zahn77, ZB89}.  In a follow-up paper \\citet{NZ04} analyzed another set of 510 lightcurves selected from the 2580 eclipsing binaries discovered in the LMC by the OGLE team \\citep[][]{lukas2003}.  Again, they derived the ratio between the stellar radii and the binary separation for these systems, and examined the dependence of the binary eccentricity on this ratio.  The previous analyses used only a small fraction of the sample of eclipsing binaries found in the OGLE data. The goal of the present study is to analyze the whole sample of 2580 lightcurves discovered by OGLE in the LMC, and to derive some statistical properties of the short-period binaries, after correcting for observational selection effects. Such a correction is essential for the derivation of the period distribution, for example, because the probability of detecting an eclipsing binary is a strong function of the orbital period. In order to apply an appropriate correction one needs a complete homogeneous data set of lightcurves, all discovered by the same photometric survey of a well defined sample of stars. These requirements were exactly met by the OGLE data set for the LMC, enabling such analysis for a large sample of eclipsing binaries for the first time.  A completely automated algorithm is needed to analyse the large set of eclipsing binaries at hand. Two such codes were developed recently. \\citet{WW1} have constructed an automatic scheme based on the Wilson-Divenney (=WD) code to analyze the OGLE lightcurves detected in the SMC in order to find eclipsing binaries suitable for distance measurements.  At the last stages of writing this paper another study with an automated lightcurve fitter --- DEBiL, was published \\citep{devor2005}.  DEBiL was constructed to be quick and simple, and therefore has its own lightcurve generator, which does not account for stellar deformation and reflection effects. This makes it especially suitable for detached binaries. We will use in this work our EBAS, which was presented in Paper I (Tamuz, Mazeh \\& North 2005) and is based on the EBOP code.  The complexity of EBAS is in between DEBiL and the automated WD code of Wyithe \\& Wilson.  To facilitate the search for global minima in the convolved parameter space, EBAS performs two parameter transformations. Instead of the radii of the two stellar components of the binary system, measured in terms of the binary separation, EBAS uses the total radius, which is the sum of the two relative radii, and their ratio. Instead of the inclination we use the impact parameter --- the projected distance between the centres of the two stars in the middle of the primary eclipse, measured in terms of the total radius.  The set of parameters of the EBAS version used here includes the bolometric reflection of the two stars, $A_p$ and $A_s$. When $A_p=1$ the primary star reflects all the light cast on it by the secondary. Together with the tidal distortion of the two components, which is mainly determined by the mass ratio of the two stars, the reflection coefficients $A_p$ and $A_s$ determine the light variability of the system outside the eclipses. Paper I discussed the reliability of the values of these two parameters as found by EBAS.  To simplify the analysis, the present version of EBAS assumes there is no contribution of light from a third star and that the mass ratio is unity.  Paper I discusses the implication of these choices on the parameters' values, showing that the values of the total radius and the surface brightness ratio are only slightly modified by these two assumptions. The only parameter which is systematically modified by the third-light assumption is the impact parameter, which is directly associated with the orbital inclination of the binary.  Note, however, that we are not interested in one specific system but aim, instead, at deriving the gross characteristics of the short-period binaries.  As the analysis of this paper does not use the inclinations of the eclipsing binaries for the derivation of the statistical features of the short-period binaries, we regard the resulting distributions as probably correct.  Paper I introduces a new 'alarm' statistic, $\\mathcal{A}$, to replace human inspection of the residuals.  EBAS uses the new statistic, which is sensitive to the correlation between neighbouring residuals to decide automatically whether a solution is satisfactory. In this work we consider only the 1931 systems that yielded solutions with low enough alarm value.  To check the reliability of our results we compare our geometrical elements with those derived by \\citet[][hereafter MiP05]{michalska05} with the WD code for 85 binaries, based on EROS, MACHO and OGLE data. The comparison is reassuring, as it shows that the geometrical parameters of EBAS are close to the ones derived by the more sophisticated WD approach.  To derive the orbital distributions of the short-period binaries we had to trim the sample, to get a homogeneous sample which we could correct for selection effects. We were left with 938 main-sequence binaries with periods shorter than 10 days and system magnitudes between 17 and 19 in the $I$ band, most of which are binaries with B-type primaries. This makes the range of the derived period distribution quite narrow.  Nevertheless, the data yielded somewhat surprising distributions, which might have implications on binary population studies.  Section~\\ref{elements} presents the resulting orbital elements of the OGLE LMC eclipsing binaries and compares the derived elements with those of MiP05. Section~\\ref{trimming} details the procedure to focus on a well defined homogeneous subsample and Section~\\ref{statistics} derives the statistical features of the short-period binaries, after correcting for the observational selection effects. Section~\\ref{discussion} discusses the new findings, and Section~\\ref{summary} summarizes the paper. ",
        "conclusions": "% \\label{discussion}  \\subsection{The period distribution}  Our analysis suggests that the binaries with B-type primaries in the LMC have flat log-period distribution between 2 and 10 days. This distribution, which was obtain after correcting for observational selection effects, is very different from the distributions of the whole samples of the SMC and the LMC eclipsing binaries \\citep{udalski1998, lukas2003}.  Evidently, correcting for the selection effects reveals a substantially different period distribution.  The period distribution derived here is probably not consistent with the period distribution of \\citet{DM91}, who adopted a Gaussian with $\\overline{\\log P}=4.8$ and $\\sigma=2.3$, $P$ being measured in days. Their distribution would provide within a log P interval almost twice as many systems at $\\log P = 1.0$ than at $\\log P = 0.3$ (more precisely, the factor would be 1.7), while the new derived distribution is probably flat.  It is interesting to note that \\citet{heacox1998} reanalysed the data of \\citet{DM91} and claimed that $f(a)$, the distribution of the G-dwarf orbital semi-major axis, $a$, is $f(a) \\propto a^{-1}da$, which implies a flat log orbital separation distribution. This is equivalent to the present probable result, although the latter refers to LMC binaries with B-type primaries, and is limited only to a very small range of orbital separation. On the other hand, \\citet[][hereafter HaMUA03]{halbwachs2003} analysed the spectroscopic binaries found within the G- and K-stars in the solar neighbourhood and the ones found in the Pleiades and Praesepe, and found that the log-period distribution is ``clearly rising until about 10 days''.  However, HaMUA03's sample included only a few systems with periods longer than 2 and shorter than 10 days. Actually, there were 7 such systems in the ``extended sample'' of nearby stars and 5 additional binaries in the Pleiades and Praesepe. In fact, if we divide these 12 binaries into two bins with equal log width we find 6 in the short-period bin and 6 in the long-period one. It is true that the correction factor of the spectroscopic binaries of these two bins is somewhat different, and therefore the distribution of this sample might still be consistent with the conclusion of HaMUA03. However, the small number of binaries with periods shorter than 10 days in the G- and K-binaries is also consistent with flat log distribution below 10 days. We verified that by building the cumulative distribution of the $\\log P$ values of the 12 G- and K-binaries, and comparing it with the theoretical flat log distribution. The one-sided Kolmogorov-Smirnov test gives a probability of 0.7 that the observed distribution is consistent with the flat log distribution.  The flat log-period distribution implies a flat log orbital-separation distribution for a given total binary mass. Such a distribution indicates that there is no preferred length scale for the formation of short-period binaries \\citep{heacox1998}, at least in the range between 0.05 and 0.16 AU, for a total binary mass of 5 $M_{\\odot}$. Alternatively, the results indicate that the specific angular momentum distribution is flat on a log scale at the $10^{19}$ $cm^2s^{-1}$ range.  A note of caution is in order here, as the present analysis refers to {\\it all} eclipsing binaries in a given absolute magnitude range. Even if the overwhelming majority of the binaries lie on or close to the main sequence, a small unknown number of them may have gone through a mass exchange phase, which would have affected both the orbital period and the stellar radii. This is very important for the interpretation of the new result, since the most interesting quantity is probably the {\\it primordial} period distribution, which was not affected by the subsequent binary evolution.  In fact, \\citet{harries2003} and \\citet[][hereafter HiHoHa05]{hilditch2005} studied in details 50 eclipsing binaries in the SMC and found more than half of them to be in a semi-detached post-mass-transfer state. However, the binaries studied in the SMC, with limiting brightness of $B < 16$, are brighter than the systems in the trimmed sample, and are composed mainly of O to B1-type primaries, while most of the primaries in the present sample are cooler B-type stars. Apparently, the fractional radii of the stars studied by HiHoHa05 is substantially larger than the typical fractional radii of the present sample. We therefore suggest that while a large fraction of the HiHoHa05 sample has gone through mass exchange, most of the binaries studied here are still on the main sequence.  We therefore suggest that the short-period B-type binaries in the LMC might have had a primordial flat log-period distribution between 2 and 10 days, and a similar feature {\\it could} have been found in binaries with G-type primaries in the solar neighbourhood. Although this probable result refers to a very narrow period range, a binary formation model should account for this scale-free distribution, which might be common in short-period binaries.  \\subsection{The binary fraction}  It would be interesting to compare the frequency of B-star binaries found here for the LMC with a similar frequency study of B stars in our Galaxy. However, such a large systematic study of eclipsing binaries is not available. Instead, a few radial-velocity and photometric searches for binaries in relatively small Galactic samples were performed. In what follows we compare the frequency derived here with results of radial-velocity searches for binaries, in samples of Galactic B stars as well as within the nearby K and G stars. In addition, we discuss the binary frequency found by the HST photometric monitoring of 47 Tuc.  \\subsubsection{Radial-velocity searches for binaries}  \\citet{wolff78} studied the frequency of binaries in sharp-lined B7--B9 stars.  Her Table~1 presents 73 such stars with $V\\sin i < 100$~km\\,s$^{-1}$, among which 17 are members of binary systems with $P_{\\mathrm orb} < 10$~days. This represents a percentage of $23$\\%, much larger than the frequency we find in the LMC.  On the one hand, the frequency of \\citet{wolff78} may be considered a lower limit, since her sample was biased toward sharp-lined stars, because her purpose was to determine what fraction of the intrinsically slowly rotating B stars are members of close binary systems. This introduced a bias toward low $V\\sin i$ values and therefore small orbital inclinations, hence lower detection probabilities, assuming the spin and orbital axes are aligned. On the other hand, the sample is magnitude limited, but the correction for the Branch bias \\citep{branch76} was not applied. The most extreme correction factor for the bias towards the more luminous binary systems, which holds for systems with two identical components, is $f=0.35$. Applying this factor to the rate of late B-type binaries with $P_{\\mathrm orb} < 10$~days leaves a minimum frequency of $8\\pm 2$\\%.  This remains one order of magnitude larger than our frequency in the LMC.  Similar and even higher values of binary frequency were derived for B stars in Galactic clusters and associations.  \\citet{ML91} have determined the rate of binaries among B stars (a few O and A-type stars are also included in the sample) in the Orion OB1 association, and obtained an overall frequency of binaries with $P_{\\mathrm orb} < 10$~days of $26\\pm 6$\\%. There is a slight trend toward a higher rate of binaries among the early B stars than among the late ones. For B4 and earlier stars they found 60\\% binaries, while only $\\sim 27$\\% of the B5 and later stars were found to be RV variables, which is reminiscent of what we find in the LMC. However, this is hardly significant as the latter frequency is based only on 3 variables.  \\citet{GM01} found a rate of binaries as high as $82$\\% among stars hotter than B1.5 in the very young open cluster NGC~6231, though in a sample of 34 stars only. Most orbits have periods shorter than 10 days.  \\citet{R96} estimated a rate of binaries of at least $52$\\% for 36 B1-B9 stars in the same cluster, but his data did not allow him to determine the orbital periods.  These studies show that binaries hosting B-type stars tend to be more frequent in young Galactic clusters than in the field. However, even though our LMC binaries are representative of the LMC field rather than of LMC clusters, the frequency we derived appears much lower than that of the binaries in the Galactic field.  One survey of a complete, volume limited, sample of binaries in the solar neighbourhood is that of HaMUA03, who carefully studied a complete sample of K- and G-type stars, and not B stars.  Out of a bias free sample of 405 in the solar neighbourhood, they found 52 binaries, 11 of them with periods shorter than 10 days. This results in a fractional frequency of $(2.7\\pm0.8)$\\% binaries, if one takes just the Poisson distribution to estimate the error.  We therefore suggest that the fractional frequency of K- and G-type stars in the solar neighbourhood found in binaries with periods shorter than 10 days is higher by a factor of about $3.9\\pm1.9$ than the corresponding frequency of B-type stars in the LMC. The error estimate of the factor between the two fractional frequencies was derived by assuming normal distribution, and is only a coarse estimate, because of the small number statistics of the K- and G-type binaries. In fact, the difference between the two fractional frequencies is at about the $2\\sigma$ level.  \\subsubsection{HST Photometric monitoring of 47 Tuc}  A few photometric studies were devoted to the search of eclipsing binaries in Galactic globular clusters \\citep[e.g.,][]{yan96, kaluzny97}. One of the recent photometric studies is the HST 8.3-day observations of 47 Tuc \\citep{albrow01}, which monitored 46,422 stars in the central part of the cluster and discovered 5 eclipsing binaries with periods longer than about 4 days. Assuming a flat log-period distribution of binaries, \\citet{albrow01} concluded that the primordial binary frequecy up to 50 years was $(13\\pm6)$\\%. This estimate translates into $(2\\pm1)$\\% for periods between 2 and 10 days. This value, derived for late-type stars, is much smaller than the frequency derived by the radial-velocity surveys of B stars, and is close to the $(0.7\\pm0.4)$\\% frequency we derived for the LMC.  \\subsubsection{Comparison between the two approaches}  We conclude that the frequency of binaries we find in the LMC is substantially smaller than the frequency found by Galactic radial-velocity surveys. The detected frequency of B-type binaries is larger by a factor of 10 or more than the frequency we find, while the binary frequency of K- and G-type stars is probably larger by a factor of four. On the other hand, the binary frequency found by photometric searches in 47 Tuc is only slightly higher and still consistent with the frequency we deduced for the LMC. It seems that the frequency derived from photometric searches is consistently smaller than the one found by radial-velocity observations.  We are not aware of any observational effect that could cause such a large difference between the radial velocity and the photometric studies.  The sensitivity of the two approaches depends on the ability to discover binaries with secondaries of small radii, in the photometry case, and small masses, in the case of the radial-velocity searches.  In order to account for the large differences between the photometric and the spectroscopic searches we have to assume that the photometric searches detect only ten percent of the binaries and missed all the others, that could have been discovered by radial-velocity techniques. Although \\citet{wolff78} suggested that many short-period binaries in her radial-velocity sample have low-mass secondaries, this effect can not explain a factor of ten in the detected frequency. The discussion of the radii distribution of our sample also indicates that if such an effect is present in the LMC binaries it must be small. After all, the radial-velocity searches are not sensitive to small-mass secondaries, and they too suffer from selection effects which depends on their precision. Therefore, the large difference in the binary frequencies is probably real and remains a mystery. Obviously, it would be extremely useful and interesting to have studies of eclipsing binaries similar to the present one in our Galaxy, as well as in other nearby galaxies.  \\subsection{The stellar radii and the surface brightness ratios}  The distributions of the two fractional radii depend strongly on the derivation of the ratio of radii, which is not well determined from the lightcurves. Still, a statistical analysis of the distributions is meaningful, given the high correlation between the true and derived $k$ values, as shown in the simulations in Paper I.  The distributions of the radii show sharp peaks at the same value of 0.1, as can be seen in Fig~\\ref{fig:param_distr_corrected}, and in Fig~\\ref{fig:param_distr_corrected_long} in particular. This suggests, but admittedly not prove, that our sample is mostly composed of stars with similar radii. The distribution of $J_s$ tends to reinforce this interpretation, since it suggests a mass ratio --- hence a ratio of radii --- close to one.  For the primary, the rise between $0.05$ and $0.1$ is due essentially to the limit imposed on the orbital period and to the minimum radius of a star on the ZAMS. On the other hand, the sharp rise of the secondary radius distribution up to a radius of about 0.1 is probably real. This feature of the secondary radius distribution could have been the result of a selection effect --- too small a secondary results in an eclipse too shallow to be detected by the OGLE photometry. We suggest that this is not the case. The detectability limit of the sample is such that any binary with eclipse deeper than about 0.1 mag, corresponding to 10\\% drop in the system brightness, is included in the sample. Such a drop can be caused by a secondary with a radius which is about 0.3 of the primary radius, for inclination of $90^{\\circ}$ and $J_s < 1$. If the primary radius peaks at $\\sim 0.1$, the sample is then complete for secondary radius down to about 0.03.  The peaks at about 0.1 in Fig~\\ref{fig:param_distr_corrected_long} indicate that a substantial fraction of the primary and the secondary stars in the sample have radius of 2.5--3 $R_{\\odot}$, assuming a typical period of 7 days in Fig~\\ref{fig:param_distr_corrected_long} and a typical mass of 5 $M_{\\odot}$. The sharpness of the peaks is probably caused by the magnitude limits of the trimmed sample.  \\subsection{The eccentricity distribution}  Fig~\\ref{fig:ecosw}, with the eccentricity as a function of the sum of the two radii, displays upper and lower envelopes that can be approximated with two straight lines, with a slope of about 1.7 in absolute value. Although both the shape and slope of these envelopes may be plagued by selection biases, it is interesting to remark that the populated area corresponds to the condition \\begin{equation} $$r_t=r_p+r_s\\lesssim 0.59\\,(0.85-|e\\cos\\omega|) \\ \\ . \\end{equation} This roughly generalizes the condition $r_{average}\\lesssim 0.25$ to the case of eccentric systems, which cannot keep an eccentricity larger than some limiting value. This limiting value of the eccentricity depends on the stellar radii, whose sum can not be larger than a given critical fraction of the distance {\\it at periastron}. Because of the $\\cos\\omega$ factor, it is reasonable to assume that the real relation is rather $$r_p+r_s\\lesssim 0.5\\,(1-e) \\ .$$ Systems with eccentricities larger than that will undergo partial tidal circularization until they satisfy this inequality; after that, tidal effects become almost inefficient. The boundary is clear-cut even though the systems do not share the same age, due to the very steep dependence of circularization time on relative radius: $t_{\\rm circ}\\propto (R/a)^{-21/2}$ \\citep[][]{zahn75, zahn77}.  Note also that there are quite a few binaries with small radii for which the parameter $e \\cos \\omega$ is very close to zero. It is difficult to assume that all these system have $\\cos \\omega=0$. Probably those systems have circular orbits. It would be interesting to find out whether these systems were circularized during their main-sequence phase, despite their small radii, or maybe their very small eccentricity is primordial."
    },
    "0601/astro-ph0601421_arXiv.txt": {
        "abstract": "We report on {\\it Spitzer} and Gemini observations of the jet of Centaurus A in the infrared, which we combine with radio, ultraviolet and X-ray data. {\\it Spitzer} detects jet emission from about 2 arcmin from the nucleus, becoming particularly bright after the jet flare point at $\\sim 3.4$ arcmin. Where X-ray and infrared emission are seen together the broad-band data strongly support a synchrotron origin for the X-rays. The jet flare point is marked by a broad, diffuse region of X-rays which may be associated with a shock: we discuss possible physical mechanisms for this. The infrared jet persists after the flare point region although X-ray emission is absent; it is plausible that here we are seeing the effects of particle acceleration followed by downstream advection with synchrotron losses. Gemini data probe the inner regions of the jet, putting limits on the mid-infrared flux of jet knots. ",
        "introduction": "Centaurus A is the closest radio galaxy (we adopt $D = 3.4$ Mpc, Israel 1998) as well as the closest active galaxy and large elliptical. Its proximity makes it a vital laboratory for AGN studies of all kinds. Its nucleus shows both heavily obscured and unobscured X-ray components (Evans \\etal\\ 2004), making it one of the few low-power radio galaxies to show X-ray evidence for the obscuring torus of canonical unification models (Evans \\etal\\ 2006). Its X-ray jet, one of the first to be discovered (Schreier \\etal\\ 1979) has been studied in great detail with {\\it Chandra} (Kraft \\etal\\ 2000, 2002; Hardcastle \\etal\\ 2003, hereafter H03; Kataoka \\etal\\ 2006). Because of the high spatial resolution (1 arcsec = 17 pc) compared to that available for more typical X-ray jets in more powerful FRI sources, observations of the jet provide evidence for both localized and diffuse particle acceleration processes. Finally, Cen A's SW radio lobe, in expanding through the ISM of the host galaxy, drives what is currently the only clear example of a high Mach number, attached bow shock to be observed in X-rays around a radio galaxy (Kraft \\etal\\ 2003).  Cen A's main disadvantage as a subject for broad-band studies is the strong dust lane, which obscures the inner regions of the source in the optical to ultraviolet. Partly as a result of this, and partly because the jet is relatively weak compared to the emission from stars, there have until recently been no clear detections of synchrotron emission from the jet at frequencies between radio and X-ray, although several claims of optical and infrared emission that may be related to the jet or material around it have been made (e.g.\\ Brodie \\etal\\ 1983, Joy \\etal\\ 1991, Leeuw \\etal\\ 2002). This contrasts with the situation in other well-studied low-power jets, such as M87 (e.g. Perlman \\etal\\ 2001) or 3C\\,66B (Hardcastle \\etal\\ 2001) in which data in the infrared, optical and ultraviolet support a synchrotron origin and provide constraints on particle acceleration and energy-production processes.  However, recent {\\it Spitzer} observations have detected the jet of Cen A in the mid-infrared (Brookes \\etal\\ 2006), while {\\it GALEX} detects it in the ultraviolet (Neff \\etal\\ 2003). In this letter we combine these data with new and archival radio and X-ray observations and discuss their implications for particle acceleration processes. In addition, we use observations of the nuclear regions with Gemini at 10 $\\mu$m to place constraints on the properties of the nucleus and the inner jet.  Except where otherwise stated, spectral indices $\\alpha$ are the energy indices, defined in the sense that flux $\\propto \\nu^{-\\alpha}$. ",
        "conclusions": "The {\\it Spitzer} observations confirm that the broad-band spectrum of the large-scale Cen A jet can be described with a synchrotron model, as in other FRI sources, although the detailed spectral shape almost certainly requires a multi-component model for the synchrotron emission. However, the infrared detection points up the importance of the region we have called the flare point. Both the radio and infrared brighten by a factor $\\sim 3$ here (Fig.\\ \\ref{slice}) while there is little or no X-ray emission after the extended component of region G. The short synchrotron lifetime of X-ray-emitting electrons means that X-ray emission must always be associated with high-energy particle acceleration: but could the lack of X-rays after region G imply that this region, which is roughly coincident with the entry of the jet into the lobe, represents the `last gasp' of particle acceleration in Cen A's jet? In a field strength of 3 nT, the equipartition value for this part of the jet, the electron energy loss timescale ($E/({\\rm d}E/{\\rm d}t)$) is $2 \\times 10^4$ years for electrons emitting at 4.5 $\\mu$m, and $4 \\times 10^4$ years for electrons emitting at 24 $\\mu$m. The projected distance from the end of the X-ray emission at the flare point to the end of the jet, where only radio and 24-$\\mu$m emission is detected, is roughly 2.5 kpc, implying a light travel time of $\\sim 8 \\times 10^3$ years. Thus, for particle acceleration to be absent in this region, we require jet speeds of $(0.2/\\sin\\theta)c$, where $\\theta$ is the angle to the line of sight. As speeds in the inner jet are $\\ga 0.5c$, and the angle to the line of sight may be relatively large (see discussion in H03) this is not impossible, so that it could indeed be the case that significant high-energy particle acceleration ceases at the flare point.  This motivates us to ask a further question: what is the nature of the extended X-ray emission at the flare point? Most of the X-rays (Fig.\\ \\ref{slice}) come from a region only 10 arcsec, or 170 pc, in size. This is still larger than the expected travel distance for 1-keV-emitting electrons (at most 100 pc) and in fact extended X-ray emission is seen on scales up to around 30 arcsec, making it difficult to sustain a model in which the particle acceleration here takes place at a single point if the magnetic field strength has its equipartition value, although only modest decreases in the magnetic field strength, by a factor of a few, would be necessary to make a one-shot acceleration model viable, since the loss timescale goes as $B^{-3/2}$. More puzzling in this picture is the offset between the peak X-ray, radio and infrared surface brightnesses seen in Fig.\\ \\ref{slice}. It is also not clear what the physical relationship is between the diffuse X-ray emission at the flare point and the X-ray knots G1--3. In fact there is a striking similarity in the X-ray, albeit on larger scales, between the flare point and other regions in the inner part of the jet where we see compact X-ray features followed by diffuse X-ray emission, such as the region around knot BX2 (H03).  In H03 we argued that the knots in the inner jet were related to shocks as a result of interactions between the jet fluid and small-scale obstacles in the jet. The flare point is different in that it appears to affect the whole jet. It would be tempting, since the extended emission is associated with the entry of the jet into the high-surface-brightness regions of the NE lobe, to suggest that we are seeing an extended reconfinement shock, in which case the similarity of the length of the X-ray emitting region and the width of the radio jet implies a Mach number $\\sim 2$. The roughly tapering shape of the downstream X-ray emission is consistent with this idea (Sanders 1983) if only the inward-propagating shock accelerates particles to high energies. The idea that the jet decelerates rapidly and significantly at the flare point is consistent with the observation that it becomes both broader and brighter at this position. The details of the shock structure would then depend on the velocity structure of the jet, and it may be that a detailed model taking this into account could reproduce the offsets between the emission peaks at different frequencies. An alternative model is that the jet interacts at the flare point with some large-scale external feature that is able to affect the whole jet. Gopal-Krishna and Saripalli (1984) have pointed out the coincidence between the radio and X-ray flare point and an optical `shell', seen in deep images, which they consider to be a remnant of a cannibalized galaxy. A series of shocks caused by such an interaction would be equally consistent with what we observe, if they can be distributed over the entire region of X-ray emission or, again, if the magnetic field is sub-equipartition.  Qualitatively, the bulk deceleration at the flare point in either of these scenarios is also consistent with the observed onset of bending of the jet there. However, if the jet really has a high bulk speed, as required to avoid {\\it in situ} particle acceleration after the flare point, its density must be low. For example, if we assume that the jet is bent by the ram pressure of hot external material moving at the sound speed, then, applying Euler's equation (e.g. Eilek \\etal\\ 1984) and using the parameters of Kraft \\etal\\ (2003), the jet density must be $10^{-4}$ times the external density if the speed is $0.2c$. This is still a factor 5 above the minimum possible effective jet density (from the minimum-energy condition, assuming that the jet is a pure lepton/magnetic field plasma), but if any entrainment of baryonic material takes place in the inner jet, or if there is a significant departure from equipartion, the jet will be heavier, in which case relativistic bulk speeds at the bend would be unrealistic. In that situation, the post-shock jet speed would have to be slower, and some continuing {\\it in situ} particle acceleration would be required to explain the extended infrared jet."
    },
    "0601/astro-ph0601617_arXiv.txt": {
        "abstract": "In many models of inflation, the period of accelerated expansion ends with preheating, a highly non-thermal phase of evolution during which the inflaton pumps energy into a specific set of momentum modes of field(s) to which it is coupled.  This necessarily induces large, transient density inhomogeneities which can source a significant spectrum of gravitational waves.  In this paper, we consider the generic properties of gravitational waves produced during preheating, perform detailed calculations of the spectrum for several specific  inflationary models, and identify problems that require further study.   In particular,  we argue that if these gravitational waves  exist  they will necessarily fall within the frequency range that is feasible for direct detection experiments -- from laboratory through to solar system scales. We extract the gravitational wave spectrum from numerical  simulations of preheating after $\\lambda \\phi^4$ and $m_{\\phi}^2 \\phi^2$ inflation, and find that they  lead to a gravitational wave amplitude of around $\\Omega_{gw}h^2\\sim 10^{-10}$. This is considerably higher than the amplitude of the primordial gravitational waves produced during inflation. However, the typical wavelength of these gravitational waves is considerably shorter than LIGO scales, although in extreme cases they may be visible at scales accessible to the proposed BBO mission. We survey possible experimental approaches to detecting any gravitational wave background generated during preheating. ",
        "introduction": "In recent years, the Cosmic Microwave Background (CMB) has been our primary window into the primordial universe. Scale invariant density perturbations, sourced by quantum fluctuations during inflation and processed by the photon-baryon plasma, are by far the most satisfying explanation for the observed temperature anisotropies in the CMB. Inflation also predicts the existence of a scale invariant spectrum of primordial gravitational waves, sourced by the same quantum fluctuations  that underlie the scalar perturbations. Gravitational waves are only weakly coupled to matter fields, and move freely through the universe from the moment they are produced. They are thus the deepest probe of the early universe of which we currently know \\cite{Starobinsky:1979ty,Rubakov:1982df,Fabbri:1983us,Abbott:1984fp,Allen:1987bk,Turner:1993vb,Boyle:2005se}.  In the short term, the best hope for observing primordial gravitational waves is via their contribution to the B-mode of the CMB polarization. However, foreground weak lensing puts a fundamental limit on a B-mode signal sourced by gravitational waves \\cite{Knox:2002pe,Hirata:2003ka,Seljak:2003pn}. Consequently,  attention has recently focussed upon direct detection experiments \\cite{Boyle:2005se,Ungarelli:2005qb,Crowder:2005nr,Smith:2005mm} which, while enormously challenging, may ultimately be more sensitive to a stochastic  gravitational wave background than the CMB B-mode.  These experiments are sensitive to frequencies far higher than those that contribute to the CMB, since the physical sizes of detectors  necessarily range from laboratory scales through to solar system  scales.  However, as the inflationary spectrum is almost scale-free, the amplitude at short scales is not  dramatically different from that seen at CMB scales, at least for simple models of inflation.  The primordial spectrum is not alone: gravitational waves are produced whenever there are large, time-dependent inhomogeneities in the matter distribution. These more pedestrian gravitational waves are generated ``classically'', much like radiation produced by the oscillation of an electron, whereas the inflationary  tensor perturbations are sourced by quantum fluctuations in the spacetime background.  During the evolution of the universe, there are several hypothetical mechanisms which would generate large, local inhomogeneities. These include first order phase transitions \\cite{Kamionkowski:1993fg}, brane-world models \\cite{Sahni:2001qp,Sami:2004xk}, networks of cosmic strings \\cite{Vilenkin:1981bx}, or inhomogenous neutrino diffusion \\cite{Dolgov:2001nv}. Here, we investigate a further possibility:  gravitational waves produced during preheating \\cite{Felder:2000hq,Garcia-Bellido:1997wm,Greene:1997ge,Greene:1997fu,Traschen:1990sw,Kofman:1994rk,Greene:1998nh,Giudice:1999fb,Khlebnikov:1997di,Garcia-Bellido:1998gm,Kofman:1997yn,Bassett:1998wg,Easther:1999ws,Parry:1998pn,Finelli:1998bu,Liddle:1999hq,Bassett:2005xm,Suyama:2004mz,Tsujikawa:2002nf,Finelli:2001db,Tilley:2000jh,Henriques:2003ga}, a period of non-thermal evolution following the end of inflation.  While the detailed dynamics of resonance and preheating is a complicated, nonlinear problem, the basic picture is very simple.  For a given combination of fields, parametric resonance occurs when some subset of Fourier modes have exponentially growing solutions, driven by the oscillating inflaton  field. The resonant modes are  quickly pumped up to a large amplitude $\\Phi^2 \\approx \\langle \\chi^2 \\rangle$. Resonance ends when nonlinearities render further growth kinematically expensive, leaving the  universe  far from thermal equilibrium. Thermalization occurs as the excited modes dissipate their energy over the entire spectrum via self-interaction. The rapid rise in the mode amplitude can be associated with an exponentially growing occupation number (at least for bosonic species). If one transforms into position-space, the highly pumped modes correspond to large, time dependent inhomogeneities, ensuring the matter distribution has a non-trivial quadrupole moment, sourcing the production of gravitational radiation.  This topic was first addressed by Khlebnikov and Tkachev \\cite{Khlebnikov:1997di},  and has not been widely discussed within the experimental community. We believe that the time is ripe for revisiting this question.   Since \\cite{Khlebnikov:1997di} was written, technological approaches to gravitational wave detection have advanced considerably. Moreover, there have been significant advances in the theoretical understanding of preheating in more complicated inflationary models. Finally,  numerical simulations benefit from the gains in computational power over this interval.    It will turn out that for ``typical''  cosmological parameters, gravitational radiation sourced by preheating has a peak frequency in the MHz band. Coincidentally, a new generation of detectors has been proposed which is tuned to gravitational waves in this range \\cite{Ballantini:2005am}, although their strain sensitivity would need improve over current values by around $10^6$ in order to detect the spectra we compute here. This is clearly ambitious, but probably not unreasonable when compared to the 25 year interval anticipated before BBO [Big Bang Observer] will be in a  position to detect the inflationary  gravitational wave spectrum \\cite{Phinney}.  Since preheating occurs in many (although not all) inflationary models, any gravitational wave signal associated with preheating would provide a new and currently unexplored window into inflationary physics. It is worth emphasizing that the physical principles that underlie the calculations in this paper are well established and do not rely on exotic physics -- in particular the mechanism that governs the generation of the gravitational waves is simply the usual quadrupole related emission.\\footnote{The possibility that preheating dynamics are directly affected by ``back-reaction'' from metric perturbations has also been discussed \\cite{Finelli:1998bu,Parry:1998pn,Easther:1999ws,Bassett:1998wg,Finelli:2001db,Suyama:2004mz}  but we do not address this effect in this work.} Finally, this signal carries information about the epoch at the end of inflation, opening a new window into the early universe. With the stakes this high any possible observational opportunity must be carefully explored.  Unlike the primordial spectrum, gravitational waves induced by preheating (or any subsequent phase transition) will not be scale-free; indeed their spectrum is indicative of the complicated processes that generate them. This is both a boon and a possible pitfall. A scale dependent spectrum necessarily contains more information than  a scale-free spectrum, so the detection and subsequent mapping of a gravitational wave background induced by preheating would yield a  rich trove of information.  However,  gravitational wave detectors are subject to physical limitations, since they all ultimately measure  deformations in (an array of) physical objects induced by passing gravitational waves.  Thus, even in principle, it is hard to imagine the direct detection of gravitational waves below atomic scales, or beyond solar system scales. To be sure, this is a large range but it is very much shorter than any relevant cosmological scale. Fortunately, by a happy numerical coincidence, any gravitational wave spectrum generated during preheating will peak at scale between a few meters and millions of kilometers -- which overlaps with the ``golden window'' open to direct detection experiments.  In Section (\\ref{sect:preheating}), we briefly review preheating and discuss previous work. Following that, in Section (\\ref{sect:lengthscales}) we demonstrate that preheating gravitational waves from inflation scales ranging from TeV to the GUT scale will peak around $1$ Hz to $10^8$ Hz. In Section (\\ref{sect:production}) we discuss our methodology for numerically computing the gravitational wave spectrum. In Section (\\ref{sect:results}) we present our results for $\\lambda \\phi^4$ and $m^2 \\phi^2$ inflation, confirming and extending the results of \\cite{Khlebnikov:1997di}. We note that while the inflationary dynamics of the two models are very similar, their resonant behaviors diverge considerably.   Finally, we discuss the theoretical and observational implications of our results and lay out future plans to further improve the computational methodology in Section (\\ref{sect:conclusions}). ",
        "conclusions": "\\label{sect:conclusions}    \\begin{figure*} \\myfigure{7in}{sensitivityplot.eps}% \\caption{Projected sensitivities of gravitational wave detectors (dotted lines) to a stochastic gravitational wave spectrum (full lines) in $\\log(\\Omega_{gw} h^2)$ and $\\log(f/\\mathrm{Hz})$. The limits are for \\emph{correlated} observatories, except for LISA. The primordial inflation spectrum is plotted for a scale-invariant GUT scale inflation. The gravitational wave spectra generated by the quartic and quadratic type inflation models for a GUT scale inflation are schematically shown here; see Section (\\ref{sect:results}) for more detailed plots. In addition, we have extrapolated a $f^3$ tail to the red of the simulated results, which is a conservative estimate on the degree of the Hubble scale cut-off to the gravitational wave spectrum \\cite{Liddle:1999hq}. We have also added the expected peak point of a TEV scale hybrid inflation model with high oscillation scale, calculated using the naive formula (\\ref{eqn:simpleamplitude}) \\cite{Garcia-Bellido:1998gm}.\\label{fig:sensitivityplot}} \\end{figure*}   We have carefully investigated the production of gravitational waves during preheating,   reproducing the work of Khlebnikov and Tkachev \\cite{Khlebnikov:1997di} for the $\\lambda \\phi^4$ inflation model and extending it to the $m_{\\phi}^2\\phi^2$ case.  For both models we show numerically that  preheating is  a sizable source of gravitational waves with frequencies of around $10^{6}\\sim 10^{8}$ Hz, and peak power of $\\Omega_{gw}h^2\\approx 10^{-9}  \\sim 10^{-11}$.  We present simple scaling arguments to predict the overall properties of the spectrum for a broader class of inflationary models. We see that the spectrum of gravitational waves induced by preheating peaks at a scale proportional to $1/\\sqrt{H}$, where $H$ is the Hubble parameter during preheating, and generally somewhat smaller than the scale of inflation. Thus, lowering the inflationary scale reddens the spectrum and makes it easier to observe.  This is in contrast to the primordial inflationary spectrum, which is roughly scale invariant and becomes harder to observe as inflationary scale is lowered.    We now ask what we can learn about inflation if we detect a spectrum of gravitational waves generated during preheating. Its two most basic features, the peak frequency and the amplitude, represent the reheat and the oscillation scales respectively.  As the reheat scale is lower than the inflation scale, its detection would impose a lower bound on the inflationary scale. Its usefulness as a probe of inflation is amplified if we have a separate probe of the scale of inflation, say from the  CMB B-mode observations.  The oscillation scale is harder to interpret, as it is often highly model dependent. For single scalar field inflation, such as the models considered here,  knowledge of the reheat scale would constrain the coupling constant $g^2$. More optimistically, the \\emph{structure} of the spectrum encodes information about resonance and preheating, so if we can predict the structure accurately we potentially probe the detailed mechanics of preheating. Such an endeavour will require more careful computations and simulations than we present in this paper.  Further progress on this problem can be made in two ways \\cite{inpreparation}. The first is to further refine the code to accommadate a larger class of models, particularly hybrid inflation models which have an essentially arbitrary inflationary scale. Alternatively, as alluded to in Section (\\ref{sect:production}), a more sophisticated theoretical calculation would be to directly solve the evolution equations for the off-diagonal parts of the perturbed Einstein tensor, which are sourced by $T_{ij}$.  This would avoid any ambiguity concerning the use of a formula that is only strictly applicable in flat space, and it would avoid the need to run the code for a finite number of ``boxes'', since we would only need to take a Fourier transform at the end of the computation.  Investigating these gravitational waves is timely, since there is currently considerable interest in the direct detection of gravitational waves.  At the moment, several proposals based on different technologies are being actively pursued. At the solar-system scale, the space-based interferometer LISA \\cite{LISA} will probe frequencies from $10^{-2}$ Hz, which is probably too small for the gravitational waves we are considering. The proposed BBO \\cite{Phinney} and also the Deci-hertz Interferometer Gravitational Wave observatory (DECIGO) \\cite{Seto:2001qf} missions are sensitive to frequencies on the order of $1$ Hz, and would probe gravitational waves arising from preheating after TeV scale inflation.  The array of terrestrial interferometers also probes frequencies corresponding to preheating following low-scale inflation. These experiments include GEO600 \\cite{GEO600},  LIGO  \\cite{LIGO}, TAMA \\cite{TAMA} and VIRGO \\cite{VIRGO}. These are sensitive to scales between $100\\sim 1000$ Hz and may be able to probe a stochastic background  in the interesting range $\\Omega_{gw}h^2 \\sim 10^{-10}$, if they are correlated. At even higher frequencies in the KHz range, we have a slew of resonant bars detectors   \\cite{ALLEGRO,AURIGA,EXPLORER,Blair:1997as}. Once correlated, these resonant bars have a potential to reach $\\Omega_{gw}h^2\\approx 10^{-5}$ \\cite{Maggiore:1999vm},  which puts the signals we see here out of their reach. However, theoretical studies suggest that correlating hollow spherical detectors may eventually allow us to reach $\\Omega_{gw}h^2\\approx 10^{-9}$ \\cite{Coccia:1997gy}.  At even higher frequencies from $10^{3} \\sim 10^{5}$ Hz, there has been a proposal to build a superconducting resonant cavity detector called the Microwave Apparatus for Gravitational Wave Observation (MAGO) \\cite{Pegoraro:1977uv,Reece:1984gv,Ballantini:2005am}. Although the strain sensitivity $\\tilde{h}_f \\approx 10^{-21} \\textrm{Hz}^{-1/2}$ for the prototype is expected to be comparable to resonant bar detectors at $4\\times 10^{3}$ Hz, the large $f^3$ suppression from the relation $\\Omega_{gw} h^2 \\propto \\tilde {h}_f^2 f^3$ means that at these frequencies we can only reach $\\Omega_{gw} h^2\\approx {\\cal{O}}(1)$. An improvement of $5\\sim 6$ orders of magnitude in the strain sensitivity is needed to reach $\\Omega_{gw}h^2\\approx 10^{-10}$, a possibility which may be achieved by further refinements to the prototype and/or construction of an array of such detectors \\cite{Ballantini:2005am}. By way of comparison, we note that the best hope for observing a primordial gravitational wave background is currently provided by BBO, which has a lead time of at least 20-25 years. In this context hoping for a large extrapolation of detector technologies at high frequencies is perhaps not excessively optimistic.    In Figure~\\ref{fig:sensitivityplot} we sketch the sensitivities of the leading interferometric detectors along with the expected stochastic background produced during preheating for the models we discuss in detail here.  Finally, if one was to ever make a concerted attempt to detect a gravitational wave spectrum associated with preheating, one would need to be understand other potential sources that could supply a stochastic background of gravitational waves. This includes any first order phase transition in the early universe (such as the electroweak scale), or decays from cosmic strings. In addition, the presence of a rising component in the spectrum illustrates the dangers of using a (locally) positive spectral index of any detected stochastic gravitational wave background to rule out inflation in favour of alternative cosmogenesis ideas such as ekpyrosis \\cite{Khoury:2001wf,Boyle:2003km} or pre-big bang scenarios \\cite{Gasperini:1992em}.   The potential for gravitational waves to provide a clean probe of inflation has rightfully drawn considerable attention, and strongly motivates attempts to detect the primordial gravitational spectrum. However, cosmic evolution is seldom tidy and gravitational waves are produced as long as large inhomogeneities are present. Preheating is a mechanism which will generate large inhomogeneities, and will necessarily be accompanied by the generation of a stochastic background of gravitational waves. The challenge now is to better determine their properties, and to assess possible strategies for their detection."
    },
    "0601/astro-ph0601663.txt": {
        "abstract": "We extend the model of \\citet{2004MNRAS.348..111K} for variability in black hole accretion discs, by taking proper account of the thermal properties of the disc. Because the degree of variability in the \\citet{2004MNRAS.348..111K} model depends sensitively on the ratio of disc thickness to radius, $H/R$, it is important to follow the time-dependence of the local disc structure as the variability proceeds. In common with previous authors, we develop a one-zone model for the local disc structure. We agree that radial heat advection plays an important role in determining the inner disc structure, and also find limit-cycle behaviour. When the stochastic magnetic dynamo model of \\citet{2004MNRAS.348..111K} is added to these models, we find similar variability behaviour to before.  We are now better placed to put physical constraints on model parameters. In particular, we find that in order to be consistent with the low degree of variability seen in the thermal disc component of black hole binaries, we need to limit the energy density of the poloidal field that can be produced by local dynamo cells in the disc to less than a few percent of the energy density of the dynamo field within the disc itself. ",
        "introduction": "Accretion powered x-ray sources, both on the galactic (AGN) and stellar (X-ray binaries) scales, display significant aperiodic variability, or flickering, on a broad range of timescales \\citep[see for example][]{1988MmSAI..59..239M,1990A&A...230..103B,1994ApJS...92..511V, 2003ApJ...593...96M,2003MNRAS.345.1271V,0306213}. The origin of this variability is not well understood. However, \\citet{2004MNRAS.347L..61U} and \\citet{2005MNRAS.359..345U} have pointed out that the existence of a strong linear relationship between the amplitude of the X-ray variability and the amplitude of the x-ray flux \\citep{2001MNRAS.323L..26U} -- the rms-flux relation -- can be used as a diagnostic to distinguish between various models. For example \\citet{2005MNRAS.359..345U} emphasise that the standard simple shot-noise models, where the light curve is produced by the summation of randomly occurring shots, or flares, cannot explain such a relation. Both \\citet{2004MNRAS.347L..61U} and \\citet{2005MNRAS.359..345U} note that such a relation can be produced naturally by the model of \\citet{1997MNRAS.292..679L} where the variability is produced by variations in the accretion rate occurring at different radii which propagate inwards, and so modulate the energy release in the central x-ray emitting region. The major problem with this model, as noted by \\citet{1997MNRAS.292..679L}, is the physical origin of the accretion rate variations. In order for the model to work, as also emphasised by \\citet{2001MNRAS.321..759C}, it is necessary for the timescale of the variations at each radius to be at least as long as the viscous timescale at that radius. The most obvious origin for fluctuations in an accretion disc are turbulent or hydro-magnetic, perhaps associated with the disc dynamo, and these take place on the local dynamical timescale ($\\sim \\Omega^{-1}$, where $\\Omega$ is the angular velocity), which is shorter than the local viscous timescale by a factor of $\\sim \\alpha (H/R)^2$, where $\\alpha$ ($<$1) is the viscosity parameter, and $H/R < 1$ the disc opening angle \\citep{1981ARA+A..19..137P}.  A solution to this problem, in the form of an explicit physical model, has been recently proposed by \\citet{2004MNRAS.348..111K}, based on ideas put forward by \\citet{2003ApJ...593..184L}. They suggest that the local dynamo processes in the disc can affect the accretion rate by driving angular momentum loss in the form of an outflow (wind or jet). They model the dynamo as a small-scale stochastic phenomenon, operating on roughly the local dynamical timescale. They then postulate that large-scale outflow can only occur when small-scale random processes in neighbouring disc annuli give rise by chance to a coherent large-scale field. This occurs on much longer timescales, by a factor of order $2^{R/H}$ \\citep{2003ApJ...593..184L}. This also provides an explanation for why the dominant flickering frequencies are typically much longer than the dynamical timescales at the centre of the disc where most of the energy is released.  However, \\citet{2004MNRAS.348..111K} took a very idealised model for the disc in that they assumed that the disc had a constant thickness ratio of $H/R$. This simplified the computations for two main reasons. First, local disc structure equations could be ignored, because there was no need to compute the local value of disc thickness. Second, because the local disc dynamos were assumed to be spatially independent on a radial scale of $\\Delta R \\sim H$, there was no need to vary the number of independent dynamos as the disc thickness varied with time. In this paper we take the first steps towards remedying this deficiency. For the time being we concern ourselves with a standard (optically thick, geometrically thin) accretion disc, and thus any conclusions we have will be relevant mainly to the thermally dominated (TD) state of the X-ray transients \\citep[e.g.][]{0306213}. We leave the extension to the more interesting, and more variable (low/hard) state, containing both disc and corona, to future work. In Section~\\ref{sect:methods} we explain how we generalise the work of \\citet{2004MNRAS.348..111K} by including equations to compute the local disc structure, and also explain how we vary the grid spacings to follow the radial size of the local disc dynamo cells in such a way as to minimise unwanted numerical mixing. In Sections~\\ref{sect:stationary} and~\\ref{sect:timedep} we investigate the models which result from our equations when there is a steady external accretion rate, and in Section~\\ref{sect:review} compare our results with previous work in the field. In Section~\\ref{sect:flickering} we add the variability according to the stochastic magnetic model of \\citet{2004MNRAS.348..111K}. In Section~\\ref{sect:obs}, we discuss the observational data on the variability of the thermal disc component, to which our models are relevant, and in Section~\\ref{sect:res} we show that for expected values of the model parameters our models are consistent with the data. We present brief conclusions in Section~\\ref{sect:conclusions}. ",
        "conclusions": "\\label{sect:conclusions}  We have generalised the model of \\citet{2004MNRAS.348..111K} for variability in black hole accretion discs by including proper consideration of the local disc structure. We take the disc properties to depend only on radius $R$, and so, in common with other authors, make a local one-zone approximation for the disc structure.  In the absence of the addition of a stochastic magnetic dynamo, which drives the flickering, we mainly reproduce results similar theoretical work has shown.  We agree that taking account to advective heat flow is of crucial importance in determining the structure of the inner regions, but find that the assumption of strictly Keplerian angular velocities, and the neglect of radial pressure gradients, are generally reasonable assumptions.  When a stochastic dynamo is included we find behaviour similar to that described by \\citet{2004MNRAS.348..111K}. However, because the degree of variability of the thermal disc is observed to be small (less than around 1 per cent), if it is detected at all, we can only draw limited conclusions. Our main finding is that for consistency we require that the strength of the poloidal field generated by any radial disc dynamo cell must be at least an order of magnitude smaller than the field generated by the dynamo within the disc. The simulations reported here however are of help to understand the influence of the parameters on the variability.  In this paper we have restricted ourselves to modelling just the thermal disc component, and so have had to restrict our attention to the high/soft or thermally dominant state of black hole binaries. In future work we shall begin to include consideration of a hot corona, in addition to the cooler thermal disc, so that we can begin to model the more highly variable low/soft state of the black hole binaries, as well as the X-rays from AGN \\citep{2005MNRAS.359..345U}."
    },
    "0601/astro-ph0601329_arXiv.txt": {
        "abstract": "{There is mounting evidence from observations of long duration gamma ray bursts (GRBs), supernova remnants (SNR) and the supernova (SN) explosion 1987A, that SN explosions eject highly relativistic bipolar jets of plasmoids (cannonballs) of ordinary matter. The highly relativistic plasmoids sweep up the ambient matter in front of them, accelerate it to cosmic ray (CR) energies and disperse it along their long trajectories in the interstellar medium, galactic halo and intergalactic space. Here we use the remarkably successful cannonball (CB) model of GRBs to show that bipolar jets from Galactic SN explosions can produce the bulk of the Galactic cosmic rays at energies below the ankle, while the CRs which escape into the intergalactic space or are deposited there directly by jets from SNe in external galaxies can produce the observed cosmic ray flux with energies above the ankle. The model predict well all the observed properties of cosmic rays: their intensity, their spectrum including their elemental knees and ankles, their composition and the distribution of their arrival directions.  At energies above the CR ankle, the Galactic magnetic fields can no longer delay the free escape of such ultra high energy CRs (UHECRs) from the Galaxy. These UHECRs, that are injected into the intergalactic medium (IGM) by the SN jets from our Galaxy and all the other galaxies and are isotropized there by the IGM magnetic fields, dominate the flux of UHECRs. Almost all the extragalactic UHECRs heavier than helium photo-disintegrate in collisions with the far infrared (FIR), microwave and radio background radiations. The surviving CR protons and He nuclei suffer a Greisen-Zatsepin-Kuzmin (GZK) cutoff due to pion photo-production in collisions with the FIR, microwave and radio background photons.}  \\normalsize\\baselineskip=15pt ",
        "introduction": "\\noindent Cosmic rays (CRs) were discovered by Victor Hess in 1912. Today, 93 years later, their origin is still debated.  CRs have been studied in experiments above the atmosphere, in the atmosphere, on the ground, underground and in space. Their energies cover an enormous range, from sub GeV to more than a few $10^{11}\\, GeV\\,,$ over which their differential flux decreases by roughly 33 orders of magnitude. Any successful theory of the origin of CRs must explain their main observed properties near Earth (for recent reviews see Biermann \\& Sigel 2002; Olinto 2004; Cronin 2004 Watson 2004; Hoerandel 2005; Engel 2005):  \\noindent {\\bf The CR energy spectrum} shown in Fig. 1, has been measured up to hundreds $EeV$. It can be approximated by a broken power-law, $dn/dE\\sim E^{-p}\\,,$ with a series of breaks near $ 3\\, PeV$ known as `the knee', a second `knee' near 200 PeV and an `ankle' near $3\\, EeV\\, $. The power-law index changes from $p\\sim 2.67\\, $ below the knee to $p\\sim 3.05$ above it and steepens to $p\\sim 3.2$ at the second knee. At the ankle the spectral index seems to decrease to $p\\sim 2.7\\, .$ The spectral behaviour above $50\\ EeV$ is still debated (see e.g., Olinto 2004; Cronin 2004; Watson 2004).   \\noindent {\\bf The CR elemental composition} is known well only from direct measurements on board satellites which run out of statistics well below the knee energy. The measured composition of high energy CRs is highly enriched in elements heavier than hydrogen relative to that of the solar system.  The enrichment increases with atomic number and with CR energy almost up to  the second knee beyond which it appears to decline (e.g., Kampert et al.~2004; Hoerandel, 2004 ). as shown in Figs 4,5. It is still debated at energies above the ankle (e.g. Cronin 2004; Watson 2004).  The detailed elemental composition at the knee and above it is known only very roughly (e.g. Hoerandel 2004, 2005; Cronin 2004; Watson 2004; Engel 2005).   \\noindent {\\bf The CR arrival directions} at energy below $\\sim EeV$ are isotropized by the Galactic magnetic fields. Only at energies well above $EeV\\,,$ their arrival directions may point towards their Galactic sources and only at extremely high energies they may point towards nearby extragalactic sources. Initial reports of deviations from isotropy in the arrival directions of such cosmic rays with energy above $EeV$ and of some clustering in their arrival directions are still debated (e.g. Cronin 2004).  \\noindent {\\bf The Galactic cosmic ray luminosity} has been estimated to be between $10^{41}\\, erg\\, s^{-1}$ (e.g., ) and more than $10^{42}\\, erg\\, s^{-1}\\, $ (see, e.g. Dar \\& De R\\'ujula~2001) depending on the size of the Galactic CR halo.    It is widely believed that Galactic CRs with energy below the knee are accelerated mainly in Galactic supernova remnants (SNRs). The opinions on the origin of CRs with energy between the knee and the ankle are still divided between Galactic and extragalactic origin.  CRs with energy above the ankle are generally believed to be extragalactic in origin because they can no longer be isotropized by the Galactic magnetic fields while their observed arrival directions are isotropic to a fair approximation. Yet Galactic origin is not ruled out -- they may be produced by the decay of unknown massive particles or other sources which are distributed isotropically in an extended Galactic halo (e.g. Plaga 2002). An observational proof that the CRs above the ankle are extragalactic in origin, such as arrival directions which are correlated with identified extragalactic sources or a Greisen-Zatsepin-Kuzmin (GZK) cutoff due to $\\pi$ production in their collisions with the cosmic microwave background radiation (MBR), are still lacking (Cronin 2004). Moreover, there is no single solid observational evidence which supports the SNR origin of the bulk of CRs below the knee (see, e.g. Plaga 2002 and references therein).  In fact, the evidence from $\\gamma$-ray astronomy, x-ray astronomy and radio astronomy strongly suggest that {\\bf SNRs are not the major accelerators of CRs with energy below the knee}:  \\begin{itemize} \\item{} {\\bf SNR origin cannot explain the Galactic CR luminosity:} Radio emission and X-ray emission from SNRs provide strong evidence for acceleration of high energy electrons in SNRs. Some SNRs were also detected in TeV $\\gamma$-rays which could be produced by the prompt decay of $\\pi^0$'s from collisions of high energy hadronic cosmic rays with the ambient protons and nuclei in and around the SNR. However, recent observations of SNRs inside molecular clouds (and careful analysis of Tev emission in others, such as SN1006) show that the time-integrated CR luminosity of these SNRs does not exceed a few $10^{48}\\, erg$ unless an extremely small mean baryon density is assumed for the production region. The rate of Galactic SN explosions, which is estimated to be $\\sim 1/50\\, year\\, ,$ yields a total CR luminosity which falls short by  two orders of magnitude  than the estimated luminosity of the Milky Way (MW) in CRs, $L_{CR}[MW] > 10^{41}\\, erg\\, s^{-1}\\, .$ \\item{} {\\bf SNR origin cannot explain the diffuse Galactic GBR:} The interactions of CR electrons with ambient photons and of CR nuclei with ambient nuclei in the interstellar medium produce a diffuse background of $\\gamma$-rays. Such a diffuse gamma background radiation (GBR) has been detected by EGRET and Comptel on board the Compton Gamma Ray Observatory (CGRO). However, the scale length of the distribution of SNR in the Galactic disk is $\\sim 4\\, kpc$ and cannot explain the scale length of the observed GBR, which is larger by more than an order of magnitude. In particular, energetic electrons cool rapidly by inverse Compton scattering of stellar light and of microwave background photons. Over their cooling time they cannot reach by diffusion far enough from the SNRs to explain the intensity of GBR at large Galactic distances.  \\item{} {\\bf SNR origin cannot explain the diffuse radio emission from galaxies and clusters:} Because of their fast cooling in the MBR the electrons from SNRs, which are mainly located in the galactic disk, cannot reach large distances from the galactic disk by diffusion or galactic winds. The radio emission from our Galaxy, from edge-on galaxies and from the intergalactic space in clusters of galaxies provide evidence for high energy electrons at very large distances from the galactic disks where most of the SNRs are located. \\end{itemize}  \\noindent  Although supernova remnants (SNRs) do not seem to be the main source of CRs in our Galaxy, in external galaxies and in the intergalactic medium (IGM), still SN explosions may be the main source of CRs at all energies if SNe emit highly relativistic bi-polar jets which produce the visible GRBs when they point in our directions (Dar and Plaga 1999)\\footnote{The association between GRBs and high energy CRs was first suggested by Dar et al. (1992). Waxman (1995), Vietri (1995) and Milgrom (1995) suggested that extragalactic GRBs are the main accelerators only of ultra-high energy cosmic rays (UHECRs). However, following the first observational evidence for a GRB-SN association (Galama et al 1998), Dar \\& Plaga (1999), suggested that the bulk of the cosmic rays {\\bf at all energies} are accelerated in bipolar jets that are ejected in Galactic SN explosions which produce Galactic GRBs, most of which are beamed away from Earth.}. Here, I will outline briefly a simple theory of the origin of CRs at all energies which is based on the extremely successful cannonball (CB) model of GRBs (Dar \\& De R\\'ujula~2000,2004). I will show that it explains remarkably well the main observed properties of Galactic and extragalactic CRs.  The complete theory and a rigorous derivation of the main properties of CRs from the CB model will be published elsewhere  (Dar \\& De R\\'ujula~2005b). According to this theory:  \\begin{itemize} \\item{} The bulk of the CR nuclei are accelerated by the highly relativistic bipolar jets of plasmoids (cannonballs) of ordinary matter ejected in SN explosions (which produce GRBs, most of which do not point towards Earth). These jets can accelerate swept in ISM particles to the highest observed cosmic ray energies\\footnote{Jets from microquasars and active galactic nuclei may also contribute to CR acceleration in galaxies and in clusters of galaxies, but they will not be discussed here.}.  \\item{} The elemental knees  are the maximal lab energy that ISM nuclei acquire by elastic magnetic scattering in the CBs' rest frame.  The knee in the al-particle spectrum is the maximal lab energy of CR protons and the `second' knee is that of the iron group nuclei and heavier metals.  \\item{} The Galactic cosmic rays escape outside the Galaxy by diffusion in the Galactic magnetic fields. Their energy-dependent confinement time steepens their Galactic spectrum. The ankle is the energy where the Galactic magnetic fields can no longer isotropize their directions and delay their free escape into the IGM.  \\item{} The CRs from SN explosions that have escaped the galaxies into the IGM, or were deposited there directly by the jets from galactic SNe, produce the bulk of the CRs in the intra cluster medium in galaxy clusters and in the IGM outside clusters.  They have been accumulating there over the Hubble time and were isotropized by the IGM magnetic fields. Above the ankle, the Galactic magnetic fields cannot shield it from UHECRs. Thus most of he UHECRs nuclei in the Galaxy with energy above the ankle are extragalactic in origin.   \\item{} The galactic CRs which escape into the IGM, or are directly deposited there, have the injection spectral index $p\\approx 2.17$ below the elemental knees.  \\item{} The spectrum of the extragalactic UHECRs in the IGM is modified by the cosmic expansion and by their interaction with the far infrared, microwave and radio background radiations. The effective thresholds for pair production and photo-disintegration are well below the CR ankle but the Galactic magnetic field and Galactic winds probably prevent their penetration at energies well below the ankle. At energies above the ankle, the UHECRs which cross through the Galaxy dominate the CR flux. Most of the UHECR nuclei heavier than $He$ are destroyed by photo-disintegration in less than a Hubble time. The bulk of the UHECR protons and $He$ nuclei in the IGM suffer a GZK suppression due to $\\pi$ production and only those from nearby extragalactic sources may reach the Galaxy.  \\item Direct CR deposition in the IGM, by CBs and their associated CR jets and by CR diffusion and galactic winds, stir it up and generate the turbulent IGM magnetic fields (Dar and De R\\'ujula 2005a).  \\item{} The CR electrons, in external galaxies and in the IGM produce the bulk of the extragalactic diffuse gamma-ray background radiation (GBR) by inverse Compton scattering of the cosmic microwave background (MBR) (Dar and De R\\'ujula 2000). \\end{itemize}  \\noindent ",
        "conclusions": "\\noindent  The CB model of GRBs was used here to demonstrate that the bipolar jets from SN explosions can be the main source of cosmic rays at all energies. The model correctly predicts, within the experimental uncertainties, the observed all-particle CR flux, its energy spectrum and its elemental composition at all energies.  In particular, the model predicts that the CRs above the ankle are extragalactic in origin, are mostly protons $+He$ nuclei, are highly isotropic and their spectrum exhibit the GZK suppression due to their interaction with the microwave and radio background radiations. It appears like the 93 years old puzzle of the origin of galactic CRs has been solved. However, more precise CR data on their energy spectrum and composition, in particular near the CR knees and well above the GZK cutoff, as well as astrophysical data on the elemental abundances in SNRs and SBs are needed before a firm conclusion can be made."
    },
    "0601/astro-ph0601435_arXiv.txt": {
        "abstract": "We present mid-infrared $N$- and $Q$-band photometry of the Galactic Center from images obtained with the mid-infrared camera VISIR at the ESO VLT in May 2004. The high resolution and sensitivity possible with VISIR enables us to investigate a total of over 60 point-like sources, an unprecedented number for the Galactic Center at these wavelengths. Combining these data with previous results at shorter wavelengths \\citep{viehmann2005} enables us to construct SEDs covering the $H$- to $Q$-band regions of the spectrum, i.e. 1.6 to 19.5~$\\mu$m. We find that the SEDs of certain types of Galactic Center sources show characteristic features. We can clearly distinguish between luminous Northern Arm bow-shock sources, lower luminosity bow-shock sources, hot stars, and cool stars. This characterization may help clarify the status of presently unclassified sources. ",
        "introduction": "\\label{INTRO}  The Galactic Center is known as a bright source of near- and mid-infrared radiation since the late 1960s \\citep{becklin1968, becklin1969, low1969}. The main source of near-infrared radiation is photospheric emission from a dense stellar cluster, i.e.\\ a crowded field of point sources, while almost all of the mid-infrared radiation originates from extended gas and dust features as well as dust emission from the circumstellar regions of a dozen individual sources interacting with the more extended Galactic Center interstellar medium.  The nature of the brightest mid-infrared sources is more or less clear -- IRS~1W, 2, 5, 10W, and 21 are bow-shock sources, caused by bright stars with strong winds plowing through the ambient gas and dust of the Northern Arm \\citep{tanner2002, tanner2003, tanner2005, rigaut2003, eckart2004, geballe2004}. Of the fainter sources, IRS~7 is an M2 supergiant, however, most of the others remain enigmatic.  Combined photometry at near- and mid-infrared wavelengths is an ideal tool to investigate the nature of the bright mid-infrared sources, since it gives information on both the mid-infrared sources themselves and the stellar sources that could be powering them. We therefore present $N$- and $Q$-band photometry of the Galactic Center, obtained with the ESO VLT, in order to further investigate the nature of the mid-infrared sources. Additional data at shorter wavelengths is taken from \\citet{viehmann2005} to provide a more complete picture.  Previous studies at mid-infrared wavelengths have been severely limited in resolution (\\citealt{becklin1978} $\\sim$2\\arcsec, \\citealt{gezari1985} $\\ga$1\\arcsec, \\citealt{stolovy1996} $\\sim$0.7\\arcsec), however, since the availability of the mid-infrared camera VISIR at the ESO VLT, the Galactic Center can be studied in great detail in the  $N$- and $Q$-bands (8-13~$\\mu$m and 17-24~$\\mu$m), with a resolution of 0.3\\arcsec to 0.6\\arcsec (the diffraction limit of an 8~m telescope), provided that the seeing is good enough. This corresponds to a linear resolution of 2500 to 5000~AU at the distance of 8~kpc to the Galactic Center. ",
        "conclusions": "\\label{conc}  We have presented mid-infrared photometry of the Galactic Center from images obtained with the VISIR camera at the ESO VLT. In combination with shorter wavelength data \\citep{viehmann2005}, we were able to construct SEDs for an unprecedented number of sources (for the Galactic Center) in this wavelength regime. We find that different source types show characteristic differences in the shape of their SEDs, which should help in the preliminary classification of compact mid-infrared sources. In particular, we find that additional, i.e. intrinsic, absorption at approximately 10~$\\mu$m is a sign of dust formation. Many of the brightest Galactic Center mid-infrared sources, however, appear to be bow-shock powered and do not show dust formation features in their SEDs. In addition, we find that hot and cool stars show different $M-m_{8.6\\mu\\mathrm{m}}$ colors, with the cool stars showing redder colors, as is to be expected.  We have reported on four fairly bright $N$-band point sources located to the east of the bright Northern Arm source IRS~5. The most likely explanation for their unusual appearance at mid-infrared wavelengths is that these sources are bow-shock sources of lower luminosity than the typical Northern Arm sources.  The main uncertainty in our results is due to calibration uncertainties as a result of the lack of suitable reference observations. Therefore, mid-infrared observations of the Galactic Center remain necessary in order to obtain more complete information on the dusty mid-infrared sources in this intriguing environment."
    },
    "0601/astro-ph0601717.txt": {
        "abstract": "We present the first multi-parameter analysis of the spatial clustering properties of nearby galaxies in which accretion activity is present to different degrees.  We spectroscopically identify and classify active galactic nuclei (AGN) from the Sloan Digital Sky Survey (SDSS) DR2 main galaxy sample, which yields the most precise measurement to date of AGN clustering. % Estimates of the redshift space two-point correlation function (CF) reveal that Seyferts are clearly less clustered than normal galaxies, while the clustering amplitude of LINERs is consistent with that of the parent galaxy population. The similarities of the distributions of host properties (color and concentration index) of Seyferts and LINERs suggest that the difference in their clustering amplitudes is not driven by the morphology-density relation.  We find that the luminosity of the [\\ion{O}{1}] emission shows the strongest influence on AGN clustering, with low $L_{\\rm [O I]}$ sources having the highest clustering amplitude, $s_0$.  This trend is much stronger than the previously detected dependence on $L_{\\rm [O III]}$, which we confirm. % The fact that objects of given spectral types are clustered differently seems correlated with a variety of their physical properties including $L_{\\rm [O I]}$ , $L_{\\rm [O III]}$, the electron density of the emitting gas $n_e$, and the obscuration level.  LINERs, which exhibit high $s_0$, show the lowest luminosities and obscuration levels, and relatively low $n_e$, suggesting that these objects harbor black holes that are relatively massive yet weakly active or inefficient in their accretion, probably due to the insufficiency of their fuel supply.  Seyferts, which have low $s_0$, are very luminous in both [\\ion{O}{1}] and [\\ion{O}{3}] and show large $n_e$, suggesting that in these systems the black holes are less massive but accrete quickly and efficiently enough to clearly dominate the ionization. Star-forming galaxies -- the H {\\sc ii}'s -- are weakly clustered; because these systems are hosted mainly by blue, late type galaxies, this trend can be understood as a consequence of both the morphology-density and star formation rate-density relations. However, the spectral properties of the H {\\sc ii}'s suggest that these systems hide in their centers, amidst large amounts of obscuring material, black holes of generally low mass whose activity remains relatively feeble. % Our own Milky Way may be a typical such case. ",
        "introduction": "Clustering analysis can be a useful tool to employ in investigating the formation and evolution mechanisms of various astrophysical objects.  In the case of AGN, the pattern of spatial clustering and its implications for the environments of AGN provides constraints on galaxy evolution models, and in particular on how well AGN trace the normal, non-active galaxy distribution.  Comparisons of the AGN environments with those of control samples of regular galaxies may indicate differences that reveal what determines the activity in galaxy nuclei, and even gauge its strength and time cycle, thus yielding a phenomenological link between black hole growth and galaxy formation and evolution processes.  Estimates of the two-point correlation function (CF) provide important diagnostics for the black hole (BH) masses and duty cycle of fueling in AGN.  If structure formation is hierarchical, then the rarest and most massive virialized dark matter (DM) halos cluster the most strongly \\citep{kai84}, thus the clustering strength of a given object type constrains the mass and number density of the halos in which they reside. The DM halo mass can further be connected to the physical properties of the resident galaxies and their central black holes through empirical scaling relations: the black hole mass -- stellar velocity dispersion $M_{\\rm BH} - \\sigma_*$ relation that holds for normal galaxies \\citep{geb00, fer00} as well as for the active ones \\citep{geb00b, fer01, nel04, onk04, gre05}, the correlations between bulge properties and $M_{\\rm BH}$ \\citep{mag98, gra01}, and the correlation between circular velocity and central velocity dispersion \\citep{fer02, bae03} that links dark matter halos and supermassive BHs.  The number density of the DM halos can be used to infer the duty-cycle, or the accretion life-time, as this is given by the inverse ratio of this number density to that of the observed objects \\citep{mar01}.  Such predictions for the AGN's lifetime, and its implications for the efficiency of BH formation \\citep{hai00}, can test the fundamental assumptions of semi-analytical models of AGN and galaxy formation \\citep{kau00, mat03}.  Studying the spatial clustering properties of different spectroscopically-classified AGN species may link their ionization mechanisms to properties of their hosts and environment.  Through comparison of the clustering amplitudes of accreting-type sources such as Seyferts and the starlight-ionized H {\\sc ii} galaxies, it is possible to contrast properties of the dark matter halos that host different kinds and levels of galactic nuclear activity.  Particularly interesting is the clustering of the ubiquitous Low Ionization Emission-Line Regions (LINERs, e.g. Heckman 1980), whose physical origin, whether thermal or non-thermal, remains ambiguous despite exhaustive analysis of their properties \\citep{ho99,ho03,fil04,mao05,con05}.  Galaxy mergers are predicted to play a crucial role in triggering AGN activity \\citep{bar92, kau00, mat05}.  If galaxy nuclei become more AGN-like when major mergers occur, then active galaxies of different spectral signatures, corresponding to different degrees of contributions from a non-thermal ionization process, may cluster differently, reflecting variation in their environments. Thus, a comparison of the clustering properties of the active and non-active galaxies is an important tool for testing the merger-accretion connection.  Through its wealth of both spectroscopic and photometric data, SDSS \\citep{yor00} is ideal for investigating the dependence of clustering properties on a variety of emission characteristics.  The SDSS observations offer the unique opportunity to examine simultaneously the dependence of environment on the spectral classification of emission-line galaxies, and on various other features that characterize or even determine a detectable accretion process, in conjunction with properties of their host galaxies.  Previous studies with similar aims \\citep{kau03, kau04, mil03, wak04} compared passive galaxies with AGN as a general class and did not investigate the properties of sub-populations such as Seyferts and LINERs separately. In these studies, the optical classification of an object as an AGN relies, especially among the narrow-line emitters, mainly on the [\\ion{N}{2}]/H$\\alpha$ vs. [\\ion{O}{3}]/H$\\beta$ diagnostic diagram. There are suspicions that such rather lenient definitions may not best represent the accretion dominated population ~\\citep{hao05}. Using such definitions, the AGN population is found by \\citet{mil03} and \\citet{wak04} to occupy a uniform fraction of galaxies, to be common to a large range of galactic environments, and to follow the distribution of the whole galaxy population, and therefore to be unbiased tracers of mass in the universe.  However, when divided by their ([\\ion{O}{3}]) luminosities, the active nuclei show markedly different clustering properties \\citep{wak04}, different star-formation histories in their hosts \\citep{kau03}, and different environmental preference \\citep{kau04}.  Combined together, these findings seem to suggest that Seyferts and LINERs, that represent quite distinct distributions in $L_{\\rm [O III]}$ \\citep{kau03}, might also cluster differently.  Such inferences are intriguing and thus merit further clarification.  It is thus crucial to establish and quantify the connections between the intrinsic physical properties of the nuclear line-emitting regions and those of their host galaxies and dark matter halos. We pursue these issues here by including important spectral diagnostics that probe the astrophysics of AGN.  In this paper, we use SDSS DR2 data in an attempt to detect and analyse differences in the clustering properties of (1) active and non-active galaxies, (2) objects of distinct types of activity, i.e., accretion versus star-formation, and (3) different spectroscopic types of AGN, e.g., Seyferts versus LINERs.  In order to carry out such analysis with a good understanding of the possible biases introduced by detection and spectral identification, we present in Section ~\\ref{data} a systematic examination of the different AGN selection criteria and their sensitivity to the dominance of an accreting-type nuclear power source.  In Section ~\\ref{csi} we estimate and compare the redshift-space CF for subsamples of objects of different spectral classes (Section ~\\ref{class}), and investigate the degree to which these measurements are influenced by the level of nuclear activity, as traced by the [\\ion{O}{3}] and [\\ion{O}{1}] luminosities (Section ~\\ref{csio1o3}), properties of the line-emitting gas, e.g., intrinsic obscuration, as traced by the neutral hydrogen column densities $N_H$ and the electron density $n_e$ (Section ~\\ref{ne}), and host properties (Section ~\\ref{host}).  We discuss in Section ~\\ref{discussion} possible implications of our findings on the features that distinguish Seyferts from LINERs, on the peculiar nature of the Transition objects, that border the definition of AGN (Seyferts or LINERs) and H {\\sc ii}, and on the potentially great similarity between the Galactic Center and a typical H {\\sc ii} system. Section ~\\ref{summary} summarizes the findings and conclusions of this work.     %% differences in clustering detected with relatively high precision. ",
        "conclusions": "\\label{summary}   % summary:  We investigate the spatial clustering of low-luminosity AGN in the nearby universe, using spectroscopic classification and measurement of related physical parameters to examine this clustering as a function of their nuclear emission properties.  We estimate and compare the CFs for sub-samples selected on these various properties, and analyze these results in light of the well-known dependence of clustering on halo mass, as well as the empirical relation between black hole mass and galaxy mass.  This study reveals a strong connection between the galaxy clustering amplitudes, hence the properties of their dark matter halos, and the intrinsic physical properties of these objects' line-emitting regions. A brief overview of our main results is given in Table ~\\ref{tbl-summary}, and we list our conclusions as follows.  (i) The amplitude of the two-point correlation function is strongly dependent on galactic nuclear emission properties.  By employing detailed spectroscopic classification schemes we find that, contrary to previous claims, bona fide low luminosity AGN, spectroscopically classified as Seyfert galaxies, are clearly less clustered than the average galaxies, and therefore do not tracing the same underlying structures.  LINERs, on the other hand, exhibit a clustering amplitude higher than that of Seyferts, and consistent with that of the whole galaxy sample.  (ii) AGN defined using only the [\\ion{N}{2}]/H$\\alpha$ vs. [\\ion{O}{3}]/H$\\beta$ diagnostic diagram are dominated by the LINER-like objects, and include a significant fraction of objects with conspicuously strong nuclear star-formation activity.  Two, or better, three-dimensional diagnostic diagrams that employ forbidden lines whose emissivity is less sensitive to abundance effects, like [\\ion{S}{2}] and/or [\\ion{O}{1}], are necessary to accurately separate sources powered by accretion from those in which star-forming activity is dominant.  It is, therefore, important to emphasize that more lenient AGN definitions are strongly biased toward LINER-like behavior, and consequently, do not represent the traditional AGN activity and physical characteristics.  LINERs differ from Seyferts in many of their intrinsic characteristics.  In LINERs, the emitting gas density is lower than in Seyferts, and the obscuration is generally at very low levels.  As suggested by the large differences we observe in the [\\ion{O}{1}] and [\\ion{O}{3}] line luminosities, the accretion activity is significantly reduced in LINERs while in Seyferts is generally high.  (iii) Emission parameters such as $L_{\\rm [O III]}$ and, in particular, $L_{\\rm [O I]}$, which are sensitive to the level of nuclear accretion activity, play an important role in differentiating systems of different clustering properties. Highly luminous and, therefore, more active objects are significantly less clustered than faint, less active ones; the latter types are strongly represented by LINERs.  (iv) The amount and nature of the fuel available for nuclear activity, whether accretion or star-formation, is also connected to the spatial clustering characteristics.  The less obscured systems and the galaxies with moderately high emitting gas densities are more clustered.  Such sources are highly represented by LINERs.  (v) Host galaxy properties and, consequently, the morphology-density relation, do not strongly influence the difference in the spatial clustering found between Seyferts and LINERs, and play only a secondary role in the large differences between the clustering amplitudes of the low and high $L_{\\rm [O I]}$ and/or $L_{\\rm [O III]}$ subsets.  (vi) Based on the strong relationship between host halo mass, as given by the clustering amplitude, and the spectral properties of the line-emitting galaxies, we speculate that: 1) Seyfert activity arises from active accretion onto small BHs that is abundantly fueled, while 2) LINERs' emission-line properties originate from ionization by slowly accreting massive BHs whose circumnuclear material supply is generally low. % MSV: again, I don't understand the difference between slow and inefficient % accretion. How can their be fast accretion unless it is efficient? %anca: I got carried away...  but I corrected now. In this scenario, H {\\sc ii} systems may be interpreted as harboring relatively small, weakly accreting BHs whose activity remains enshrouded within the surrounding star-forming activity.  The similarity of the Milky Way nuclear system to that of a typical H {\\sc ii} galaxy in this picture provides a means of integrating the Galactic Center into the larger context of galaxy nuclei.  %(vii) The strength of the central activity in the local galaxies %appears related to the crowdness of the environment, on small scales."
    },
    "0601/astro-ph0601573_arXiv.txt": {
        "abstract": "The latest results in the research of forming planetary systems have led several authors to compile a sample of candidates for searching for planets in the vicinity of the sun. Young stellar associations are indeed excellent laboratories for this study, but some of them are not close enough to allow the detection of planets through adaptive optics techniques. However, the existence of very close young moving groups can solve this problem. Here we have compiled the members of the nearest young moving groups, as well as a list of new candidates from our catalogue of {\\it late-type stars possible members of young stellar kinematic groups}, studying their membership through spectroscopic and photometric criteria. ",
        "introduction": "\\label{sec:intr}  In recent years, a series of young stellar kinematic groups (clusters, associations, and moving groups) of late-type stars with similar space motion and ages ranging from 8 to 50 Myr \\citep[see][and references therein]{zuc04a} has been discovered in our neighbourhood: TW~Hya, $\\beta$~Pic, AB~Dor, $\\eta$~Cha, $\\epsilon$~Cha, Tucana and Horologium associations. In addition, several more distant young associations such as MBM~12 \\citep{hea00}, Corona Australis \\citep{qua01}, and possibly the group of stars with a motion similar to that of HD~141569 \\citep{wei00} have also been identified. In the Galactic velocity space, they situate inside the boundaries of the \\object{Local Association} (see Fig.~\\ref{fig:uv}), a mixture of young stellar complexes ---~OB and T-associations~--- and clusters with different ages \\citep{egg75, egg83a, egg83b, m01}. These associations of very young stars are excellent laboratories for investigations of forming planetary systems \\citep{zuc04}. Nevertheless, they are generally situated at distances above 50 pc, which makes them less accessible to adaptive optics systems even on large telescopes.  It is well-known that tightly bound, long-lived open clusters can account for only a few per cent of the total galactic star formation rate \\citep{wie71}. Therefore, either most clusters and associations disperse very quickly after star formation has started or most are born in isolation \\citep{wic03}. The existence of very young moving groups (MGs) with a few dozens of stars showing the same spectroscopic properties ---~i.e. age, metallicity, level of magnetic activity~--- is in agreement with the first explanation. Small associations of stars may be dispersed by galactic differential rotation since they are not gravitationally bounded enough, taking into account that their nucleus consist of only a few stars as in the case of the Ursa Major MG (see King et al. 2003, for a recent review) or the recently discovered AB~Doradus MG \\citep{zuc04}. The location of these young MGs inside the Local Association and its proximity in the $UV$-plane can be explained as the result of the juxtaposition of several star forming bursts in adjacent cells of the velocity field (see Montes et al. 2001, and references therein) or dynamical perturbations caused by spiral waves \\citep{sim04,fam05,qui05}. Thus, one expects to find groups of coeval stars with similar space motion in our neighbourhood.  In 2001, the $\\beta$~Pic~MG \\citep{zuc01} --- a group of stars with an age of $\\sim$~12~Myr \\citep{zuc01, ort04} at a mean distance of $\\sim$~35~pc co-moving with the well-known young star $\\beta$~Pic --- was confirmed to be the closest kinematic group up to date. More recently, \\citet{zuc04} have identified a new group of stars co-moving with the also well-known young star AB~Dor, at a mean distance of $\\sim$~30~pc, and with an age of $\\sim$~50~Myr. Nevertheless, the existence of a nearer association of a few stars was proposed by \\citet{gai98} and studied in detail by \\citet{fur04}, though its existence is quite controversial. Here we discuss about the fact of the nearest MGs using both spectroscopic and photometric criteria of membership for a sample of stars that includes the proposed members from the literature and our list of young cool stars possible members of young stellar kinematic groups \\citep{m01, lop05}. ",
        "conclusions": "\\label{sec:conc}  In Table~\\ref{tab:MGs} we list the stars belonging to the nearest moving groups: Hercules-Lyra Association and AB~Dor MG, and those being part of the Local Association B4 subgroup. For the Hercules-Lyra Association, a division between certain members and candidates with doubtful classification or non members has been made. In the three groups, new candidates from our catalogue of {\\it Late-type Stars Possible Members of Young Stellar Kinematic Groups} \\citep{m01} have been selected because of their kinematics (see $\\S$~\\ref{sec:herc} and $\\S$~\\ref{sec:abdor}). A total of 75 stars including the known members and the new candidates selected by us have been analysed. Kinematic, spectroscopic and photometric criteria have been utilized to discriminate non members from the rest of candidates of the Hercules-Lyra Association, and to distinguish between the members of the AB~Dor MG and those of the B4 subgroup.  In the velocity space, Hercules-Lyra is clearly distinguishable from the rest of the sample (see Figs.~\\ref{fig:uv} and~\\ref{fig:uvw_MGs}). The dispersion in $U$, $V$, and $W$ is comparable with that of other coeval MGs such as Castor and Ursa Major \\citep[e.g.][]{m01}, and compatible with its age (see $\\S$~\\ref{sec:herc}). On the other hand, AB~Dor~MG and B4 subgroup are mixed up and age-dating criteria are necessary to distinguish between the members of both groups. Nevertheless, the dispersion in $W$ for AB~Dor MG is quite smaller than the one of B4 subgroup. Age-dating criteria are also necessary to discriminate non members of Hercules-Lyra from the certain ones. The results of applying them are summarized in Table~\\ref{tab:MGs}: the Hercules-Lyra Association is formed by 10 certain members situated at a mean distance of $\\sim$~20~pc and show values of $EW$(Li~{{\\sc i}}) (Fig.~\\ref{fig:li_MGs}) and a position in the $V-I$ CMD (Fig.~\\ref{fig:VI}) compatible with an age of 150~--~300~Myr; the members of AB~Dor MG are situated at a mean distance of $\\sim$~30~pc and show lithium abundances typical of stars with 30~--~50~Myr (Fig.~\\ref{fig:li_MGs}), which is in agreement with their position in the $M_{\\rm V}$ vs. $V-I$ diagram (Fig.~\\ref{fig:VI}); finally, a set of stars with $EW$(Li~{{\\sc i}}) and positions in the CMD compatible with an age of 80~--~120~Myr are mixed with Hercules-Lyra and AB~Dor MG, and have been classified as other members of the Local Association B4 subgroup (see $\\S$~3). Note that the age estimated using the position of the members of Hercules-Lyra in the CMD is a lower limit since the 80~Myrs isochrone is overlapped with the ZAMS for spectral types earlier than about K5. On the other hand, the age estimated using the equivalent width of the Li~{\\sc i} line $\\lambda$6707.8 \\AA \\ is more robust since the 50\\% of the stars classified as members have measurements of the $EW$(Li~{{\\sc i}}): the Li indicator is useful only for spectral type later than G0, but only three of the 25 candidates of the initial sample are F stars.  Stars in these three subgroups form an excellent list of young cool stars for studying how planets are formed, since they cover a range of ages between 30 and 200 Myr, characteristic of the period during which the Solar System was formed, and they are close enough to be accessible to adaptive optics. In addition, they can be taken as targets for direct imaging detection of sub-stellar companions ---~brown dwarfs and extra-solar giant planets~--- \\citep{neu00,mar03,mas05,low05} and for cold dust, debris disks \\citep{gai04,met04,liu04,che05}."
    },
    "0601/astro-ph0601090_arXiv.txt": {
        "abstract": "Conditional luminosity function (CLF) of dark matter halos supersedes simple models on the number of galaxies as a function of the halo mass, or the so-called halo occupation number, by assigning a luminosity distribution to those galaxies. Suggestions have now been made that both the clustering strength and properties of hosted galaxies of a given dark matter halo depends also on the age and the environment of that halo in addition to halo mass. Based on simple CLF models, we find that at least one suggestion that made use of central galaxy properties may be affected by uncertainties in the group catalog from which clustering properties were measured to address the age dependence. Establishing how galaxy properties  of fixed mass halos change with the environment or the age is challenging, if not impossible, given uncertainties in group and cluster catalogs. It may be possible to address this issue through statistics based on a combined study of luminosity and color distributions, and luminosity- and color-dependent galaxy clustering, all of them as a function of the galaxy overdensity, though no single statistic is likely to provide the ultimate answer. The suggestion that the age dependence of the halo bias invalidates the analytical models of the galaxy distribution is premature given the improvements associated with modeling approaches based on conditional functions. ",
        "introduction": "Recently, several claims have been made that the galaxy properties of a given dark matter halo may depend on the environment in which the halo resides (Gao et al. 2005; Harker et al. 2005; Yang et al. 2005b). A connection is generally made to the age of the halo since age may be the primary factor that determines the environment in which the halo forms. Numerical simulations of the dark matter distribution suggests that older halos are more strongly clustered than younger halos of a given mass (Gao et al. 2005). Such an age dependence on the halo bias is not yet present within successful analytical descriptions of the halo distribution and clustering properties (e.g., Mo \\& White 1996; Sheth, Mo \\& Tormen 2001). These results invalidate analytical models of the galaxy distribution since these models assume that the galaxy properties are only determined by the halo mass (Kravtsov et al. 2004; Cooray \\& Sheth 2002 for a review).  Beyond simulations, whether galaxy properties of a given halo mass vary with the  environment is yet to be established observationally. The galaxy morphology--density (Dressler 1980; Dressler et al. 1997) and the color--density (Blanton et al. 2004) relations do indicate that galaxy properties change with the environment: red, early-type, galaxies are usually found in denser environments than the environments where blue, late-type, galaxies are found. It could, however, be that the increasing fraction of early-type galaxies with increasing density is simply a reflection of a higher fraction of early-type galaxies in more massive halos (e.g., Cooray 2005a; Weinmann et al. 2005). Thus, interpretations of the density-morphology or density-color relations are generally inconclusive with some suggesting support for the halo model (Blanton et al. 2004) while others not (Goto et al. 2003; Balogh et al. 2004).  To study the difference between galaxy properties of same mass halos when either the environment or the age is varied one must find fixed mass halos while a second variable, such as the environment, is varied. The test related to the environment could be achieved with large catalogs of galaxy groups, but since one may not be able to reliably assign a mass to an individual group, it is unlikely that there is an easy way to uniquely identify whether one is comparing halos of similar mass. Since the expected differences are likely to be at the few to, at most, ten percent level (see e.g., discussion in Zheng \\& Weinberg 2005), the mass assignments must be accurate to the same precision. Statistics that measure properties after randomizing galaxies can partly address the question on the environmental dependence by testing whether the surrounding large-scale environment is responsible for properties within a halo. It could be that the internal structure of a halo, such as the dark matter concentration parameter, changes with the environment and galaxy properties reflect those changes.  The second approach to address this question involves statistical study of a large sample of galaxies, such as clustering, binned in terms of an estimator of age for galaxies. Using a catalog of galaxy groups and clusters from the 2-degree Field Galaxy Redshift Survey  (2dFGRS; Colless et al. 2001), Yang et al. (2005b) recently considered the cross-correlation function between the central bright galaxies in galaxy groups and the general population of galaxies above a certain luminosity.  For groups with halo masses above $M_\\star$, they found the groups with a lower luminosity central galaxy to be more strongly clustered than the ones with a higher luminosity galaxy. Furthermore, galaxies that are more passive in a given mass bin were found to be more correlated. Thus, a claim was made that the cross-clustering measurement in Yang et al. (2005b) indicates age and environmental dependence of halo bias.  Using models of the conditional luminosity function (CLF; Yang et al. 2003, 2005a; Cooray \\& Milosavljevi\\'c 2005b), we make a simple model of the cross-correlation between central galaxies of groups and the galaxy distribution and suggest that the measurement in Yang et al. (2005b) may be contaminated by uncertainties in their group catalog involving mass assignments. We find that, for example,  if one selects roughly 10\\% to 20\\% fraction of satellite galaxies of larger mass halos mistakenly as the central galaxy in lower mass halos, then the bias factor for smaller halos with fainter luminosities is boosted.  With a simple relation between the starformation rate and the galaxy luminosity, the clustering results binned in terms of starformation rate may also be explained without requiring a second parameter. It is unlikely that Yang et al. (2005b) measurements provide a definitive answer on the age dependence of galaxy properties of fixed halo mass with age captured by the either the star-formation rate or the luminosity of the central galaxy.  Here, we also consider if there is an appropriate statistical test on the density and age dependence of galaxy properties within a given halo that is affected less by systematics. While we have yet to come up with a simple test to provide the ultimate answer a combination of measurements on the luminosity and color distributions, combined with clustering, all as a function of galaxy density may be utilized to address this issue. Any single statistical measurement is affected by two potential variations: the mass function that varies the halo distribution with the environment and galaxy properties within a fixed halo mass that change with environment: \\begin{equation} P(L,\\delta_{\\rm gal})  \\propto \\int dM\\; \\frac{dn(\\delta_{\\rm gal})}{dM} \\Phi(L|M,\\delta_{\\rm gal}) \\end{equation} where both the mass function,  $dn/dM$, and the CLF, $\\Phi(L|M)$, varies as a function of the environment captured by $\\delta_{\\rm gal}$, which is defined in terms of the density of galaxies in some fixed volume.  To address the question whether $\\Phi(L|M,\\delta_{\\rm gal})=\\Phi(L|M)$ observationally, the statistical methods must allow a separation of effects coming from $dn/dM(\\delta_{\\rm gal})$ with $\\Phi(L|M,\\delta_{\\rm gal})$ (see, e.g., Mo et al. 2004; Cooray 2005a). This degeneracy is also what affects interpretation of the density-morphology or density-color relations.  The {\\it Letter} is organized as following: In the next section, we briefly summarize a model for the cross-clustering between central galaxies in galaxy groups and the general galaxy distribution. We compare our model predictions with those of Yang et al. (2005b) and suggest that their observations can be explained in terms of a small contamination in the group catalog. Furthermore, we comment on potential tests to address whether  $\\Phi(L|M,\\delta_{\\rm gal})=\\Phi(L|M)$ or not. Finally, as a side note, even if it turns out that the CLF has a second parameter, we will argue that this does not imply that the analytical halo model is wrong, but rather it can be easily improved by including the environmental   dependence when describing galaxy properties through conditional functions.  \\begin{figure*}[!t] \\centerline{\\psfig{file=bosch.eps,width=3.6in,angle=0} \\hspace{0.1cm} \\psfig{file=biaseta.eps,width=3.6in,angle=0}} \\vspace{0.3cm} \\caption{The product of bias factor  of central galaxies in groups $b_c(L_c)$ and galaxy bias for a volume limited sample above a $b_J$ magnitude of $-19.45+5\\log h$ with measurements subdivided to mass bins in galaxy groups that host the central galaxy. The measurements are from Yang et al. (2005b) and we show their data for two mass bins from  10$^{12}$ to $10^{12.5}$ (open symbols) and $10^{13.5}$ to 10$^{14}$ M$_{\\sun}$ (filled symbols). In the left panel we consider measurements as a function of central galaxy luminosity while in the right panel we consider their measurements as a function of star-formation rate. The solid and long-dashed lines show the expected bias factor for central galaxies in these mass bins as a function of the central galaxy luminosity. The dot-dashed line shows a simple modification to the bias factor calculation where we include a small correction to central galaxy sample by allowing satellites in larger mass halos ($> 10^{14}$ M$_{\\sun}$) to be identified (wrongly) as central galaxies  and be assigned to a group with a mass in the range between $10^{13.5}$ M$_{\\sun}$ to 10$^{14}$ M$_{\\sun}$. The contaminant fraction $f$ (see text) needed to explain left and right panels differ by a factor of 2 and could indicate extra scatter in mapping between luminosity, which is the primary ingredient in the CLF-based models, to the star-formation rate of central galaxies. The simple analytical description, which does not make any modifications to the halo bias as a function of age, suggests that these measurements with the observed galaxy distribution do not conclusively establish an age dependence to galaxy clustering properties.} \\label{fig:sloan} \\end{figure*} ",
        "conclusions": ""
    },
    "0601/astro-ph0601059_arXiv.txt": {
        "abstract": "We discuss the instrumental and data reduction techniques used to suppress speckle noise with the Simultaneous Differential Imager (SDI) implemented at the VLT and the MMT.  SDI uses a double Wollaston prism and a quad filter to take 4 identical images simultaneously at 3 wavelengths surrounding the 1.62 $\\mu$m methane bandhead found in the spectrum of cool brown dwarfs and gas giants.  By performing a difference of images in these filters, speckle noise from the primary can be significantly attenuated, resulting in photon noise limited data past 0.5''.  Non-trivial data reduction tools are necessary to pipeline the simultaneous differential imaging. Here we discuss a custom algorithm implemented in IDL to perform this reduction. The script performs basic data reduction tasks but also precisely aligns images taken in each of the filters using a custom shift and subtract routine. In our survey of nearby young stars at the VLT and MMT (see Biller et al., this conference), we achieved H band contrasts $>$25000 (5$\\sigma$ $\\Delta$F1(1.575 $\\mu$m) $>$ 10.0 mag, $\\Delta$H$>$11.5 mag for a T6 spectral type object) at a separation of 0.5\" from the primary star. We believe that our SDI images are among the highest contrast astronomical images ever made from ground or space for methane rich companions. ",
        "introduction": "In theory, adaptive optics (AO) systems that are ``photon noise limited'' can detect an object up to 10$^{5-6}$ times fainter its primary at separations of $\\sim$1'' -- sufficient to detect a young, self-luminous giant extrasolar planet.  However, overcoming the large contrast difference between star and planet is not the only obstacle in directly detecting extrasolar planets -- all AO also systems suffer from a limiting ``speckle noise'' floor (Racine et al. 1999).  Within 1'' of the primary star, the field is filled with speckles left over from instrumental features and residual atmospheric turbulence after adaptive optics correction.  These speckles vary as a function of time and color. For photon noise limited data, the signal to noise S/N increases as t$^{0.5}$, where t is the exposure time.  For speckle-noise limited data, the S/N does not increase with time past a specific speckle-noise floor (limiting contrasts to $\\sim$10$^3$ at 0.5'').  This speckle-noise floor is considerably above the photon noise limit and makes planet detection very difficult.  Interestingly, space telescopes such as HST also suffer from a somewhat similar limiting speckle-noise floor due to imperfect optics and ``breathing'' (Schneider et al. 2003).  Direct detection of extrasolar giant planets requires special instrumentation to suppress this speckle noise floor and produce photon noise limited images.  Simultaneous Differential Imaging is an instrumental method which can be used to calibrate and remove the ``speckle noise'' in AO images, while also isolating the planetary light from the starlight.  This method was pioneered by Racine et al. (1999), Marois et al. (2000), Marois et al. (2002) and Marois et al. (2005). It exploits the fact that all cool (T$_{eff}$ $<$1200 K) extra-solar giant planets are thought to have strong CH$_4$ (methane) absorption redwards of 1.62 $\\mu$m in the H band infrared atmospheric window (Burrows et al. 2001, Burrows et al. 2003).  Our SDI device obtains four images of a star simultaneously through three slightly different narrowband filters (sampling both inside and outside of the CH$_4$ features).  These images are then differenced.  This subtracts out the halo and speckles from the bright star to reveal any massive extrasolar planets orbiting that star. Since a massive planetary companion will be brightest in the F1(1.575 $\\mu$m) filter and absorbed in the rest, while the star is bright in all three, a difference can be chosen which subtracts out the star's light and reveals the light from the companion.  Thus, SDI also helps eliminate the large contrast difference between the star and substellar companions.  \\begin{figure} \\begin{center} \\begin{tabular}{cc} \\includegraphics[height=3cm]{fig1a_xv.ps} & \\includegraphics[height=3cm]{fig1b_xv.ps} \\\\ \\end{tabular} \\end{center} \\caption[Raw VLT SDI data] { \\label{fig:SDIRAW} Left: Two minutes of raw SDI data from NACO SDI's 1024$\\times$1024 Aladdin array in the CONICA AO camera (Lenzen et al. 2004).  Right: Same dataset, slightly processed.  Apertures have been selected around each filter image.  In order to reveal the speckle pattern, a heavily smoothed image was subtracted from the raw images (unsharp masking). The resulting speckle patterns are very similar between the 4 simultaneous images which means that an effective subtraction of speckles can be obtained between the filters. } \\end{figure} ",
        "conclusions": ""
    },
    "0601/astro-ph0601509_arXiv.txt": {
        "abstract": "We have obtained high angular resolution, high signal-to-noise spectra of the Calcium triplet absorption lines on the photometric axes of the stellar spheroid in the polar disk galaxy NGC~4650A.  Along the major axis, the observed rotation and velocity dispersion measurements show the presence of a kinematically decoupled nucleus, and a flat velocity dispersion profile. The minor axis kinematics is determined for the first time: along this direction some rotation is measured, and the velocity dispersion is nearly constant and slightly increases at larger distances from the center. The new high resolution kinematic data suggest that the stellar component in NGC~4650A resembles a nearly-exponential oblate spheroid supported by rotation.  The main implications of these results on the previous mass models for NGC~4650A are discussed. Moreover, the new kinematic data set constraints on current models for the formation scenarios of Polar Ring Galaxies (PRGs), supporting a slow accretion rather then a secondary strong dissipative event. ",
        "introduction": "Studies of galaxy formation and evolution both in the Local Group and in the high-redshift universe (Conselice et al. 2003) provide increasing evidence that mergers play a role in the formation of early-type galaxies, in clusters and in the field (Schweizer 1999, Hibbard \\& Yun 1999). The merging of two disk galaxies produces a spheroidal remnant with physical properties, such as density profiles, mean velocity dispersion and surface brightness, similar to those observed for elliptical and early-type disk galaxies (Toomre \\& Toomre 1972; Gerhard 1981; Barnes \\& Hernquist 1992; Bekki 1998a; Naab \\& Burkert 2003; Bournaud et al. 2005). These kind of mergers may also produce a PRG (Bekki 1998b; Bournaud \\& Combes 2003), like NGC~4650A (Fig.\\ref{fig1}).  PRGs generally contain a central featureless stellar spheroid and an elongated structure, the ``polar ring'', which in projection appears close to the inner spheroid's meridian plane. In all PRGs, the \\HI\\ gas is associated with the polar structure and not with central stellar spheroid; furthermore the stars in the central spheroid and the gas and stars in the polar structure rotate in two nearly perpendicular planes. The studies of PRGs promise to yield detailed information about some of the processes at work during galaxy merging and on the shape of dark matter halos in galaxies (Schweizer, Whitmore \\& Rubin 1983, Sackett \\& Sparke 1990, Sackett et al. 1994, Combes \\& Arnaboldi 1996, Iodice et al. 2003).  NGC~4650A is the prototype for PRGs. Its luminous components, inner spheroid and polar structure, have been studied with optical and near-infrared (NIR) photometry, optical spectroscopy, and in the radio (\\HI\\ 21 cm line and continuum).  The surface brightness profile of the central spheroid is described by an exponential law, with a small exponential nucleus; its integrated optical vs. NIR colors are similar to those of an intermediate age stellar population (Iodice et al. 2002; Gallagher et al. 2002). Previous absorption line spectroscopy at optical wavelengths showed a substantial rotation along the major axis, with $v_{max} \\simeq 100$ \\kms (Whitmore et al. 1987, Sackett et al. 1994), while the velocity dispersion measurements were plagued by systematic errors, caused by both low angular and low spectral resolution. Best estimates gave $\\sigma_0 \\simeq 60$ \\kms.  The polar structure in NGC~4650A has been shown to be a disk, rather than a ring. Its stars and dust can be reliably traced inward within the stellar spheroid to $\\sim 1.2$ kpc radius from the galaxy nucleus (Iodice et al. 2002; Gallagher et al. 2002). Emission and absorption line optical spectra show that it is an extended stellar disk rather than a narrow ring (Swaters \\& Rubin 2003): both rotation curves show a linear inner gradient and a plateau, as expected for a disk in differential rotation. Furthermore the \\HI\\ 21 cm observations (Arnaboldi et al. 1997) show that the gas is 5 times more extended than the luminous polar structure, with a position-velocity diagram very similar to those observed for edge-on gaseous disks.  The question about the shape of the dark halo of NGC~4650A is still open. Because of the large error-bars in the velocity dispersion profile, dynamical models by Sackett et al. (1994) and Combes \\& Arnaboldi (1996), which differ in the orientation of the main axes and flattening of the dark halo, were both compatible with the observed data. The polar disk luminous mass distribution along the meridian plane may also affect the inner stellar kinematics and induce a triaxial-like perturbation to the axisymmetric stellar mass distribution. Kinematic information along independent position angles on the sky may therefore be required for a reliable mass model. Furthermore, moving to longer wavelength may help reducing the contamination from the dust (Arnaboldi et al. 1995).  An 8 meter class telescope, with a spectrograph like FORS2 and a high efficiency holographic grism in the 1 micron wavelength range shall allow us to study the stellar motions of the NGC~4650A inner spheroid, via the absorption line spectroscopy of the Calcium triplet (CaT) lines, at 8498 \\AA, 8542 \\AA, 8662 \\AA.  These lines are the strongest features in the stellar continuum for a large variety of stellar types (Cenarro et al. 2001).  In this work, we shall present our measurements for the radial velocity, the velocity dispersion profile and the Gaussian Hermite parameters $H_3,\\, H_4$ along the main photometric axes of the central spheroid in NGC~4650A.  We then discuss their implication on the proposed mass models and formation scenario for PRGs. In what follow, we shall adopt 2905 \\kms as the heliocentric systemic velocity of NGC~4650A (Arnaboldi et al. 1997) and H$_0 = 70$\\kms Mpc$^{-1}$, which implies $1''= 201$ pc. ",
        "conclusions": "We have obtained new high angular resolution, high signal-to-noise spectra along the photometric axes of the stellar spheroid in NGC~4650A. Compared to previous kinematics studies, the new observations show {\\it i)} a kinematic signature of a decoupled core, and {\\it ii)} a non zero rotational velocity and an increasing velocity dispersion along the minor axis.  In the following we discuss them and address the main conclusions about the structure and formation of NGC~4650A.  \\subsection{Comparison with previous observations and dynamical models}  We now discuss the current rotation velocities and velocity dispersion measurements with those previously published in the literature (Whitmore et al. 1987, Sackett et al. 1994, Swaters \\& Rubin 2003). Our data represent a significant improvement in angular resolution, which is needed for the understanding of the nuclear regions, and no velocity dispersion measurements were available for the spheroid minor axis prior to this work.  The rotation curve along the polar disk at PA$=158^\\circ$ was measured from the H$\\alpha$ emission line (Whitmore et al. 1987), and from optical absorption lines at PA$=155^\\circ$ (Swaters \\& Rubin 2003), with slits aligned separately along North and South part of the polar disk, so not through the center of the spheroid; both were obtained at a slightly different PA from the one of the minor axis spheroid (PA$=152^\\circ$; Iodice et al. 2002).  A direct comparison of the major axis kinematics data, $v$ and $\\sigma$ is shown in Fig.~\\ref{conf}; previous published data show a larger scatter for the velocity dispersion measurements, while the agreement is good for the rotation curves. Our kinematic data are consistent within the errors: the new rotation curve is in good agreement with the previous ones and our $\\sigma$ profile has clearly benefited from a substantial increase in S/N, thanks to the collecting area of an 8 meter telescope. The new $\\sigma$ measurements are in agreement with those by Whitmore et al. (1987), while they are systematically larger than the measurements from Sackett et al. (1994).  The average difference in this case is about $15$ \\kms, similar to the 19 \\kms, which Sackett et al. (1994) subtracted off the data to account for their systematics.  The new high resolution CaT spectra are better tracers of the kinematics for the NGC~4650A spheroid than was available before. The higher angular resolution and the net increase in S/N allow us to 1) measure a {\\it flat} velocity dispersion profile along the spheroid major axis, while previous $\\sigma$ measurements were too scattered to reliably establish any trend with radius, 2) detect a decrease of the velocity dispersion in the center, and 3) hint of increasing $\\sigma$ in the NE, as previously noticed by Sackett et al. (1994).  The measured flat velocity dispersion profile along the spheroid major axis shows that both the linear decreasing fit $\\sigma_r(r) = \\sigma_r(0) - K\\times r_d $ proposed by Sackett et al. (1994) and the exponential empirical law $ \\sigma_r(r) = \\sigma_r(0)\\exp[-(r/4r_d)^4]$ proposed by Combes \\& Arnaboldi (1996) do not reproduce the observed trend with radius. Previous conclusions that the same authors drew were based on data that are no longer valid, and the dynamical model for NGC 4650A must be revised.  \\subsection{The spheroid kinematics}\\label{sphkin}  {\\it The stellar spheroid $v/\\sigma - \\epsilon$ relation} - The light distribution of the central stellar spheroid in NGC~4650A is fit by a nearly exponential surface brightness profile, with $r_e = 6''.7$ and $\\epsilon = 0.46$ in the I band, and nearly edge-on as inferred from the presence of a thin disk aligned with the spheroid's major axis (Iodice et al. 2002). This is a low luminosity spheroid, with $M_B \\sim -18.1$ (Gallagher et al. 2002) based on HST photometry.  The observed kinematics along the major axis is consistent with that of a spheroidal galaxy supported by rotation: at $r\\geq 2r_e$, where the ellipticity $\\epsilon \\sim 0.5$, we estimate $(v/\\sigma)^* \\sim 0.96$.  The anisotropy parameter $(v/\\sigma)^*$ is defined as the ratio of the observed value of $(v/\\sigma)$ and the theoretical value for an isotropic oblate rotator $(v/\\sigma)_{oblate}=[\\epsilon/(1-\\epsilon)^{1/2}]$ (Binney 1978). For a total B magnitude of $M_B \\sim -18.1$ (Gallagher et al. 2002), in the plane $\\log(v/\\sigma)^* - M_B$ the central spheroid of NGC~4650A is located in the region of low luminosity elliptical galaxies flattened by rotation (see Fig.18 in Bender et al. 1994).  {\\it Faber-Jackson relation and the Fundamental plane } - We have investigated the position of the stellar spheroid in NGC~4650A with respect to the {\\it Faber-Jackson relation (FJR)} (Faber \\& Jackson 1976) and {\\it Fundamental Plane (FP)} (Djorgovski \\& Davies 1987) of elliptical galaxies. In Fig.\\ref{FJFP} (top panel) we compare the absolute B magnitude and the central velocity dispersion for the spheroid in NGC~4650A with those for ellipticals in the Virgo and Coma clusters (in the sample studied by Dressler et al. 1987).  For the observed velocity dispersion, the spheroidal component in NGC~4650A is more luminous with respect to the predicted value by the best-fit of FJR (which did not include the 5 galaxies with $log(\\sigma)<2$).  Respect to the B-band\\footnote{The effective radius and the effective surface brightness used to derive the FP for the spheroid in NGC~4650A were estimated by fitting the whole HST B-band light profiles with a Sersic law ($I(r)\\propto r^{1/n}$), where $n \\sim 2$, $\\mu_e \\sim 23.00 \\pm 0.02$ and $r_e \\sim 7.1 \\pm 0.1$ arcsec.}  FP of elliptical galaxies and bulges (by Bender et al. 1992; Jorgensen et al. 1996; Falcon-Barroso et al. 2002), the spheroid in NGC~4650A falls slightly above the average relation for Es and S0 in the Coma cluster (see bottom panel of Fig.\\ref{FJFP}). Forbes \\& Ponman (1999) showed that bluer early-type galaxies have larger deviations from the FJR: the observed deviations of the NGC~4650A spheroid from both the FJR and FP, toward higher luminosity, is consistent with the stellar component being younger and bluer (Gallagher et al. 2002; Iodice et al. 2002) than the one in standard early-type galaxies.  {\\it Kinematics on the spheroid minor axis} - The new kinematics observed along the minor axis is rather complex and needs to be discussed in some details.  In the inner regions ($r\\le 4''-6''$), where the spheroid's light dominates, the observed kinematics is consistent with an isotropic oblate spheroid, as suggested by the major axis kinematics. The LOS velocities are close to zero and the velocity dispersions measurements are very similar to the constant values measured along the major axis slit. Within this regions, the slightly asymmetric circular velocity and dispersion profiles, with respect to the galaxy center, may reflect the contamination by the ``inner polar disk arms'' (Fig.\\ref{fig2}). On the SE side respect to the galaxy center, along the minor axis slit, the light includes most of the inner polar disk arms (Gallagher et al. 2002, Iodice et al. 2002) and both the LOS velocity and the velocity dispersion at $r \\ge 0.5''$ are larger than those observed on the NW side (Fig.\\ref{RCmin}).  At larger radii (for $r \\ge 6''$), the dispersion increases and larger rotational velocities are observed. Three possible interpretations are currently viable: {\\it i)} the observed profiles are tracing the intrinsic spheroid kinematics; {\\it ii)} these effects are produced by contamination of the stellar motion in the polar disk or {\\it iii)} by stronger Paschen lines in a mixed population (Sec.\\ref{Paschen}).  If the stellar spheroid has indeed intrinsic minor axis rotation (Fig.\\ref{RCmin}, with $\\mu \\sim 0.23$\\footnote{According to Binney (1985), the amount of minor axis rotation is parametrized as $\\mu = v_{min}/\\sqrt{(v_{maj}^{2} +v_{min}^{2})}$.} at $r\\sim 0.5 r_e$), and the presence of isophotal twist (Fig.\\ref{fig2}), both suggest that the system may be triaxial (Wagner, Bender \\& Moellenhoff 1988; Franx, Illingworth \\& de Zeeuw 1991). If the flat or slightly increasing velocity dispersion profile at large radii is intrinsic, it may trace the presence of larger masses along the meridian plane.  This would be consistent with the presence of a massive polar disk (Arnaboldi et al. 1997; Iodice et al. 2002) and with the empirical evidences for a dark halo flattened along this direction from the Tully-Fisher relation of PRGs (Iodice et al. 2003).  The potential generated by a bar can also induce minor axis rotation (Athanassoula 1992), but the presence of a bar in the stellar spheroid of NGC 4650A is not supported by the photometry.  \\subsection{The nuclear regions of NGC~4650A}\\label{nuc}  The stellar spheroid of NGC~4650A is known to contain a very luminous, compact nucleus, whose extension is about 20 pc ($\\sim 0.1''$) and $L_I \\sim 8 \\times 10^7 L_{\\odot}$, which is superimposed on a more extended and rounder structure, of about 60 pc (Gallagher et al. 2002), see also Sec.~\\ref{result}. The high angular resolution of our spectra allows us to reveal the kinematic signature of the nucleus: within a $0''.5$ radius along the major axis, we observe a decoupled component in the circular velocity profile, associated with a lower velocity dispersion.  Such features indicate the presence of a small core, kinematically distinct from the main central spheroid, whose extension ($\\sim 100$ pc) is consistent with the photometry. This is also comparable with the sizes of the kinematically decoupled cores (KDC) observed in ``normal'' early-type galaxies, which varies from some hundred kpc (Kropolin \\& Zeilinger 2000) to about $30 - 60$ pc, for the very small cores detected in HST images (Krajnovi\\'{c} \\& Jaffe 2004). The HST data for KDC in early-type galaxies show that the isophotes in these regions are rounder and have a different PA with respect to the outer regions (Krajnovi\\'{c} \\& Jaffe 2004): this behavior is very similar to what is observed in the central $0''.5$ radius in NGC~4650A (Fig.~\\ref{fig2}).  The observed KDC kinematics in early-type galaxies has very similar features to those measured in the nuclear regions of NGC~4650A; on the whole, the velocity profile is characterized by a central asymmetry, with a flat or a decreasing velocity dispersion profile at the correspondent nuclear radii.  An upper limit to the mass of the KDC observed in NGC~4650A is $\\sim 10^8 M_{\\odot}$: here we have assumed that the circular-orbit speed at $0.5''$ from the center is less than the squared sum of the velocity dispersion and the measured circular velocity. The nuclear mass estimate for NGC~4650A is higher than the masses of the star cluster nuclei derived for late-type spirals, which varies from $10^5$ to $10^7 M_{\\odot}$ (Matthews \\& Gallagher 2002; B$\\ddot{o}$ker et al. 2001): as already suggested by Gallagher et al. (2002), this is consistent with a more luminous star cluster nucleus in NGC~4650A, for which we estimate a mass-to-light ratio $M/L \\sim 1.5 M_{\\odot}/L_{\\odot}$.  The moderate age for the star cluster nucleus in NGC~4650A (Gallagher et al. 2002), based on the $B - I$ color, is consistent with the absence of Paschen lines in the nuclear regions (Sec.~\\ref{Paschen}).  KDC are often related to secondary event in a galaxy evolution (Bertola \\& Corsini 1999). Numerical simulations indicate that kinematically peculiar nuclei may result from the merging of two disk galaxies, where the size of the KDC becomes smaller (less than the effective radius) as the mass ratio of the two merging galaxies is smaller than one (Balcells \\& Gonzalez 1998). KDC are also present in barred galaxies, and are characterized by a low velocity dispersion in the center (Emsellem et al. 2001). In this case, they are the result of a secular evolution, rather than of an interaction, because the gravitational torques by the bar drives the gas toward the central regions and this then forms new stars with lower velocity dispersion (Wozniak et al. 2003), but in NGC 4650A, no evidence for a bar-like structure was detected (Gallagher et al. 2002; Iodice et al. 2002), and the ellipticity and PA profiles do not show the typical features observed for barred galaxies (Wozniak et al. 1995; Erwin \\& Sparke 1999).  The formation scenarios for PRGs, either via major merging of two disk galaxies (Bekki 1998b) or accretion of external material (Bournaud \\& Combes 2003), predict that some gas flows to the galaxy center and quickly forms a small star cluster. This can easily lead to the formation of a KDC in these systems.  \\subsection{Implications for the formation of NGC~4650A}  As discussed in Sec.~\\ref{sphkin}, the observed kinematics and the photometry suggest that the central galaxy in NGC~4650A is a spheroid with a nearly exponential surface brightness profile, supported by rotation.  In a merging scenario with low relative velocity and low impact parameter (Bekki 1998b, Bournaud \\& Combes 2003), a high mass ratio is required to form a massive polar disk, as observed in NGC~4650A: the 'victim' mass should exceed the 'intruder' mass. According to simulations of galaxy mergers (e.g. Bournaud et al. 2005) this would convert the intruder into an elliptical-like, not rotationally supported, stellar system. Differently, in the tidal accretion scenario, with large relative velocity and large impact parameter, for a particular orbital configuration and a gas-rich donor, a polar ring and/or disk may form both around a disk or an elliptical galaxy (Bournaud \\& Combes 2003). External gas can be also accreted from the cosmic web filaments (Dav\\'e et al. 2001; Maccio et al. 2005): in this formation scenario, there is no limits to the mass of the accreted material, thus a very massive polar disk may develop either around a stellar disk or a spheroid.  The new kinematic data obtained for NGC~4650A suggest that the accretion scenarios are favored, nevertheless more investigations are needed to solve the discrepancies between the observations and the theoretical predictions."
    },
    "0601/astro-ph0601023_arXiv.txt": {
        "abstract": "The majority of nearby early-type galaxies contains detectable amounts of emission-line gas at their centers. The nuclear gas kinematics form a valuable diagnostic of the central black hole (BH) mass. Here we analyze and model HST/STIS observations of a sample of 27 galaxies; 16 Fanaroff \\& Riley Type I radio galaxies and 11 (more) normal early-type galaxies. We focus here on what can be learned from the nuclear velocity dispersion (line width) of the gas as a complement to the many studies dealing with gas rotation velocities. We find that the dispersion in a STIS aperture of $\\sim 0.1''$--$0.2''$ generally exceeds the large-scale stellar velocity dispersion of the galaxy. This is qualitatively consistent with the presence of central BHs, but raises the questions whether the excess gas dispersion is of gravitational or non-gravitational origin and whether the implied BH masses are consistent with our current understanding of BH demography (as predicted by the $M-\\sigma$ relation between BH mass and stellar velocity dispersion). To address this we construct purely gravitational axisymmetric dynamical models for the gas, both thin disk models and models with more general axis ratios and velocity anisotropies. For the normal galaxies the nuclear gas dispersions are adequately reproduced assuming disks around BHs with masses that follow the $M-\\sigma$ relation. In contrast, the gas dispersions observed for the radio galaxies generally exceed those predicted by any of the models. We attribute this to the presence of non-gravitational motions in the gas that are similar to or larger than the gravitational motions. The non-gravitational motions are presumably driven by the active galactic nucleus (AGN), but we do not find a relation between the radiative output of the AGN and the non-gravitational dispersion. Given the uncertainties about the dynamical state of the gas, it is not possible to uniquely determine the BH mass for each galaxy from its nuclear gas dispersion. However, for the sample as a whole the observed dispersions do not provide evidence for significant deviations from the $M-\\sigma$ relation. ",
        "introduction": "\\label{s:intro}  The mass of a central black hole (BH) in a galaxy can be directly weighed using gravitational 'test particles' moving around it (see e.g., Kormendy \\& Richstone 1995). Commonly used test particles are stars, optical emission-line gas and maser clouds. These move around in the combined potential well of the stellar mass and the central black hole mass.  A secure direct dynamical measurement of the BH mass requires that the distribution and kinematics of the test particles are measured in the immediate vicinity of the BH, i.e., the 'black hole's sphere of influence', so that its gravitational potential has a measurable effect on the kinematics in addition to the effect of the stellar mass potential. Ground-based optical telescopes can resolve the BH's sphere of influence only for very nearby galaxies, with our Milky Way being a spectacular nearby example (e.g., Ghez \\etal 2003; Sch{\\\"o}del \\etal 2003). The Hubble Space Telescope (HST) can provide the required spatial resolution for galaxies with distances up to several tens of Mpc (e.g., Kormendy \\& Gebhardt 2001). The Very Long Baseline Interferometer (VLBI) can in principle probe BH masses out to even larger distances provided that the galaxies display nuclear maser emission. Together these methods have led to direct dynamical measurements of BH masses $\\Mbh$ in the range $\\Mbh \\sim 10^6 - 5 \\times 10^9 \\Msun$ in several tens of galaxies in the nearby Universe (see e.g., Tremaine \\etal 2002; Marconi \\etal 2003 for listings of BH mass measurements). The BH masses are mostly based on dynamical modeling of kinematics of either stars or emission-line gas. The detected black hole masses $\\Mbh$ correlate well with both host spheroid luminosity $L_B$, global stellar velocity dispersion $\\sigma_{\\rm s}$ and spheroid mass (Gebhardt et al., 2000; Ferrarese \\& Merritt, 2000; Tremaine \\etal 2002; Marconi \\etal 2003; H\\\"aring \\& Rix 2004). These correlations roughly represent our current knowledge of the local BH demography.  To explore these correlations further, i.e., their scatter and extent in BH mass, dynamical modeling using emission-line gas kinematics in early-type galaxies is convenient and in some respects even crucial. First of all, more than 50\\% of the early-type galaxies in the nearby Universe contain gas at their centers (e.g., Goudfrooij \\etal 1994; Ho, Filippenko \\& Sargent 1997). Second of all, the most massive black holes currently found (i.e., $\\Mbh \\gta 10^9 \\Msun$) are often in giant ellipticals. The central stellar surface brightness of these galaxies is typically too low to obtain accurate stellar kinematics with HST and maser emission has not been detected (e.g., Barth \\etal 2004).  At the centers of galaxies, the collisional gas is expected to settle quickly into a disk, if unperturbed by forces other than gravity (e.g., Habe \\& Ikeuchi 1985). The frequent detection with HST of disk-like structures in dust and gas in early-type galaxies suggests that settling actually takes place (e.g., van Dokkum \\& Franx 1995; Tran \\etal 2001; Verdoes Kleijn \\etal 1999; Laine \\etal 2003). Therefore the extended gas velocities in galactic nuclei have been modeled assuming a thin disk in circular rotation. Such models have been successful in explaining the observed gas velocities.  However, in the few cases where both gas velocities and gas dispersions are modeled, it is often found that the nuclear velocity dispersion exceeds the prediction for a thin disk (van der Marel \\& van den Bosch 1998; Barth \\etal 2001; Maciejewski \\& Binney 2001; Cappellari et al.~2002; Verdoes Kleijn et al.~2002). Many of these targets are AGN and display regular disks of dust and gas.  The origin of this 'excess' velocity dispersion is unknown. Two common ad-hoc assumptions are that: (a) the excess dispersion is due to non-gravitational 'turbulent' forces of unknown origin which do not affect the mean circular rotation of the gas (e.g., van der Marel \\& van den Bosch 1998; Verdoes Kleijn \\etal 2002); or (b) the excess dispersion is purely gravitational and affects the rotation of the gas through the asymmetric drift equation (e.g., Barth \\etal 2001).  A main aim of this paper is to determine which early-type galaxies do and which do not display an excess in nuclear gas dispersion.  If there are galaxies for which the disk remains thin all the way to the nucleus, then any velocity dispersion observed through a nuclear aperture is caused by differential rotation over this aperture. A nuclear aperture samples gas which is much closer to the BH than a series of apertures which samples the extended rotation curve. Hence, the nuclear gas dispersion can be sensitive to much lower black hole masses than the rotation curve. This is relevant as the sensitivity in terms of the minimum detectable BH mass in early-type galaxies with stellar dynamical methods or with gas rotation velocities is typically not much below the values predicted by the correlations between BH mass and spheroid properties. A relevant example in this case is the claim that active spiral galaxies classified as Narrow Line Seyfert 1s harbor BHs with masses smaller than those predicted by the $\\Mbh-\\sigma$ relation (e.g., Grupe \\& Mathur 2004).  For the galaxies with nuclear gas dispersions in excess of that expected from a thin disk, we want to determine the origin of this excess.  On the one hand, if it has a gravitational origin, this implies that the gas should rotate less fast than the circular speed and the thin-disk approximation does not hold. The gas might then have a distribution somewhere in the range from a thin disk to a purely spherical distribution. Previous BH mass estimates based on the assumption of thin disks would then be underestimates, and this would affect the observed correlations between black hole mass and large-scale spheroid properties.  On the other hand, if the excess dispersion has a non-gravitational origin, we would like to quantify how it depends on galaxy parameters, such as nuclear activity. This in turn could shed light on the accretion process in these nearby, typically low-luminosity, AGN.  In summary, this paper studies the nature of the line widths of emission-gas in nearby early-type galaxy nuclei and its implications for the BH demography and BH accretion processes. A similar study into the line-widths of gas for late-type galaxies has been performed by Sarzi \\etal (2002).  The outline of the paper is as follows. Section~\\ref{s:sampledata} describes the sample selection, the data and the determination of the central gas velocity dispersions. Section~\\ref{s:excess} compares the observed dispersions to those expected from a thin circular disk. Section~\\ref{s:spheroid} compares the observed dispersions to those expected from a gas distribution which is more spheroidal than a disk. Section~\\ref{s:nongrav} constrains the direction of the non-gravitational dispersion. Section~\\ref{s:cave} describes the caveats of our modeling. Finally, Section~\\ref{s:conc} provides a discussion of the results and lists the main conclusions of the paper. Throughout the paper we use a Hubble constant $H_0$= 75 kms$^{-1}$Mpc$^{-1}$. ",
        "conclusions": "  \\begin{itemize}  \\item{For the normal galaxies in the sample (i.e., without large-scale radio-jets) the nuclear gas dispersions are adequately reproduced by models that have purely gravitational motion and BHs that follow the $\\Mbh-\\sigma$ relation. Among the purely gravitational models we cannot generally discriminate between thin disk models and vertically extended models. The former might seem preferred theoretically because they represent a longer-term equilibrium configuration for the gas. However, in some cases it is unclear if the observed line profiles are consistent with the double-peaked structure that is generally predicted by thin disk models.}  \\item{The nuclear gas dispersions observed for the radio galaxies generally exceed those predicted by models with only gravitational motions in either a thin disk or a more spheroidal gas distribution. We attribute this to the presence of non-gravitational motions in the gas that are similar to or larger than the gravitational motions. The non-gravitational dispersion is consistent with being either isotropic or perpendicular to the extended gas and dust disks. The non-gravitational motions are presumably driven by the active galactic nucleus (AGN), but we do not find a relation between the radiative output of the AGN and the non-gravitational dispersion.}  \\item{Given the uncertainties about the dynamical state of the gas, it is not possible to uniquely determine the BH mass for each galaxy from its nuclear gas dispersion. However, for the sample as a whole the observed dispersions do not provide evidence for significant deviations from the $\\Mbh-\\sigma$ relation in either active or non-active elliptical galaxies. In no case is the observed nuclear gas dispersion so low that it puts an upper limit on the BH mass that is significantly below the $\\Mbh-\\sigma$ relation.}  \\end{itemize}  \\noindent For the normal, i.e., non-radio galaxies, one should note that the success of purely gravitational models does not imply that there are no alternative models that can fit the data. For example, it cannot be ruled out that in reality there is a non-gravitational component to the gas motions, and that the BH masses in these galaxies are below the $M-\\sigma$ relation. Some models in the literature, such as those published for IC 1459 (Cappellari \\etal 2002) fall in this category. Nevertheless, if one is willing to assume {\\it a priori} that the BH masses of all elliptical galaxies follow the $\\Mbh-\\sigma$ relation, then one can turn the argument around and conclude that normal elliptical galaxies do not have a very significant non-gravitational component in their gas motions. The absence of excess dispersion in non-radio galaxies is corroborated by the fact that the BH mass estimates for this sample by itself yield a relation in good agreement with the $\\Mbh-\\sigma$ relation (e.g., Figure~\\ref{f:mbhdisk}). It is interesting to note in this context that a similarly good correspondence between BH masses from gas dispersions and from the $\\Mbh-\\sigma$ relation was found for non-active early-type {\\it spiral} galaxies by Sarzi et al.~(2002).  For the low-luminosity radio galaxies in our sample there is also other evidence, besides the dynamical evidence presented here, for non-flat gas distributions with turbulent motion. The core emission-line luminosity of these galaxies correlates with the radio and optical core luminosity which are both most likely due to synchrotron emission from the jet (e.g., Chiaberge \\etal 1999; Verdoes Kleijn \\etal 2002). This correlation could mean that the emission-line luminosity is driven by jet photo-ionization. It then implies covering factors of $\\sim 0.3$ and hence thick disks (Capetti 2005). Alternatively, the correlation could indicate that the gas is excited by shocks induced by jet-gas interactions.  Although it is mostly the radio galaxies that show significant excess dispersion over a purely gravitational model, there is one normal galaxy for which this is the case too: IC 989. It should be noted though that for this particular galaxy the error of $33\\kms$ on the stellar dispersion $\\sigma_{\\rm s} = 176\\kms$ is larger than typical for the sample. The predicted gas dispersion can be increased by a factor $\\sim 4$ by simply varying $\\sigma_{\\rm s}$ within its 1-$\\sigma$ confidence range. Also, the predicted gas dispersion could go up by up to a factor $\\sim 2$ if the unknown inclination of the purported gas disk is closer to edge-on than the default value of $60\\deg$ that we assumed. So the evidence for non-gravitational gas motions in the normal galaxy IC 989 is not strong.  Finally, we note that it is the sample of relatively close to face-on ($i<40\\deg$) gas disks which reveal most clearly an excess dispersion compared to a purely gravitational model. Unfortunately, this face-on group only contains radio galaxies. However, a statistically significant difference between the excess dispersion in radio and non-radio galaxies remains when one excludes galaxies with $i<40\\deg$ gas disks from the analysis. Nevertheless, a good test of the dichotomy found in this paper will be the modeling of gas disks at lower inclinations also in non-radio galaxies. Such data is not available currently.  The evidence for non-gravitational gas motions in radio galaxies makes the gas dispersion unsuited for dynamical mass estimates in this class of galaxies. It would be interesting to establish more generally whether dynamical mass estimates from gas velocities, including rotation curves, are robust. This would require comparisons of BH masses from independent indicators (e.g., stellar and gas kinematics) within the same galaxy. Such a comparison has been performed for a radio galaxy (NGC 4335, Verdoes Kleijn \\etal 2002) and a non-radio galaxy (Cappellari \\etal~2002). Unfortunately, in both cases the stellar dynamical evidence is too inconclusive to reach firm conclusions. But it does appear in both cases that the BH mass implied by thin disk models for the gas rotation curves is lower than would otherwise have been expected, either from stellar kinematics or from the $\\Mbh-\\sigma$ relation. This suggests that the gas might be moving slower than the circular velocity, as would be expected for example if there were asymmetric drift in the gas. This occurs for gas distributions that are dynamically hotter than a thin disk. More independent BH mass evaluations from gas and other tracers, both for normal and active galaxies, would be quite valuable to shed more light on these issues."
    },
    "0601/astro-ph0601679_arXiv.txt": {
        "abstract": "The Milky Way is made up of a central bar, a disk with embedded spiral arms, and a dark matter halo. Observational and theoretical constraints for the characteristic parameters of these components will be presented, with emphasis on the constraints from the dynamics of the Milky Way gas. In particular, the fraction of dark matter inside the solar radius, the location of the main resonances, and the evidence for multiple pattern speeds will be discussed. ",
        "introduction": "In our current understanding of the universe galaxies evolve within dark matter halos during the expansion of the universe.  Very successfully, the theory of hierarchical clustering and merging of dark matter halos explains the structure on large scales and the variety of galaxy masses and types observed. Yet, on smaller scales the details of dissipation, star formation, and secular evolution, which also depend on non-gravitational forces, become more important and need to be addressed.  Observations and detailed modeling is required to help understand the processes of galactic evolution.  Our own galaxy is particularly important in this respect, because it allows much more detailed observations, down to stellar proper motions around the central black hole and planet formation around young stars. Many new large surveys of the visible matter inventory of the Milky Way are under way: the survey RAVE hopes to measure radial velocities and chemical abundances of the brightest 50 million stars, the space interferometry mission SIM will accurately determine the distance to the galactic center and the local standard of rest (LSR) velocity, and GLIMPSE, a Spitzer Legacy Science Program, already built a catalog of 47 million sources in the longitude range $10^\\circ\\leq|l|\\leq 64^\\circ$ within the galactic plane $|b|\\leq 1^\\circ$, to be extended to fill the gap within $|l|\\leq 10^\\circ$. The goal of these surveys is to understand the Milky Way's present state and formation history \\citep{GLIMPSE,SIM,RAVE}.  Detailed modeling of the Milky Way is required to make sense of these data sets. Because of our position within the galactic plane, we cannot grasp the large-scale structure easily.  With certain symmetry assumptions, however, it is possible to recover the present 3-dimensional structure. In the following, we review the steps necessary for deriving such a model. ",
        "conclusions": "In this review, we gave an overview of the results obtained from reconstruction of the present state of the Milky Way using near-IR luminosity and gas kinematics complimentary data sets. From the luminosity data alone, one could not infer anything about the dark matter component. Likewise, the observed gas kinematics alone is difficult to interpret and has lead historically to wrong conclusions, such as steep drops in the inferred rotation curve, large dark matter components, and explosions in the galactic center. We avoid using parametric models for describing the important bulge region which is best described as a barred bulge.  This bar in the galactic center explains the observed asymmetry in the near-IR light distribution within the bulge region as well as the high non-circular gas motion inside 3 kpc galactocentric radius.  The inferred mass density is in good agreement with micro-lensing optical depth observations which provide an important independent test. The molecular ring material, is interpreted as four tightly wound spiral arms in approximate circular rotation outside the corotation radius of the bar. Models indicate that the spiral arm pattern rotates much slower than the bar, allowing the arms to extend beyond the solar orbit. Most imporantly, the model has provided evidence for the maximum disk hypothesis and the amount of dark matter present inside the solar radius.  Future and presently ongoing observational surveys will provide even more detailed data which will allow us to improve the model further. The next revision of the Milky Way models will probably allow more detailed studies of the spiral arms and locate more resonances."
    },
    "0601/gr-qc0601056.txt": {
        "abstract": "We investigate the conditions under which cosmological variations in physical `constants' and scalar fields are detectable on the surface of local gravitationally-bound systems, such as planets, in non-spherically symmetric background spacetimes. The method of matched asymptotic expansions is used to deal with the large range of length scales that appear in the problem. We derive a sufficient condition for the local time variation of the scalar fields driving variations in 'constants' to track their large-scale cosmological variation and show that this is consistent with our earlier conjecture derived from the spherically symmetric problem. We perform our analysis with spacetime backgrounds that are of Szekeres-Szafron type. They are approximately Schwarzschild in some locality and free of gravitational waves everywhere. At large distances, we assume that the spacetime matches smoothly onto a Friedmann background universe. We conclude that, independent of the details of the scalar-field theory describing the varying `constant', the condition for its cosmological variations to be measured locally is almost always satisfied in physically realistic situations. The very small differences expected to be observed between different scales are quantified. This strengthens the proof given in our previous paper that local experiments see global variations by dropping the requirement of exact spherical symmetry. It provides a rigorous justification for using terrestrial experiments and solar system observations to constraint or detect any cosmological time variations in the traditional `constants' of Nature in the case where non-spherical inhomogeneities exist.  PACS Nos: 98.80.Es, 98.80.Bp, 98.80.Cq $\\ $ ",
        "introduction": "Over the past few years there has been a resurgence of observational and theoretical interest in the possibility that some of the fundamental `constants' of Nature might be varying over cosmological timescales \\cite% {webb}. In respect of two such `constants', the fine structure constant, $% \\alpha $, and Newton's `constant' of gravitation, $G$, the idea of such variations is not new, and was proposed by authors such as Milne \\cite{milne}% , Dirac \\cite{dirac}, and Gamow \\cite{gam} as a solution to some perceived cosmological problems of the day \\cite{btip}. At first, theoretical attempts to model such variations in constants were rather crude and equations derived under the assumption that constants like $G$ and $\\alpha $ are true constants were simply altered by writing-in an explicit time variation. This approach was first superseded in the case of varying $G$ by the creation of scalar-tensor theories of gravity \\cite{jordan}, culminating in the standard form of Brans and Dicke \\cite{bd} in which $G$ varies through a dynamical scalar field which conserves energy and momentum and contributes to the curvature of spacetime by a means of a set of generalised gravitational field equations. More recently, such self-consistent descriptions of the spacetime variation of other constants, like $\\alpha $ \\cite{bek, bsm}, the electroweak couplings \\cite{ewk}, and the electron-proton mass ratio, $\\mu $% , \\cite{bm} have been formulated although most observational constraints in the literature are imposed by simply making constants into variables in formulae derived under the assumption that are constant.  The resurgence of interest in possible time variations in $\\alpha $ and $\\mu $ has been brought about by significant progress in high-precision quasar spectroscopy. In addition to quasar spectra, we also have available a growing number of laboratory, geochemical, and astronomical observations with which to constrain any local changes in the values of $\\ $these constants \\cite{reviews}. Studies of the variation of other constants, such as $G$, the electron-proton mass ratio, $\\mu =m_{e}/m_{pr}$, and other standard model couplings, are confronted with an array of other data sources. The central question which this series of papers addresses is how to these disparate observations, made over vastly differing scales, can be combined to give reliable constraints on the allowed global variations of $% \\alpha $ and the other constants. If $\\alpha $ varies on cosmological scales that are gravitationally unbound and participate in the Hubble expansion of the universe, will we see any trace of this variation in a laboratory experiment on Earth? After all, we would not expect to find the expansion of the universe revealed by any local expansion of the Earth. In Paper I \\cite% {shawbarrow1}, we examined this question in detail for spherically symmetric inhomogeneous universes that model the situation of a planet or a galaxy in an expanding Friedmann-Robertson-Walker (FRW)-like universe. In this paper we relax the strong assumption of spherical symmetry and examine the situation of local observations in a universe that contains non-spherically symmetric inhomogeneity. Specifically, we use the inhomogeneous metrics found by Szekeres to describe a non-spherically symmetric universe containing a static star or planet. As in Paper I, we are interested in determining the difference (if any) between variations of a supposed 'constant' or associated scalar field when observed locally, on the surface of the planet or star, and on cosmological scales.  When a `constant', $\\mathbb{C}$, is made dynamical we can allow it to vary by making it a function of a new scalar field, $\\mathbb{C}\\rightarrow \\mathbb{C}(\\phi )$, that depends on spacetime coordinates: $\\phi =\\phi (\\vec{% x},t)$. It has become general practice to combine take all observational bounds on the allowed variations of $\\mathbb{C}(\\phi )$. This practice assumes implicitly that any time variation of $\\mathbb{C,}$ on or near the Earth, is comparable to any cosmological variation that it might experience, that is to high precision \\begin{equation} \\dot{\\phi}(\\vec{x},t)\\approx \\dot{\\phi}_{c}(t),  \\label{wettcond} \\end{equation}% \\noindent for almost all locations $\\vec{x}$, where $\\phi _{c}$ is the cosmological value of $\\phi $. This assumption is always made without proof, and there is certainly no \\emph{a priori} reason why it should be valid. Strictly, $\\phi $ mediates a new or `fifth' force of Nature. If the assumed behaviour is correct then this force is unique amongst the fundamental forces in that its value locally reflects its cosmological variation directly.  In this series of papers we are primarily interested in theories where the scalar field, or `dilaton' as we shall refer to it, $\\phi $, evolves according to the conservation equation \\begin{equation*} \\square \\phi =B_{,\\phi }(\\phi )\\kappa T-V_{,\\phi }(\\phi ), \\end{equation*}% where $T$ is the trace of the energy momentum tensor, $T=T_{\\mu }^{\\mu }$, (with the contribution from any cosmological constant neglected). We absorb any dilaton-to-cosmological constant coupling into the definition of $V(\\phi )$. The dilaton-to-matter coupling $B(\\phi )$ and the self-interaction potential, $V(\\phi )$, are arbitrary functions of $\\phi $ and units are defined by $\\kappa =8\\pi G$ and $c=\\hslash =1$. This covers a wide range of theories which describe the spacetime variation of `constants' of Nature; it includes Einstein-frame Brans-Dicke (BD) and all other, single-field, scalar-tensor theories of gravity \\cite{bd, bsm, poly, posp}. In cosmologies that are composed of perfect fluids and a cosmological constant, it will also contain the Bekenstein-Sandvik-Barrow-Magueijo (BSBM) theory of varying $\\alpha $, \\cite{bsm}, and other single-dilaton theories which describe the variation of standard model couplings, \\cite{posp}. We considered some other possible generalisations in \\cite{shawbarrow1}. It should be noted that our analysis and results apply equally well to any theory which involves weakly-coupled, `light', scalar fields, and not just those that describe variations of the standard constants of physics.  In first paper of this series, \\cite{shawbarrow1}, we determined the conditions under which condition \\ref{wettcond} would hold near the surface of a virialised over-density of matter, such as a galaxy or star, or a planet, such as the Earth, under the assumption of spherical symmetry. We chose to refer to this object as our `star'. In Paper I, matched asymptotic expansions were employed to analyse the most general, \\emph{% spherically-symmetric}, dust plus cosmological constant embeddings of the `star' into an expanding, asymptotically homogeneous and isotropic spherically symmetric universe. We proved that, independent of the details of the scalar-field theory describing the varying `constant', that \\ref% {wettcond} is almost always satisfied under physically realistic conditions. The latter condition was quantified in terms of an integral over sources that can be evaluated explicitly for any local spherical object.  In this paper we extend that analysis, and our main result, to a class of embeddings into cosmological background universes that possess \\emph{no} Killing vectors i.e. \\emph{no} symmetries. The mathematical machinery that we use to do this is, as before, the method of matched asymptotic expansions, employed in \\cite{shawbarrow1}, where the technical machinery is described in detail. A summary of the results obtained there can also be found in \\cite{shawbarrowlett}.  This paper is organised as follows: We shall firstly provide a very brief summary of the method of matched asymptotic expansions used here. In section II we will introduce the geometrical set-up that we will use. We will be working in spacetime backgrounds of Szekeres-Szafron type \\cite{szek, szafron}.  We describe theses particular solutions of Einstein's equations briefly in section II and then in greater detail in section III. In section IV we extend the analysis of \\cite{shawbarrow1} to include non-spherically symmetric backgrounds of Szekeres-Szafron type.  In section V, we consider the validity of the approximations used, and state the conditions under which they should be expected to hold. In section VI we perform the matching procedure (as outlined below), and extend the main result of \\cite% {shawbarrow1} to Szekeres-Szafron spacetimes. We consider possible generalisations of our result in section VII before considering the implications of the results in the section VIII.  We will employ the method of matched asymptotic expansions \\cite{hinch, Death}. We solve the dilaton conservation equations as an asymptotic series in a small parameter, $\\delta $, about a FRW background and the Schwarzschild metric which surrounds our star. The deviations from these metrics are introduced perturbatively. The former solution is called the \\emph{exterior expansion} of $\\phi $, and the latter the \\emph{interior expansion} of $\\phi $. The exterior expansion is found by assuming that the length and time scales involved are of the order of some intrinsic exterior length scale, $L_{E}$. Similarly in the interior expansion we assume all length and time scales we be of the of $L_{I}$, the interior length scale. Neither of the two different expansions will be valid in both regions. In general, we define $\\delta :=L_{I}/L_{E}\\ll 1$. This means that in general only a subset of our boundary conditions we will be enforceable for each expansion, and as a result both the interior and exterior solutions will feature unknown constants of integration. To remove this ambiguity, and fully determine both expansions, we used the formal matching procedure. The idea is to assume that both expansions are valid in some intermediate region, where length scales go like $L_{int}=L_{I}^{\\alpha }L_{E}^{1-\\alpha } $, for some $\\alpha \\in (0,1)$. Then by the uniqueness property of asymptotic expansions, both solutions must be equal in that intermediate region. This allows us to set the value of constants of integration, and effectively apply \\emph{all} the boundary conditions to both expansions. A fuller discussion of this method, with examples, and its application in general relativity is given in \\cite{shawbarrow1}. ",
        "conclusions": "In this paper we have extended the analysis of \\cite{shawbarrow1} to include a class of dust-filled spacetimes without any symmetries provided by the Szekeres-Szafron metrics. Again, we have used the method of matched asymptotic expansions to link the evolution of the dilaton field, $\\phi $, in an approximately Schwarzschild region of spacetime to its evolution in the cosmological background. By these methods, we have provided a rigorous construction of what has been simply assumed about the matching procedures in earlier studies \\cite{early}. We have also analysed, more fully, the conditions that we need the background spacetime to satisfy for the matching procedure to be valid, and we have interpreted these conditions in terms of their requirements on the local energy density. Finally, we have conjectured a generalisation of our result to more general spacetime backgrounds than those considered here.  By combining the results found here with those of the previous paper, we conclude that, in the class of quasi-spherical Szekeres spacetimes in which the matching procedure is valid, the local time variation of the dilaton field will track its cosmological value whenever: \\begin{equation} \\left\\vert \\frac{B_{,\\phi }\\left( \\phi _{c}\\right) \\int_{\\infty }^{r}\\mathrm{% d}r^{\\prime }R_{,r}\\kappa \\Delta (R_{,t}\\varepsilon _{s}(r^{\\prime },t))+\\frac{2% }{3}B_{,\\phi }\\left( \\phi _{c}\\right) R\\int_{\\infty }^{r}\\mathrm{d}r^{\\prime }R_{,r}R_{,t}\\frac{\\kappa \\varepsilon _{ns}(r^{\\prime },t)}{R}+\\frac{1}{3}% B_{,\\phi }\\left( \\phi _{c}\\right) RR_{,t}\\kappa \\varepsilon _{ns}(r,t)}{\\dot{% \\phi}_{c}}\\right\\vert \\ll 1.  \\label{condition1} \\end{equation}% When the cosmological evolution of $\\phi $ is dominated by its matter coupling: $\\dot{\\phi}_{c}\\sim \\mathcal{O}(B_{,\\phi }H_{0}^{-1}\\kappa \\varepsilon _{c}),$ this condition is equivalent to: \\begin{equation*} \\left\\vert H_{0}\\int_{\\infty }^{r}\\mathrm{d}r^{\\prime }R_{,r}\\frac{\\Delta (R_{,t}\\varepsilon _{s}(r^{\\prime },t))}{\\varepsilon _{c}(t)}+\\frac{2}{3}% H_{0}R\\int_{\\infty }^{r}\\mathrm{d}r^{\\prime }R_{,r}R_{,t}R^{-1}\\frac{% \\varepsilon _{ns}(r^{\\prime },t)}{\\varepsilon _{c}(t)}+\\frac{1}{3}H_{0}R_{,t}\\frac{% \\varepsilon _{ns}(r,t)}{\\varepsilon _{c}(t)}\\right\\vert \\ll 1. \\end{equation*}  In the other extreme, when the potential term dominates the cosmic dilaton evolution, the left-hand side of the above condition is further suppressed by a factor of $B_{,\\phi }(\\phi _{c})/V_{,\\phi }(\\phi _{c})\\ll 1$. As in our previous paper, \\cite{shawbarrow1}, we can see that for a given evolution of the background matter density, condition (\\ref{wettcond}) is more likely to hold (or will hold more strongly) when $\\left\\vert B_{,\\phi }(\\phi _{c})\\kappa \\varepsilon _{c}/V_{,\\phi }(\\phi _{c})\\right\\vert \\ll 1$. We reiterate our previous statement that: \\emph{domination by the potential term in the cosmic evolution of the dilaton has a homogenising effect on the time variation of} $\\phi $.  The non-spherically symmetric parts of energy density enter into the expression differently. The magnitude of the terms on the left-hand side of eq. (\\ref{condition1}) is, as in the spherically symmetric case, still $% \\left\\langle H_{0}R\\Delta R_{,t}\\varepsilon /\\varepsilon \\right\\rangle (R,t)$ where $\\left\\langle \\cdot \\right\\rangle (R,t)$ represents some `average' over the region outside the surface $(R,t)=const$. We should note that, given the condition on $\\kappa \\varepsilon $ that has been required for matching, the leading-order contribution to $\\kappa \\varepsilon _{ns}$ is everywhere of dipole form and this is responsible for the special form of the average over the non-spherically symmetric terms. We can also see that, as a result of form of eqn. (\\ref{condition1}), peaks in $\\kappa \\varepsilon _{ns}$ that occur outside of the interior region will, in the interior, produce a weaker contribution to the left-hand side of eqn. (\\ref{condition1}% ) than a peak of similar amplitude in a spherically symmetric energy density $\\kappa \\varepsilon _{s}$. This behaviour would continue if we were also to account for higher multipole terms in $\\kappa \\varepsilon _{ns}$. The higher the multipole, the more `massive' the mode, and the faster it dissipates.  If we are interested in finding a sufficient condition (as opposed to a necessary and sufficient one) for (\\ref{wettcond}) to hold locally, then in most circumstances we will be justified in averaging over the non-spherically symmetric modes in the same way as we average over the spherically symmetric ones. In most cases, this will over-estimate rather than under-estimate the magnitude of the left-hand side of our condition, (% \\ref{condition1}). This reasoning leads us to the statement that for $\\dot{% \\phi}(\\mathbf{x},t)\\approx \\dot{\\phi}_{c}(t)$ to hold locally it is sufficient that: \\begin{equation} \\mathcal{I}:=\\int_{\\gamma (R)}\\mathrm{d}lH_{0}R^{\\prime }\\frac{\\Delta (v\\varepsilon )}{\\varepsilon _{c}}\\ll 1  \\label{suff} \\end{equation}% \\noindent where $\\mathrm{d}l:=\\mathrm{d}rR_{,r}$, $v=R_{,t}$ is the velocity of the dust particles, $\\lim_{R\\rightarrow \\infty }v=H_{0}R$. We make the same generalisation that we did in Paper I by taking $\\gamma (R)$ to run from $R$ to spatial infinity along a past, radially-directed light-ray. In this way, we incorporate the limitations imposed by causality. We should also assume that the above expression includes some sort of average over angular directions; to be safe we could replace $\\varepsilon $ by its maximum value for fixed $R$ and $t$. This sufficient condition, (\\ref{suff}), is precisely the generalised condition proposed in our first paper on this issue. The inclusion of deviations from spherical symmetry, therefore, has little effect of the qualitative nature of the conclusions that were found in \\cite{shawbarrow1}. If anything, we have seen that the non-spherical modes dissipate faster and, as a result, will produce smaller than otherwise expected deviations in the local time derivative of $\\phi $ from the cosmological ones.  On Earth we should expect, as before, that the leading-order deviation of $% \\dot{\\phi}$ from $\\dot{\\phi}_{c}$ is produced by the galaxy cluster in which we sit, and that for a dilaton evolution that is dominated by its coupling to matter, this effect gives $\\mathcal{I}\\approx 6\\times 10^{-3}\\Omega _{m}^{-1}h\\ll 1$, where $\\Omega _{m}\\approx 0.27$ and $h\\approx 0.71$. If the cosmic dilaton evolution is potential dominated then $\\mathcal{I}$ is even smaller. We conclude, as before, that irrespective of the value of the dilaton-to-matter coupling, and what dominates the cosmic dilaton evolution, that \\begin{equation*} \\dot{\\phi}(\\mathbf{x},t)\\approx \\dot{\\phi}_{c}(t) \\end{equation*}% will hold in the solar system in general, and on Earth in particular, to a precision determined by our calculable constant $\\mathcal{I}$. We also conclude, as before, that whenever $\\mathcal{I}\\ll 1$ near the horizon of a black hole, there will be no significant gravitational memory effect for physically reasonable values of the parameters \\cite{memory, jacobson}.  Our result relies on one major assumption: the physically realistic condition that the scalar field should be weakly coupled to matter and gravity -- in effect, the variations of 'constants' on large scales must occur more slowly than the universe is expanding and so their dynamics have a negligible back-reaction on the cosmological background metric. In this paper we have removed the previous condition of spherical symmetry at least in as far as the spacetime background is well described by Szekeres-Szafron solution. We have therefore extended the domain of applicability our general proof: that \\emph{terrestrial} and \\emph{solar system} based observations can legitimately be used to constrain the \\emph{cosmological} time variation of supposed `constants' of Nature and other light scalar fields."
    },
    "0601/astro-ph0601165_arXiv.txt": {
        "abstract": "{ } {We determined the $K_s$ band luminosity function (LF), and inferred the corresponding stellar mass function, of cluster galaxies at redshift $z\\simeq 1.2$,  using near--infrared images of three X--ray luminous clusters at $z=1.11,1.24,1.27$.} {The composite LF was derived down to M$^{*}$+4, by means of statistical background subtraction,  and is well described by a Schechter function with $K_s^{*}=20.5^{+0.4}_{-1}$ and $\\alpha=-1.0^{+0.2}_{-0.3}$. Using available X--ray mass profiles we determined the M/L ratios of these three clusters, which tend to be lower than those measured in the local universe. Finally, from the $K_s$ band composite LF we derived the stellar mass function of cluster galaxies.} {With these data, no significant difference can be seen between the cluster galaxies LF and the LF of field galaxies at similar redshift. We also found no significant evolution out to $z \\simeq 1.2$ in the bright ($<$ M$^{*}$+4) part of the LF probed in this study, apart from a brightening of  $\\simeq 1.3$ mag of the characteristic magnitude of the high redshift LF. We confirm, and extend to higher redshift, the result from previous work that the redshift evolution of the characteristic magnitude M$^{*}$ is consistent with passive evolution of a stellar population formed at $z>2$. } {The results obtained in this work support and extend previous findings that most of the stars in bright galaxies were formed at high redshift, and that $K_s$--bright ($M>10^{11} M_{\\odot}$) galaxies were already in place at $z\\simeq 1.2$, at least in the central regions of X--ray luminous clusters. Together with recent results on the field galaxies stellar mass function,  this implies that most of the stellar mass is already assembled in massive galaxies by z $\\simeq 1$, both in low and high density environments.} ",
        "introduction": "Galaxy clusters are rare systems forming in the highest density peaks of large scale structure. In these special regions galaxy formation and evolutionary processes are expected to be faster with respect to the low density fields, thus making galaxy clusters a biased environment. On the other hand, clusters of galaxies, particularly at high redshift ($z \\simeq 1$), provide a very convenient place for studying the evolution of massive galaxies.  Not only do they contain high numbers of such objects, but these objects turn out to be so evolved (already at $z \\simeq 1$) that they show a colour--magnitude sequence as clear as at lower redshifts.  Thus evolved galaxies in distant clusters can be easily identified even without complete spectroscopic follow--up.  The study of massive galaxies has a relevant role in constraining galaxy formation and evolution models, as different models provide different predictions for their assembly (in particular in the redshift range [$0 \\div 1$]).  They could have rapidly formed their stars at high redshift and at the same time assembled their stellar mass, and then simply evolved passively as their stars aged. Alternatively, massive galaxies could have assembled on a longer time scale in a process of continuous merging of smaller units until redshift $< 1$.  Comparison of these different scenarios has proven to be a difficult task: even if merging galaxies are observed, the relevance of the merging process in galaxy evolution and especially the epoch at which major mergers occur is still debated.  Colours and spectra of massive galaxies at $z \\simeq 1$ show that there is a significant population of such systems already hosting mainly evolved stellar populations, both in the field (see below) and in clusters (\\citet{stanford1997, stanford1998, vandokkum1998, benitez1999, depropris1999, rosati1999, stanford2002, blakeslee2003, vandokkum2003, kodama2004, lidman2004, delucia2004, holden2005b, holden2005c, tanaka2005}). In such studies, the presence of evolved stellar populations is generally inferred from fundamental plane or colour--magnitude sequence evolution. These studies indicate that most of the stars in massive galaxies were formed at $z >2$. At the same time several of them point out that the less massive the galaxy is,  the more likely is the presence of a younger component in its stellar population -- the so called ``downsizing'' \\citep{cowie1996} in galaxies hosting star formation.   It is not possible, however, on the basis of spectrophotometric analysis only, to rule out the possibility that these galaxies formed via merging of smaller galaxies with already evolved stellar populations even a short time before being observed.  For instance, a passively evolving zero--point of the colour--magnitude relation does not imply that the galaxies formed long ago, but that the stars in the galaxies formed at high redshift, possibly in smaller progenitors.  In other words, while the underlying stellar populations can place constraints on the details of the star-formation history, they cannot tell when a galaxy assembled.  In fact, even if these massive galaxies appear to be passively evolving, several studies have noted that to some extent they can still be forming (or recently have formed) stars (for instance \\citet{nakata2001,vandokkum2003,holden2005,demarco2005,jorgensen2005}), likely implying merging (as observed for instance by \\citet{vandokkum1999} and \\citet{tran2005}) or anyway subsequent episodes of star formation. Such secondary episodes of star formation are probably correlated to cluster--related processes (accretion of field galaxies or groups and cluster merging), since galaxies which exhibit these features are often located outside of the cluster core, or in regions of lower X--ray luminosity.  Moreover, the redshift evolution of the brightest cluster galaxies (BCGs) is peculiar and exhibits a large scatter at increasingly high redshift, so that at least in some cases merging could be required to make them evolve into local BCGs (see for instance \\citet{ellisejones}). Even if it is difficult to trace a common ``BCG evolutionary path'', due to the intrinsically peculiar nature of these galaxies, some high redshift clusters (including Cl1252 \\citep{rosati2004b,blakeslee2003} and Cl0848 \\citep{vandokkum2001}) show signs of interactions or clear ongoing merging between few massive galaxies which could lead to the formation of a cD.   However, even if the bulk of the stars had similar ages in the two formation scenarios (i.e. star formation occurring at the same early epoch in both), the epoch of assembly of the final mass observed locally in massive galaxies is different in the two cases.  If merging is a relevant process in the look--back time range that we can probe with our observations, looking at progressively higher redshift one should see the number of massive objects decreasing as they break up into their progenitors; this would cause an evolution in the shape of the mass function of galaxies.   The differences in the predictions of the two formation scenarios have recently become less extreme, partly due to the higher redshift peak of the merging activity in $\\Lambda$CDM models as compared to standard CDM initially considered, but also to different ad--hoc recipes for the star formation adopted in the hierarchical merging models.  It has recently become more evident that, both in clusters (references mentioned above) and in the field (\\citet{glazebrook2004,mccarthy2004,fontana2004,saracco2004, cimatti2002k20, franx2003}), a significant population of massive galaxies is already in place at $z \\simeq 1 \\div 2$. However, see also \\citet{vandokkum2005}, finding that a considerable fraction of a nearby bulge--dominated galaxy sample, recently experienced a merging episode involving more than 20\\% of its final mass. The stellar mass function of bright galaxies shows only a mild evolution up to redshift 1 , close to the prediction of simple pure luminosity evolution \\citep[e.g.][]{fontana2004}.  The comparison with recent semi--analytical models however shows that different renditions predict very different evolution, especially at higher redshift (i.e. results are very sensitive to the chosen model ingredients), and most of them under--produce very massive galaxies (more severely the higher the redshift) even when reproducing the stellar mass function around M$_{*}$ -- however, see also recent results from \\citet{bower2005}. At the same time, \\citet{nagamine2005} show that with recent hydrodynamical simulations they can account for $\\simeq 70\\%$ of the total stellar mass at $z =0$ already being formed by $z=1$.    Since a direct measure of the mass function is too difficult at high redshift for a reasonably large galaxy sample including faint objects, the near--infrared galaxy luminosity function (LF) represents an useful cheaper surrogate. The galaxy LF is a first order description of a galaxy population (density of galaxies as a function of their luminosity). Despite (or because of) its conceptual simplicity the LF has been for many years one of the most popular tools for the interpretation of galaxy observations at all redshifts and in very different environments. The comparison of the LF at different redshifts constrains models of galaxy formation and evolution \\citep{kauffmannecharlot1998b}, while the comparison of the LF in low and high density environments probes the relevance of the environmental effects on the galaxy populations.  The LF is historically best studied in rich clusters of galaxies, which provide large numbers of galaxies at the same distance and, at low redshift, with high contrast against the background, allowing an efficient identification of cluster members with small contamination from background galaxies.  At higher redshift, the faint luminosities and the substantial background contamination makes the LF  determination more uncertain. However, the steadily increasing data quality, and the quest for strict constraints on galaxy evolutionary models, have made the study of the LFs in high redshift clusters and fields a popular topic.  In this work, we determine the LF of distant ($z >1$) cluster galaxies in the near--infrared (NIR).  NIR galaxy samples are particularly well suited for studying galaxy evolution.  Apart from advantages such as the smaller effect of dust extinction (as compared to bluer wavelengths), and the k--corrections relatively insensitive to galaxy type, they provide a relatively good estimate of the stellar mass in galaxies up to redshift $z \\sim 2$.  Therefore, near--infrared luminosity functions can trace the stellar mass function more effectively than bluer band LFs,  which are more sensitive to the star formation histories of the galaxies.   While LFs for cluster galaxies at low redshift ($z \\leq 0.2 \\div 0.3$) have been determined for a large number of clusters, allowing detailed discussion of the features and the separate contributions of different galaxy populations down to very faint magnitudes, the determination of the LF with comparable accuracy at high redshift is clearly more difficult.   The NIR LF of cluster galaxies at high redshift ($z \\geq 0.8$) has been measured by \\citet{trenthamemobasher1998, depropris1999, nakata2001, kodamaebower, toft2003, ellisejones}, and \\citet{toft2004}. The evolution of the characteristic magnitude $M^{*}$ was first studied by \\citet{depropris1999} from low redshift up to $z\\simeq 0.9$, finding that it is consistent with pure luminosity evolution of a stellar population formed at $z >2$; this result has been confirmed by subsequent studies. The evolution of the faint--end slope was only studied by \\citet{toft2003} and \\citet{toft2004}, who found a flatter slope at higher redshift compared to the local value.  The adopted cosmology in this paper is H$_{0} = 70$ km s$^{-1}$ Mpc$^{-1}$, $\\Omega_{M}=0.3$, $ \\Omega_{\\Lambda}=0.7$ unless otherwise stated. Magnitudes are in the AB system. ",
        "conclusions": "We have studied the near infrared luminosity function of high redshift cluster galaxies in three X--ray luminous clusters that are among the most distant discovered so far ($1.1<z<1.3$).  These clusters bear the strongest leverage on evolutionary studies as they probe a redshift regime when the Universe was less than half of its present age.  By measuring the K band luminosity function (and galaxy stellar mass function) in these systems, and by comparing it with that at $z \\simeq 0$, we can set valuable constraints on the galaxy evolution in dense environments in the redshift range [0 $\\div$ 1].  We should note, however, that derived quantities still suffer from large uncertainties, due to small--number statistics and possibly biases in the galaxy populations (due to probing only the cluster cores).   The LF resulting from this work is consistent with previous determinations at similar or lower redshifts, and consolidates a scenario where the evolution of the characteristic magnitude M$^{*}$ ($\\Delta M^{*} \\simeq -1.3$) is consistent with predictions of passive evolution for a stellar population formed at $z >2$. Moreover, we find that the overall shape of the high redshift LF matches the one of the local cluster galaxies LF, once such a brightening is taken into account.  Similarly, the evolution of the K band LF of field galaxies has been found to be consistent with passive evolution up to $z = 1$, with a density evolution lower than 30\\% and a brightening of $\\Delta M^{*} \\simeq 0.5 \\div 0.7$ mag (e.g., \\citet{pozzetti2003,drory2003} -- however, see also \\citet{dahlen2005}).  A direct comparison of the field and cluster LFs at redshift zero, has revealed a significant difference both in the B and in the K band (e.g. \\citet{balogh2001,depropris2003}) amounting to $\\simeq 0.3$ mag in M$^{*}$ and $\\simeq 0.1\\div 0.2$ in $\\alpha$. While there is a hint of exceeding very bright galaxies in our distant clusters with respect to the field, our error bars do not allow us to make such a claim.  For Cl1252, for which best quality data and extended wavelength coverage are available, we attempted a separation of the contributions of late and early type galaxies to the bright end of the LF. We find that, already at $z \\simeq 1.2$, the bright end of the LF appears to be  dominated by early type galaxies selected either on the basis of their morphological appearence or of their spectrophotometric properties.   Using the individual LFs for the three clusters, we calculated the K band cluster M/L ratios making use of X--ray mass profiles derived in \\citet{ettori2004}. The M/L ratio tends to be smaller than the typical value at redshift zero, as expected on the basis of pure luminosity evolution of the cluster galaxies stellar populations. However, a much larger sample would be needed for a detailed investigation of the evolution of the cluster M/L ratio.   Finally, from the composite K$_{s}$ band LF we have estimated the stellar mass function of cluster galaxies.  The observed K$_{s}$ band light at $z \\simeq 1.2$ corresponds approximately to the restframe {\\it z} band light, and is considered as a good tracer of the stellar mass.  We have outlined in the introduction how the determination of the stellar mass function at high redshift sets strong constraints on the evolution of massive galaxies.  While the early formation epoch of the bulk of the stars in massive galaxies is now generally established by several observations, the epoch of the major mass assembly can only be assessed by studying the redshift evolution of the mass function. Our study shows that the massive ($M_{stars} > 10^{10} M_{\\odot}$) galaxy populations in massive high redshift clusters have not significantly changed since $z \\simeq 1$, apart from passive evolution of their stars, thus extending previous results at lower redshifts. The shape of the stellar mass function at $z \\simeq 1.2$ is found to be consistent with the one observed in local clusters, within our uncertainties.  The high-mass end of the LF, made of giant ($M_{\\rm stars}>10^{11} M_{\\odot}$) E/S0 galaxies, is already in place at $z\\simeq 1.2$.  This points toward an early assembly of the galaxy mass, mostly completed before $z \\simeq 1$, thus implying that the bulk of merging activity for massive galaxies in clusters has to occur at much earlier epochs.  This might appear not surprising as field studies have found evidence of a similar early assembly of massive galaxies with most of the stellar mass already assembled in systems more massive than the local characteristic mass by redshift 1 (e.g., \\citet{fontana2004,conselice2005}). Since galaxy evolution in clusters is expected to be faster than in the field in hierarchical galaxy formation scenarios, an even milder evolution of the mass function of cluster galaxies is expected.  The similarity between  the shape of the MF we have found at $z\\simeq 1.2$ in rich clusters and in the field sets upper limits on the difference of formation time--scales of stellar populations in high and low density environments.  For example, if we use the star formation histories for early type galaxies in different environments, for different stellar masses, as derived from the fossil record data \\citep{thomas2005}, our results would imply an age difference of $\\leq 2$Gyr.   The situation is probably different for lower mass galaxies, however it remains difficult to probe the mass function significantly lower than $10^{10} M_{\\odot}$ at $z>1$.  Among the caveats in our study, we should mention the so--called progenitor bias. Since we have considered the whole cluster galaxy population, our work is in principle not affected by the progenitor bias referred to when dealing with selected populations of galaxies (generally early--types \\citep{vandokkum2001b}). It remains true that galaxy populations in clusters at high redshift might not be directly comparable to local ones (e.g., \\citet{kauffmannecharlot1998a}), and that the high--redshift clusters we are observing might not be the progenitors of the local X--ray luminous clusters \\citep{kauffmann1995}, leading to an underestimated evolution.  With the aid of multicolour photometry, including Spitzer/IRAC bands, we can now directly estimate the stellar masses of high redshift cluster galaxies, as well as approximate ages of their stellar populations, and push these studies out to $z=1.4$ (\\citet{mullis2005}, Stanford et al. (2005)), thus probing an epoch which is thought to be crucial for the formation of massive clusters. This work will stimulate significant progress in discriminating between different formation scenarios."
    },
    "0601/astro-ph0601486_arXiv.txt": {
        "abstract": "Extrasolar planet surveys have begun to detect gas giant planets in orbit around M dwarf stars. While the frequency of gas giant planets around M dwarfs so far appears to be lower than that around G dwarfs, it is clearly not zero. Previous work has shown that the core accretion mechanism does not seem to be able to form gas giant planets around M dwarfs, because the time required for core formation scales with the orbital period, which lengthens for lower mass stars, resulting in failed (gas-poor) cores unless the gaseous protoplanetary disk survives for $> 10$ Myr. Disk instability, on the other hand, is rapid enough ($\\sim 10^3$ yrs) that it should be able to form gas giant protoplanets around even low mass stars well before the gaseous disk disappears. A new suite of three dimensional radiative, gravitational hydrodynamical models is presented that calculates the evolution of initially marginally gravitationally unstable disks with masses of 0.021 to 0.065 $M_\\odot$ orbiting around stars with masses of 0.1 and 0.5 $M_\\odot$, respectively. The models show that gas giant planets are indeed likely to form by the disk instability mechanism in orbit around M dwarf stars, the opposite of the prediction for formation by the core accretion mechanism. This difference offers another observational test for discriminating between these two theoretical end members for giant planet formation. Ongoing and future extrasolar planet searches around M dwarfs by spectroscopy, microlensing, photometry, and astrometry offer the opportunity to help decide between the dominance of the two mechanisms. ",
        "introduction": "M dwarf stars dominate the nearby stellar population: within 10 pc of the sun, there are least 236 M dwarfs, but only 21 known G dwarfs (Henry et al. 1997). In spite of this fact, spectroscopic extrasolar planet surveys have concentrated on G dwarfs, in hopes of finding Solar System analogues. Gravitational microlensing surveys at first set only upper bounds ($< 45\\%$) on the frequency of multiple-Jupiter-mass planets orbiting between 1 AU and 7 AU around M dwarf stars toward the Galactic bulge (Gaudi et al. 2002), but recently the first microlensing detections of Jupiter-mass companions orbiting M dwarfs appear to have been accomplished (Bond et al. 2004; Udalski et al. 2005), and more should be on the way. Spectroscopic planet surveys have begun to focus more on M dwarfs, though a few M dwarfs have been included for several years. One of these, GJ 876, has a pair of gas giant planets (Marcy et al. 1998, 2001) as well as an even lower mass planet of uncertain composition (Rivera et al. 2005). Butler et al. (2004) and Bonfils et al. (2005a) have found evidence for giant planets of Neptune-mass orbiting two more M dwarfs, GJ 436 and Gl 581. There are hints from other spectroscopic surveys as well that M dwarfs may have a significant frequency ($\\le 13 \\%$) of short period giant planets (Endl et al. 2003).  Given that M dwarfs have masses as low as $\\sim 0.1$ that of G dwarfs, their most massive planets might be expected to be similarly lower in mass and therefore harder to detect than those orbiting G dwarfs, whose planets range up to masses of $\\sim 10 M_J$ (Jupiter masses). GJ 876 was the first M dwarf found spectroscopically to have planets, implying that its planetary spectroscopic signatures are relatively large. Hence, GJ 876's gas giants may be indicative of the most massive planets to be found around M dwarfs, at least on relatively short period orbits. GJ 876's two gas giant planets have minimum masses of $\\sim 2.1 M_J$ and $\\sim 0.56 M_J$ (Rivera et al. 2005), consistent with this argument, given GJ 876's mass of $0.32 M_\\odot$. If the range of gas giant masses is shifted downward by a factor of $\\sim 3$ or more for M dwarfs compared to G dwarfs, it will take considerably more effort to obtain as complete a census of M dwarf planets compared to G dwarf planets, even though the lower mass of the star itself works in favor of detection by either spectroscopy or astrometry. The intrinsic faintness of M dwarfs also adds to the observational burden, particularly for spectroscopic surveys, which are photon-limited in their precision.  Young M dwarf stars appear to have dust disks similar to those around T Tauri stars (Kalas, Liu, \\& Matthews 2004; Liu et al. 2004; Liu 2004; Mohanty, Jayawardhana, \\& Basri 2005), so there is no {\\it a priori} observational reason to believe that M dwarfs should not be able to form planets in much the same manner as their somewhat more massive siblings.  Relatively little theoretical work has been done on the formation of planets around low mass stars, because of the general obsession with explaining the origin of our own planetary system. Wetherill (1996), however, found that Earth-like planets were just as likely to form from the collisional accumulation of planetesimals around M dwarfs with half the mass of the Sun as they were to form around G dwarfs. Boss (1995) studied the thermodynamics of protoplanetary disks around stars with masses from 1.0 to 0.1 solar mass, and found that the location of the ice condensation point only moved inward by a few AU at most when the stellar mass was decreased through this range. In the core accretion model of giant planet formation (Mizuno 1980), this result implied that gas giant planets should be able to form equally well around M dwarfs, though perhaps at somewhat smaller orbital distances.  Core accretion, however, is a process once thought to have required times of order 10 Myr to form a core massive enough to accrete a gaseous envelope orbiting a G dwarf star (Pollack et al. 1996). Theorists have worked hard to shorten the core accretion time scale to as little as $\\sim 4$ Myr at 5.2 AU (Inaba, Wetherill, \\& Ikoma 2003), given the evidence for typical disk lifetimes of $\\sim 3$ Myr (Haisch, Lada, \\& Lada 2001; Eisner \\& Carpenter 2003) or less in regions of high mass star formation (Bally et al. 1998; Bric\\~eno et al. 2001). The time scale for core growth can be shortened to $\\sim 1$ to 2 Myr by assuming that inward Type-I migration hastens the accumulation of solid cores (Alibert et al. 2005), but Type-I migration appears to be more of a fatal danger to core accretion in the presence of gas than an aid to gas giant planet formation (Kominami, Tanaka, \\& Ida 2005): in spite of efforts to lengthen them, Type-I migration time scale estimates remain considerably shorter than core growth time scales and gaseous disk lifetimes, and considerably shorter than was assumed to be the case by Alibert et al. (2005).  In a turbulent, magnetorotationally unstable disk, Type-I migration can become at least in part a stochastic process, resulting in a component of random walk in semimajor axis and hence a distribution of time scales for inward migration (Laughlin, Steinacker, \\& Adams 2004; Nelson 2005), possibly allowing some cores to accrete gaseous envelopes before migrating inward. However, such turbulence is limited to ionized regions of the disk, whereas the planet-forming midplane from $\\sim$ 1 AU to $\\sim$ 10 AU is likely be magnetically dead (e.g., Gammie 1996), leaving Type-I migration as a significant danger for core accretion.  Core accretion is expected to take even longer to produce a gas giant planet around an M dwarf, as the lower stellar mass leads to longer orbital periods at a given distance from the star. Core accretion thus appears to be too slow to produce Jupiter-mass planets in orbit around M dwarfs before the disk gas disappears (Laughlin et al. 2004). Laughlin et al. (2004) found that core accretion required significantly more than $\\sim 10$ Myr to form a gas giant in orbit around a $0.4 M_\\odot$ star, compared to a similar calculation where a gas giant formed in $\\sim 3$ Myr around a $1 M_\\odot$ star. Roberge et al. (2005) found no evidence for significant gas in the debris disk of the M1 star AU Microscopii, with an age of $\\sim 12$ Myr. Laughlin et al. (2004) concluded that gas giant planets should be rare around M dwarfs, but that failed cores (i.e., roughly Neptune-mass cores that grew too slowly to accrete a significant gaseous envelope) should be common.  Metallicities are notoriously difficult to measure in M dwarfs, but a recent study of wide binary systems including M dwarf secondaries by Bonfils et al. (2005b) has permitted a photometric calibration for M dwarf metallicities, assuming that the M dwarf secondaries have the same metallicities as their F, G, or K primaries. Bonfils et al. (2005b) thereby determined that the metallicities of two of the M dwarfs (Gl 876 and Gl 436) known from Doppler spectroscopy to have giant planets are very close to solar. They also found that the metallicities of 47 nearby M dwarfs are slightly lower on average than those of nearby F, G, or K dwarfs, consistent with the M dwarfs being somewhat older on average. These results suggest that M dwarfs are able to form giant planets even with solar metallicities, though the sample size obviously needs to be increased greatly.  In this paper we examine the possibility of forming gas giant planets around M dwarfs by the competing mechanism of disk instability (Cameron 1978; Boss 1997, 1998, 2003, 2004, 2005; Mayer et al. 2002, 2004). In the disk instability mechanism, a marginally gravitationally unstable protoplanetary disk forms spiral arms that can lead to the formation of gravitationally-bound clumps of Jupiter-mass on a time scale of $\\sim 10$ orbital periods, typically $\\sim 10^3$ yrs or less for a protoplanetary disk at orbital distances of $\\sim 10$ AU. Disk instability represents a potentially rapid mechanism for gas giant planet formation around M dwarfs, if M dwarfs have suitable marginally gravitationally unstable disks. ",
        "conclusions": "These models suggest that there is no reason why M dwarf stars should not be able to form gas giant protoplanets rapidly, if a disk instability can occur. This prediction differs fundamentally from that of the core accretion mechanism, which does not appear to be able to form gas giant planets around M dwarf on time scales shorter than typical disk lifetimes (Lauglin et al. 2004). Disk instability may even be able to form planetary-mass objects in orbit around brown dwarfs with suitably unstable disks, though this has not been demonstrated by the present set of models.  Ongoing and future extrasolar planet searches will answer the question of whether or not gas giant planets form frequently around low mass stars. If M dwarfs turn out not to have very many gas giant planets, this could be interpreted as either a failure of disk instability to produce gas giant planets, or as a failure of M dwarfs to have disks massive enough to be marginally gravitationally unstable. If M dwarfs do turn out to have a significant frequency of gas giant planets, this would imply that disk instability is responsible for their formation, given the apparent inability of core accretion to form gas giants fast enough (Laughlin et al. 2004).  Regardless of the situation for the gas giant planets, the Neptune-mass planets that have been found around M dwarfs are unlikely to have formed by disk instability, unless massive gaseous envelopes can be removed by stellar heating and irradiation, a process that may be reasonably efficient for solar-mass stars (Vidal-Madjar et al. 2004; Hebrard et al. 2005). Rather, such planets are likely to have formed by the same collisional accumulation process that led to the terrestrial planets in our Solar System. Such a finding would not preclude the existence of gas giant planets orbiting at greater distances from M dwarfs, as is the case for GJ 876, planets which presumably formed by disk instability.  Given the extreme imbalance in the numbers of the closest stars of different spectral types, M dwarfs are a natural choice for astrometric planet searches, where closeness is a primary virtue. The low mass of the primary helps as well to enable astrometric detections of even lower mass planets. Ground-based efforts are underway at Carnegie's Las Campanas Observatory to detect brown dwarf and gas giant planet companions to M, L, and T dwarfs by astrometry (Boss et al. 2006). Space-based astrometric searches by NASA's {\\it Space Interferometry Mission (SIM)} should lower the detectable planet mass to Earth-masses or below around M dwarfs. Space-based transit missions (COROT, Kepler) will be able to detect many giant planets orbiting M dwarfs, as well as the ongoing ground-based microlensing and spectroscopic searches. We can expect that all of these efforts will combine to determine the gas giant planet frequency around M dwarfs, and thereby help to determine the most likely mechanism for the formation of their giant planets.  I thank Fred Adams for his valuable comments on the manuscipt and Sandy Keiser for her critical computer systems expertise. This research was supported in part by NASA Planetary Geology and Geophysics grant NNG05GH30G and by NASA Astrobiology Institute grant NCC2-1056. The calculations were performed on the Carnegie Alpha Cluster, the purchase of which was partially supported by NSF Major Research Instrumentation grant MRI-9976645.  \\clearpage"
    },
    "0601/astro-ph0601308_arXiv.txt": {
        "abstract": "The history of the Milky Way Galaxy is written in the properties of its stellar populations.  Here we analyse stars observed as part of surveys of local dwarf spheroidal galaxies, but which from their kinematics are highly probable to be non-members. The selection function -- designed to target metal-poor giants in the dwarf galaxies, at distances of $\\sim 100$~kpc -- includes F-M dwarfs in the Milky Way, at distances of up to several kpc. The stars whose motions are analysed here lie in the cardinal directions of Galactic longitude $\\ell \\sim 270^\\circ$ and $\\ell \\sim 90^\\circ$, where the radial velocity is sensitive to the orbital rotational velocity.  We demonstrate that the faint F/G stars contain a significant population with $V_\\phi \\sim 100$~km/s, similar to that found by a targeted, but limited in areal coverage, survey of thick-disk/halo stars by Gilmore, Wyse \\& Norris (2002). This value of mean orbital rotation does not match either the canonical thick disk or the stellar halo. We argue that this population, detected at both $\\ell \\sim 270^\\circ$ and $\\ell \\sim 90^\\circ$, has the expected properties of `satellite debris' in the thick-disk/halo interface, which we interpret as remnants of the merger that heated a pre-existing thin disk to form the thick disk. ",
        "introduction": "The Milky Way Galaxy, once apparently satisfactorily described by Population I and Population II, is clearly a very complex system (Freeman \\& Bland-Hawthorn 2002). Substructure in kinematics and in coordinate space indicates accretion of (formerly) satellite galaxies (e.g.~the Sagittarius dwarf spheroidal: Ibata, Gilmore \\& Irwin 1994; Ibata et al.~1997; Majewski et al.~2003) and disruption of star clusters (e.g.~Odenkirchen et al.~2003).  Merger and assimilation of satellite galaxies is an inherent part of the formation of large galaxies, like the Milky Way, in the now-canonical $\\Lambda$-Cold-Dark-Matter cosmological models.  Indeed thick disks of spiral galaxies are most probably  relics from an early merger of the evolving thin disk and a smaller satellite galaxy (see Wyse 2005 for a recent discussion of the possibilities), and thus contain unique clues about the early stages of disk galaxy formation.  Models of the formation of thick disks through heating of a thin disk by (dissipationless) mergers usually invoke a satellite of 10-20\\% of the mass of the (pre-existing) stellar disk (e.g.~Velaquez \\& White 1999; Walker, Mihos \\& Hernquist 1996).  The outer parts of the satellites are removed by tidal forces, limited by the relative densities of the satellite and larger system (essentially a Roche-Jacobi criterion). The fate of this debris has been largely ignored in the models, but it is clear that material stripped by tidal forces continues on orbits close to that of the center of mass of the satellite at the time of removal. Thus, for typical satellite/subhalo initial orbital angular momentum that is close to half that of a circular orbit of the same energy (e.g.~Benson 2005; Zentner et al.~2005) one would expect an azimuthal streaming velocity of around 100km/s for satellite debris in a system like the Milky Way with a flat rotation (circular velocity) curve, with amplitude \\footnote{Of course in hierarchical clustering models the potential well of the Milky Way is not fixed, but observations constrain the bulk of the merging for normal galaxies to happen very early (e.g.~Glazebrook et al.~2004; Unavane, Wyse \\& Gilmore 1996; Wyse 2005), and theory predicts an early epoch of massive mergers (e.g.~Zentner \\& Bullock 2003).} $\\sim 220$~km/s. Identification of this satellite debris would provide strong evidence for a merger history of the thick disk.  First indications of a population that could be ascribed to debris from the satellite whose merger caused the thick disk was presented by Gilmore, Wyse \\& Norris (2002; hereafter GWN).  This population was identified in a survey of faint (apparent V-band magnitude $V \\simgt 18$) F/G stars in two lines of sight at intermediate latitude, at longitude $\\ell \\sim 270^\\circ$. The primary signature was kinematics distinct from those of the canonical thick disk, and distinct from the stellar halo, with mean azimuthal streaming velocity $\\sim 100$km/s.  Further analysis (Norris et al.~2006) has derived a characteristic metallicity about a factor of ten below the canonical thick disk, and more characteristic of stars in dwarf galaxies.  But is this a pervasive population, reasonably well-mixed as expected if indeed it resulted from a merger that took place many Gyr ago? ",
        "conclusions": "We have shown that faint F/G stars, $V \\simgt 18$, at intermediate latitudes, consistently show a population with mean azimuthal streaming that corresponds to a lag behind the Sun of $\\sim 100$~km/s. This is dissimilar to the value for the canonical thick disk and to the value for the canonical stellar halo.  We selected stars in restricted ranges of color and magnitude, in both $\\ell = 270^\\circ$ and $\\ell = 90^\\circ$, so that modulo metallicities, the stars are typically at a few kiloparsecs from the Sun, and one can infer that we have found a population of intermediate angular momentum.  This population is detected  in a wide range of Galactic latitude and longitude. Indeed, the K-giant sample of MFF also shows indications of this population, particularly for [Fe/H] $\\simlt -1.5$ (see their Figure~7 (d) and (g)). As noted in the introduction, this mean rotational velocity is what one would expect for debris from a typical shredded satellite.  The thick disk may well have been caused by the most massive merger suffered by the Milky Way since the onset of star formation in the thin disk.  In this case, the debris from the satellite implicated in this `minor merger' would dominate over other debris. We propose that this is indeed the case, and interpret our observations as tracing the remains of the satellite whose merger with the Milky Way created the thick disk.  We intend to make detailed comparisons with simulations of the merger process to constrain further this last significant merger of the Milky Way."
    },
    "0601/astro-ph0601622_arXiv.txt": {
        "abstract": "The frequency of short gamma-ray bursts (GRBs) in galaxies with distinct star formation histories can be used to constrain the lifetime of the progenitor systems. As an illustration, we consider here the constraints that can be derived from separating the host galaxies into early and late types. On average, early-type galaxies have their stars formed earlier than late-type galaxies, and this difference, together with the time delay between progenitor formation and short GRB outburst, leads to different burst rates in the two types of hosts.  Presently available data suggest, but not yet prove, that the local short GRB rate in early-type galaxies may be comparable to that in late-type galaxies. This suggests that, unlike Type Ia supernovae, at least half of the short GRB progenitors that can outburst within a Hubble time have lifetimes greater than about 7 Gyr. Models of the probability distribution of time delays, here parametrized as $P(\\tau)\\propto \\tau^n$, with $n \\gtrsim -1$ are favored. This apparent long time delay and the fact that early-type galaxies in clusters make a substantial contribution to the local stellar mass inventory can explain the observed preponderance of short GRBs in galaxy clusters. ",
        "introduction": "\\label{sec:intro}  The progenitors of short duration, hard spectrum, gamma-ray bursts (GRBs) are not yet well identified. Even with the recent localizations of a handful of short-hard GRBs \\citep{Bloom06,Gehrels05,villa,Berger05,Fox05,Berger06,Berger07}, no transient emission has been found that directly constrains the nature of the progenitor system. The current view of most researchers is that GRBs arise in a very small fraction of stars that undergo a catastrophic energy release event toward the end of their evolution \\cite[e.g.,][]{moch93,ls76,elps89,nara92,rj99,kl98,rrr03,rrrd03,srj04, Lee04,aloy,andrew,levan,LeeRR07}. Much of the current effort is dedicated to understanding the different progenitor scenarios and trying to determine how the progenitor and the burst environment can affect the observable burst and afterglow characteristics \\citep[e.g.,][]{Lee05}. The lifetime of progenitors of short bursts can be meaningfully constrained by properties of their host galaxies (e.g., \\citealt{Fox05,GalYam05}). We suggest here a method to constrain the lifetime of the progenitors by making use of the star formation histories of their host galaxies.  Based on the short burst afterglows localized so far \\citep{Berger07}, two of the three relatively nearby ($z\\lesssim 0.3$) events (GRB 050509b and GRB 050724) are plausibly associated with galaxies exhibiting characteristic early-type spectra \\citep{Berger05,Bloom06, Gehrels05,Prochaska06}. In the other case (GRB 050709), the host galaxy shows signature of a dominant stellar population with age of $\\sim$ 1Gyr \\citep{Covino06} and it also exhibits strong emission lines that indicate ongoing star formation \\citep{Covino06,Fox05,Prochaska06}. There is also independent support that, at least two of the nine short bursts localized so far \\citep{Berger07} are associated with clusters of galaxies \\citep{Bloom06,Gladders05,Ped05}.  In contrast to what is found for long-soft GRBs, for which all of the confirmed host galaxies are actively forming stars \\citep[e.g.,][]{trentham,Christensen04}, these observations clearly signify that, like Type Ia supernovae, short GRBs are triggered in galaxies of all types. What is more, it indicates that there is a time delay between short burst occurrence and the main epoch of star formation activity in the hosts, as determined by the progenitor's lifetime.  In this paper, we present a method to deduce the lifetime of short GRB progenitors from the burst rates in host galaxies of different types that have distinct star formation histories. Limited by the accuracy of how well we can separate their contribution to the star formation history, we show, with this constraint in mind, how a large sample host galaxies could be used to determine the lifetime distribution of short GRB progenitors. Here we chose to separate the formation history of stars residing in early- and late-type galaxies, but, as shown in \\S 5, our study can be easily generalized to any set of subpopulations of galaxies possessing distinct star formation histories. The layout is as follows. In \\S~2, we review the local stellar budgets in early- and late-type galaxies. The total star formation history is subsequently decomposed in \\S~3 into a sum of the stars that assembled in today's early- and late-type galaxies.  With the decomposition in place, in \\S~4, we show how the frequency of short GRBs in a well-defined sample of host galaxies of different types can be used to constrain the lifetime of the progenitors. Finally, in \\S~5 we summarize our results and outline future prospects. Throughout this paper, we assume a spatially-flat $\\Lambda$CDM cosmology with $\\Omega_m=1-\\Omega_\\Lambda=0.3$ and the Hubble constant $h=0.7$ in units of $100{\\rm km\\, s^{-1}Mpc^{-1}}$. ",
        "conclusions": "\\begin{figure*} \\plotone{f4.eps} \\caption[]{ \\label{fig:ratio} Similar to Fig.~\\ref{fig:SFR}, but the SFH is decomposed into those for galaxies with different stellar masses. This decomposition is based on the results in \\citet{Panter06}, where the SFHs of local, individual galaxies are inferred from modeling their spectra and are then subsequently grouped according to the present stellar mass content. In the left panel, the solid curve shows a spline fit to the overall SFH inferred by \\citet{Panter06}, and the dashed and dotted curves give the spline fits to the SFRs of galaxies with stellar mass below and above $10^{11}\\Msun$, respectively. For lookback time larger than $\\sim$12.5 Gyr, a cutoff profile is introduced. In the right panel, the ratio of short GRB rates in galaxies with stellar mass above and below $10^{11}\\Msun$ is plotted as a function of the index $n$ of the progenitor lifetime distribution $P(\\tau)\\propto \\tau^n$. The purpose of the figure is to illustrate that the method presented in this paper can be generalized to any sub-populations of galaxies that have distinct SFHs (see text).  } \\end{figure*}  In the local universe, based on the stellar mass functions of galaxies, about 55--70\\% of the stellar mass is in early-type galaxies, and, from the SFH we inferred, the corresponding stars mainly formed about 9 Gyr ago.  Two of the three host galaxies of local short GRBs ($z\\lesssim 0.3$) are associated with old and massive galaxies with little current or recent star formation, which makes it unlikely that short bursts are associated with massive stars. Presently available data suggests, but not yet prove, a long time delay between the formation of the progenitor system and the short GRB outburst --- for progenitors that can outburst within a Hubble time, about half of them have lifetime longer than $\\sim$7 Gyr. It is fair to conclude that, based on the current host galaxy sample, the progenitors of short GRBs appears to be longer lived than those of Type Ia supernovae. \\citet{Fox05} also reach similar conclusion by arguing that Type Ia supernovae occur more frequently in late-type, star-forming galaxies. Simply based on a comparison of Hubble types between short GRBs and Type Ia supernovae host galaxies, \\citet{GalYam05} argue that the delay time of short GRBs should be several Gyr even if the Type Ia supernovae delay time is as short as $\\sim 1$ Gyr.  The lifetime of the progenitor systems is estimated here by using the SFH of elliptical galaxies from a galaxy formation model. This allow us to separate the early- and late-type galaxy contributions to the overall cosmic SFH.  It would, however, be more self-consistent to infer SFHs of different types of galaxies by modeling the observed spectra with stellar population synthesis models. In either case, the uncertainty in the derived SFHs should be folded into the errors derived by this method for the distribution of time delays of short GRB progenitors.  In our calculation, different definitions of early-type galaxies (by color or by light concentration) introduce uncertainties in the resultant SFH. This, in principle, would not be a problem since we can choose to use the same definition for identifying the short GRB host galaxy type. The method proposed here is not limited to a separation between early and late galaxies. As long as galaxies are divided into two (or more) sub-populations that have distinct star formation histories, they can be used to constrain the lifetime distribution of GRB (or Type Ia supernova) progenitors, following the same reasoning presented here.  For example, \\citet{Shin06} recently applied our approach to field and cluster elliptical galaxies. Another example would be to divide galaxies according to their stellar mass. \\citet{Panter06} present SFH of local galaxies as a function of galaxy stellar mass by modeling their spectra. Low mass galaxies on average form their stars later than high mass galaxies, which can be used to constrain $P(\\tau)$ if the stellar mass of GRB hosts can be obtained. As an illustration, in Figure~\\ref{fig:ratio}, we show the separation of the SFR of galaxies with low and high stellar masses (below or above $10^{11}\\Msun$) using results by \\citet{Panter06}. The shape of the overall SFH used in \\citet{Panter06} is slightly different from that in Figure~\\ref{fig:SFR}, which may reflect some systematics in the derivation of the SFH based on the fossil record.  If the systematics can be well controlled, this approach can become even more powerful than the one used here by inferring SFHs of individual GRB host galaxies. A maximum likelihood method can then be used to constrain $P(\\tau)$ based solely on the SFH of host galaxies, which puts our proposed method to an extreme --- dividing the galaxy populations into individual galaxies. Furthermore, instead of assuming a functional form of $P(\\tau)$, such an application to a large number of individual host galaxies may allow constraints on a non-parametric form of the distribution [i.e., constraining $P(\\tau)$ in different $\\tau$ bins].  In this paper, we limit our study to $z\\sim 0$ galaxies, but the method can be easily generalized to galaxy populations at any redshift provided that one can accurately infer their SFHs (e.g., through spectral synthesis). The observational task can be minimized by focusing on the observations of host galaxies of individual short GRBs.  However, for applications at high redshift to be useful, the luminosity function of short GRBs needs to be understood.  Throughout this paper we have assumed the same lifetime distribution of short GRB progenitors in both early- and late-type galaxies. However, star formation processes in these two types of galaxies may not be identical. For example, elliptical galaxies can form by the merging of two gas-rich galaxies \\citep[e.g.,][]{mihos94}. Many globular clusters can form in the merging process \\citep{schweizer}, which could enhance, for example, the fraction of binary progenitors and also change the lifetime distribution (e.g., \\citealt{Grindlay06}).  The magnitude of this kind of effect on the GRB progenitors is a formidable challenge to theorists and to computational techniques. It is, also, a formidable challenge for observers, in their quest for detecting minute details in extremely faint and distant sources.  At least two short GRB host galaxies are found in cluster environments. There may exist a selection bias of detecting short GRBs in a dense medium \\citep{Bloom06}. To study the association of short GRBs with clusters, it would be useful to separate the stellar mass function into that for field galaxies and that for cluster galaxies in addition to early- and late-type galaxies. More promising for the immediate future, the preponderance of cluster environments can be investigated observationally. Important information may be gained by studying the local stellar mass inventory shown in Figure~\\ref{fig:mf}.  Approximately 50\\% of the stellar mass contents in early-type galaxies are in galaxies with $M_{\\rm star}>10^{11}\\Msun$ that typically reside in clusters.  Since it is likely that in clusters galaxies shut off their star formation process early on, a long progenitor lifetime further increases the tendency for short GRBs to happen in cluster galaxies. It is fair to conclude that the observed preponderance of cluster environments for short GRBs is consistent with an old stellar population that preferentially resides in early-type galaxies.  Detailed observations of the astrophysics of individual GRB host galaxies may be essential before stringent constraints on the lifetime of short GRB progenitors can be placed. If confirmed with further host observations, this tendency of short GRB progenitors to be relatively old can help differentiate between various ways of forming a short GRB."
    },
    "0601/astro-ph0601078_arXiv.txt": {
        "abstract": "{While low mass clouds have been relatively well studied, our picture of high-mass star formation remains unclear. Infrared Dark Clouds appear to be the long sought population of cold and dense aggregations with the potential of harbouring the earliest stages of massive star formation. Up to now there has been no systematic study on the temperature distribution, velocity fields, chemical and physical state toward this new cloud population.} {Knowing these properties is crucial for understanding the presence, absence and the very potential of star formation. The present paper aims at addressing these questions. We analyse temperature structures and velocity fields and gain information on their chemical evolution. } {We mapped the $(J,K) = (1,1)$ and (2,2) inversion transitions of ammonia in 9 infrared dark clouds. Our observations allow the most reliable determination of gas temperatures in IRDCs to date.} {The gas emission is remarkably coextensive with the extinction seen at infrared wavelengths and with the submillimeter dust emission. Our results show that IRDCs are on average cold ($T < 20 ~ {\\rm K}$) and have variations among the different cores. IRDC cores are in virial equilibrium, are massive ($M > 100 ~$\\Msol), highly turbulent (1 -- 3~$ \\rm km~s^{-1}$) and exhibit significant velocity structure (variations around 1 -- 2~$\\rm km~s^{-1}$ over the cloud). } {We find an increasing trend in temperature from IRDCs with high ammonia column density to high mass protostellar objects and hot core/Ultracompact H{\\sc ii} regions stages of early warm high-mass star formation. The linewidths of IRDCs are smaller than those observed in high mass protostellar objects and hot core/Ultracompact H{\\sc ii} regions. On basis of this sample, and by comparison of the ammonia gas properties within a cloud and between different clouds, we infer that while active star formation is not yet pervasive in most IRDCs, local condensations might collapse in the future or have already begun forming stars. } ",
        "introduction": "Infrared dark clouds (IRDCs) are cold, dense molecular clouds seen in silhouette against the bright diffuse mid-infrared (MIR) emission of the Galactic plane. They were discovered during mid-infrared imaging surveys with the Infrared Space Observatory (ISO, \\citealt{perault1996:iso}) and the Mid-course Space Experiment (MSX, \\citealt{egan1998:irdc}).  In an initial census of a $\\sim$180\\degr\\ long strip of the Galactic plane (between $269$\\degr $< l < 91$\\degr, $\\rm b = \\pm 0.5$\\degr), \\citet{egan1998:irdc} found $\\sim$2000 compact objects seen in absorption against bright mid-infrared emission from the Galactic plane. Examination of 2MASS, MSX and IRAS images of these objects reveals that they appear as shadows at all these wavelengths, although they are best identified in the $8.3 \\mu$m MSX band, because, first, the 7.7 and 8.6~$ \\mu$m PAH features associated with star-forming regions contribute to a brighter background emission and, second, the MSX $8.3 \\mu$m band is more sensitive than the satellite's other bands. Recently, we have reviewed the observational studies on IRDCs (\\citealt*{menten2005:iau}).    While low mass clouds have been relatively well studied, our picture of high-mass star formation % remains unclear \\citep[see][]{evans2002:aspc}.  IRDCs appear to be the long sought population of cold and dense aggregations with the potential of harbouring the earliest stages of massive star formation. It is likely that some of the stars forming in them are massive (luminosities of submm condensations range up to 10$^{4}$ \\Lsol). A recent study by \\citet{ormel2005:irdc} on an IRDC towards the W51 Giant Molecular Cloud (GMC) suggests that sources of $\\sim 300$~\\Msol\\ are embedded within the cores, most likely protostars. Recently, we \\citep{pillai2005a:g11} reported a detailed study of the strongest submm peak in the IRDC G11.11$-$0.12, where we find clear evidence of a heavily embedded protostar.    The salient results on IRDCs are summarised below.  The IRDCs observed so far have sizes of 1-10~pc and have mostly a filamentary morphology. On the basis of LVG calculations of mm \\ALD\\ observations \\citet{carey1998:irdc} find that typical IRDCs have gas densities of $n > 10^{6} ~ {\\rm cm}^{-3}$ and temperatures of $T < 20 ~ {\\rm K}$. Kinematic distances determined from \\ALD\\ observations using a standard Galactic rotation curve that ranges between 2.2 and 4.8~kpc (\\citealt{carey1998:irdc}) indicate that the clouds are not local. All observed IRDCs in the sample of \\citet{carey2000:irdc} contain 1 -- 4 bright sub-millimeter (submm) dust continuum emission peaks ($>$ 1 Jy/ 14 \\arcsec\\ beam at 850 $\\mu$m) surrounded by an envelope of emission which matches the morphology of the IRDC in mid-infrared extinction. The cores corresponding to the brightest submm peaks have masses 100 and 1200 \\Msol, except for two clouds in the Cygnus region that have masses around 40 \\Msol.  \\citet{hennebelle2001:irdc} in a systematic analysis of the ISOGAL images extracted about 450 IRDCs, for which they derive $15 \\mu$m optical depths of 1 to 4. \\citet{teyssier2002:irdc} reported that Large Velocity Gradient (LVG) model calculations of HC$_3$N, $^{13}$CO, and C$^{18}$O yield densities larger than $10^5$ \\percc\\ in the densest parts. The authors claim to find kinetic temperatures between 8 and 25 K based on CH$_3$CCH observations; the higher values being found toward embedded objects, however a detailed analysis is hitherto unpublished.  Observations of molecules in IRDCs so far have concentrated on just a few species.  Up to now there has been no systematic study on the temperature distribution, velocity fields, chemical and physical state toward this new cloud population.  Knowing these properties is crucial for understanding the presence, absence and the very potential of star formation.  The present paper aims at addressing these questions.  We have started with a survey of a sample of 9 IRDCs in (1,1) and (2,2) cm rotational transitions of ammonia (\\AMM). This sample has been studied before by \\citet{carey1998:irdc} and \\citet{carey2000:irdc} and was selected on the basis of the large extent and high contrast of the IRDCs against the MIR background.    Ammonia has proven to be an important tool in measuring the physical conditions in molecular clouds (\\citealt{ho1983:nh3}). Since only the lowest \\AMM\\ energy levels are expected to be populated in cool dark clouds  ($T<20$~K), their physical conditions  can be probed using the (1,1) and (2,2) inversion transitions in the metastable  ($J,K$) rotational levels of ammonia. Radiative transitions between different $K$-ladders are forbidden, therefore the lowest levels are populated only via collisions. The optical depth can be determined from the ratio of the hyperfine satellites. Thus, the population of the different levels can be estimated and hence the temperature of the gas determined. In addition, recent chemical models reveal that \\AMM\\ (and also N$_2$H$^+$), does not deplete from the gas phase for the densities observed in IRDCs ($<10^{6}$~\\percc) \\citep{bergin1997:chemistry}. Thus \\AMM\\ is an excellent tracer of the dense gas where many other molecules would have heavily depleted.  In \\S\\ref{sec:observations}, we give details of our observations with the Effelsberg 100m telescope. In \\S\\ref{sec:results}, we discuss the data reduction and present the correlation between gas emission and MIR absorption. In \\S\\ref{sec:analysis}, we derive the rotational temperature, gas kinetic temperature and \\AMM\\ column density. Furthermore, we analyse the velocity structure, estimate dust mass, virial mass and \\AMM\\ abundance. In \\S\\ref{sec:disc}, we compare the core gas properties (temperature, \\AMM\\ linewidths and column density) with that of other populations of objects that are thought to trace the early stages of high mass star formation. We do a similar comparison with % local dark clouds well-studied in \\AMM. Finally, we speculate on a possible formation mechanism of IRDCs involving  supernova remnants. ",
        "conclusions": "In this paper, we discussed ammonia observations of a selected set of IRDCs and the derived physical properties.  Our results are as summarised below.  The ammonia emission correlates very well with MIR absorption and ammonia peaks distinctly coincide with dust continuum peaks. Several cores are detected within the clouds with deconvolved sizes smaller than the 40$''$ FWHM beam size.  We can constrain the average gas temperature to between 10 and 20 K.  We observe high linewidths ($1 \\le \\Delta v/\\rm km~s^{-1} \\le 3.5$), hence turbulence plays an important role in the stability of an IRDC. There are significant velocity gradients observed between the cores. The effect of external shock/outflow tracers , on the gas kinematics is suggestive in some cases, but needs to be investigated further.  The column densities translate to extremely high $A_V$ values (55 -- $450$~mag), therefore any active star formation would be heavily embedded.  The total cloud gas mass derived from the \\AMM\\ data ranges from 10$^3$ -- 10$^4$ \\Msol.  The virial parameter is $\\sim 1$ for most of the clumps, and the cores appear to be stable against gravitational collapse.  As a result IRDCs are potential sites for star formation.  If we were to adopt the stellar IMF and a star formation efficiency of 30\\%, then every IRDC core could fragment to form $>100$ stars, with at least two high mass stars ($>8$~\\Msol).  The fractional abundance of \\AMM\\ (relative to $\\rm H_2$) is 0.7 -- $10.0\\times 10^{-8}$. This, together with the excellent correlation in morphology of the dust and gas, is consistent with the time dependent chemical model for \\AMM\\ of \\citet{bergin1997:chemistry} and implies that \\AMM\\ remains undepleted. The derived abundance is a factor 5 -- 10 larger than that observed in local dark clouds while \\ALD\\ is underabundant by a factor of $\\sim 50$. Hence, the chemistry governing these IRDCs might be complex and could be different from other parts of the dense ISM.  The time scales we derive for the clouds to disperse due to their own internal motions of a few Myrs provides an upper limit to the life time of these clouds.  We suggest that SNRs might be the trigerring mechanism responsible for the formation of an IRDC.  The comparison of the physical properties from ammonia of our IRDCs sample with other source samples -- HMPO's and UCH{\\sc ii}s strongly suggests that most of these IRDCs are the most likely candidates for pre-protostellar cores of massive star formation.                  \\begin{table*} \\begin{center} \\caption{List of IRDCs Observed in \\AMM (1,1) and (2,2)}. \\vspace*{2mm} \\begin{tabular}{llldd} \\hline \\hline Name   &        R.A.(J2000)  &     Dec.(J2000)     &  \\dlabel{$V_{\\rm LSR}$~[\\kms]}  &       \\dlabel{D~[kpc]}   \\\\ \\hline   G11.11-0.12 P1 & 18:10:29.27 & --19:22:40.3 & 29.2 &  3.6  \\\\ G19.30+0.07 & 18:25:56.78 & --12:04:25.0   &  26.3 &  2.2  \\\\ G24.72-0.75 & 18:36:21.07 & --07:41:37.7 & 56.4 & 1   \\\\ G24.63+0.15 & 18:35:40.44 & --07:18:42.3 & 54.2 & 3.6  \\\\ G28.34+0.06 P1 & 18:42:50.9 & --04:03:14 & 78.4 & 4.8 \\\\ G28.34+0.06 P2 & 18:42:52.4 & --03:59:54 & 78.4 & 4.8 \\\\ G33.71-0.01 & 18:52:53.81 & +00:41:06.4 & 104.2 & 7.2 \\\\ G79.27+0.38 & 20:31:59.61 & +40:18:26.4 & 1.2  & 1\\\\ G79.34+0.33 & 20:32:21.803 & +40:20:08.00 & 0.1 & 1\\\\ G81.50+0.14 & 20:40:08.29 & +41:56:26.4 & 8.7 & 1.3\\\\ \\hline  \\end{tabular} \\label{tab:source_list} \\end{center} Notes: Columns are name, right ascension, declination, LSR velocity, distance. Positions and Distances are taken from \\citet{carey1998:irdc}. The coordinates correspond to the reference positions of the maps.  \\end{table*}     \\begin{sidewaystable*} \\begin{minipage}{\\textheight} \\caption{\\AMM\\ (1,1) and (2,2) Map Results: Peak Position.}. \\label{tab:line_parameters} \\centering \\begin{tabular}{rclllll} \\hline\\hline Source                    & Transition & $V_{\\rm LSR}$   &  $T_{\\rm MB}$         &  FWHM              &  $\\tau_{\\rm main}$ &  1~$\\sigma$  \\\\ &           &  [\\kms]            &  [K]                  &  [\\kms]              &                   &   [K~\\kms]    \\\\ \\hline G11.11$-$0.12 \\AMM\\ P1    & 1--1 &30.44   ($\\pm$0.02) & 5.47  ($\\pm$0.96) & 1.27  ($\\pm$0.06) & 3.52  ($\\pm$0.39) &0.42 \\\\ & 2--2 &30.42   ($\\pm$0.09) & 1.97  ($\\pm$0.32) & 1.56 ($\\pm$0.30) &          &     \\\\ G19.30$+$0.07 \\AMM\\ P1    & 1--1 &26.45   ($\\pm$0.03) & 5.17  ($\\pm$0.8)  & 1.88  ($\\pm$0.08) & 1.88  ($\\pm$0.26) &0.53 \\\\ & 2--2 &26.64   ($\\pm$0.06) & 2.86  ($\\pm$0.27) & 1.57   ($\\pm$0.17) &          &     \\\\ G24.72$-$0.75 \\AMM\\ P1    & 1--1 &58.37   ($\\pm$0.07) & 2.76  ($\\pm$0.45) & 2.57  ($\\pm$0.12) & 2.81  ($\\pm$0.47) &0.25 \\\\ & 2--2 &58.47   ($\\pm$0.12) & 1.99  ($\\pm$0.23) & 2.97  ($\\pm$0.24) &          &     \\\\ G24.63$+$0.15             & 1--1 &53.11   ($\\pm$0.04) & 2.40   ($\\pm$0.34) & 1.7   ($\\pm$0.09) & 2.36  ($\\pm$0.43)  &0.16\\\\ & 2--2 &52.89   ($\\pm$0.09) & 0.98  ($\\pm$0.12) & 2.42 ($\\pm$0.24) &          &     \\\\ G28.34$+$0.06 \\AMM\\ P1    & 1--1 &79.14   ($\\pm$0.02) & 2.88  ($\\pm$0.42) & 2.67  ($\\pm$0.04) & 2.00  (0.11)  &0.16\\\\ & 2--2 &79.15   ($\\pm$0.04) & 1.37  ($\\pm$0.05) & 3.29  ($\\pm$0.11) &           & \\\\ G33.71$-$0.01 \\AMM\\ P1    & 1--1 &104.31  ($\\pm$0.09) & 2.52  ($\\pm$0.43) & 2.64  ($\\pm$0.2)  & 1.5   ($\\pm$0.52)  &0.12\\\\ & 2--2 &104.4   ($\\pm$0.17) & 1.14  ($\\pm$0.22) & 2.40 ($\\pm$0.35) &           & \\\\ G79.27$+$0.38 \\AMM\\ P1    & 1--1 &1.65    ($\\pm$0.02) & 2.28  ($\\pm$0.38) & 1.04  ($\\pm$0.05) & 2.49  ($\\pm$0.27)  &0.31\\\\ & 2--2 &1.63    ($\\pm$0.05) & 0.61  ($\\pm$0.08) & 0.98  ($\\pm$0.14) &           & \\\\ G79.34$+$0.33 \\AMM\\ P1    & 1--1 &0.31    ($\\pm$0.01) & 2.53  ($\\pm$0.30)  & 1.23  ($\\pm$0.03) & 1.45  ($\\pm$0.13)  &0.19\\\\ & 2--2 &0.25    ($\\pm$0.04) & 0.79  ($\\pm$0.07) & 1.29  ($\\pm$0.13) &           & \\\\ G81.50+0.14               & 1--1 &-4.53   ($\\pm$0.05) & 2.21  ($\\pm$0.35) & 1.35  ($\\pm$0.15) & 0.37  ($\\pm$0.46) &0.27 \\\\ & 2--2 &-4.69   ($\\pm$0.12) & 0.67  ($\\pm$0.21) & 1.00  ($\\pm$0.34) &              & \\\\  \\hline \\hline \\end{tabular}  \\vfill Notes: Columns are name, \\AMM\\ (J,K) transition, LSR velocity, (1,1) main beam brightness temperature, full linewidth at half maximum, main group optical depth and the 1~sigma noise level in the \\AMM\\ (1,1) integrated intensity map (\\ref{fig:msx_nh3}). The error estimates are given in brackets. \\end{minipage} \\end{sidewaystable*}   \\begin{sidewaystable*} \\begin{minipage}{\\textheight} \\caption{Physical Properties of Observed IRDCs.} \\centering {\\scriptsize \\begin{tabular}{ccrrrrcdddddc} \\hline \\hline Source                        & Offsets       & $T_{\\rm EX}$     & $T_{\\rm ROT}$   & $T_{\\rm KIN}$ & $N({\\rm NH_3})$ & $N({\\rm H+H_2})$ & \\dlabel{$\\chi_{\\rm NH_3}$}  & \\dlabel{Mass}  & \\dlabel{Size} & \\dlabel{$M_{\\rm vir}$} & \\dlabel{$\\alpha$} & \\dlabel{$M_{\\rm NH_3 }$} \\\\  & [arcsecs]   & [K]    & [K] & [K]    & [$10^{14}~{\\rm cm^-2}$]   &  \\dlabel{[$10^{22}~{\\rm cm^-2}$]} &  \\dlabel{[$10^{-8}$]}   &\\dlabel{$M_{\\odot}$}    & \\dlabel{[arcsecs]}& \\dlabel{[$M_{\\odot}$]} & \\dlabel{}& [$10^4$~$M_{\\odot}$] \\\\ \\hline G11.11$-$0.12  \\AMM\\ P1       & (80,40)       &   7.6 ($\\pm$0.7) & 12.3 ($\\pm$1.2) & 12.7 & 41.5 ($\\pm$ 6.4)  &  8.4 & 4.9  & 485 & 27 & 468   & 1.0    & 18.1 \\\\ G11.11$-$0.12 \\AMM\\ P2        & (0,0)         &   6.6 ($\\pm$0.4) & 13.4 ($\\pm$1.0) & 13.8 & 30.5 ($\\pm$ 3.3)  & 10.2 & 2.9  & 172 & 13 & 268   & 1.6    &   \\\\ G11.11$-$0.12 \\AMM\\ P3        & (-300,-325)   &   7.1 ($\\pm$0.3) & 14.1 ($\\pm$9.7) & 14.5 & 60.4 ($\\pm$ 11.4) &  8.7 & 6.9  & 647 & 32 & 860   & 1.3    &    \\\\ G11.11$-$0.12 \\AMM\\ P4        & (-300,-400)   &   5.7 ($\\pm$0.5) & 10.5 ($\\pm$1.0) & 10.8 & 50.9 ($\\pm$ 10.6) & 16.1 & 3.2  & 788 & 25 & 555   & 0.7    &      \\\\ G19.30$+$0.07 \\AMM\\ P1$^{n}$  & (20,20)       &   8.2 ($\\pm$0.9) & 17.6 ($\\pm$2.0) & 18.4 & 25.9 ($\\pm$ 3.2)  &      &     &     & 38 & 893   &        & 0.9 \\\\ G19.30$+$0.07 \\AMM\\ P2 $^{n}$ & (-60,-20)     &   5.9 ($\\pm$0.7) & 13.8 ($\\pm$1.8) & 14.3 & 33.3 ($\\pm$ 6.9)  &      &     &     & 45 & 823   &        &        \\\\ G24.72$-$0.17 \\AMM\\ P1 $^{n}$ & (150,220)     &   6.6 ($\\pm$0.6) & 17.4 ($\\pm$2.2) & 18.2 & 44.3 ($\\pm$ 5.7)  &      &     &     & 34 & 2450  &        & 1.4 \\\\ G24.63$+$0.15$^{n}$           & (0,0)         &   5.0 ($\\pm$0.5) & 14.3 ($\\pm$1.4) & 14.8 & 24.1 ($\\pm$ 4.4)  &      &     &     & 43 & 1358  &        & 0.7 \\\\ G28.34$+$0.06 \\AMM\\ P1        & (0,0)         &   5.7 ($\\pm$0.2) & 16.0 ($\\pm$1.2) & 16.6 & 32.1 ($\\pm$ 2.2)  &  5.8 & 5.5  &904  & 22 & 2280  & 2.5    &  25.4 \\\\ G28.34$+$0.06 \\AMM\\ P2        & (20,200)      &   5.5 ($\\pm$0.2) & 15.4 ($\\pm$1.4) & 16.0 & 38.6 ($\\pm$ 3.2)  & 32.7 & 1.2  &2310 & 45 & 2928  & 1.3    &  \\\\ G28.34$+$0.06 \\AMM\\ P3        & (-60,-60)     &   4.8 ($\\pm$0.2) & 12.8 ($\\pm$1.2) & 13.2 & 47.2 ($\\pm$ 6.1)  &  5.9 & 8    &374  & 21 & 1398  & 3.7    &      \\\\ G28.34$+$0.06 \\AMM\\ P4        & (-30,50)      &   4.8 ($\\pm$0.2) & 15.9 ($\\pm$2.2) & 16.4 & 37.4 ($\\pm$ 4.9)  &  3.7 & 10.1 &147  & 16 & 960   & 6.5    &   \\\\ G28.34$+$0.06 \\AMM\\ P5        & (50,40)       &   5.6 ($\\pm$0.3) & 13.6 ($\\pm$1.3) & 14.1 & 29.6 ($\\pm$ 3.8)  &  6.3 & 4.7  &528  & 24 & 1378  & 2.6    &     \\\\ G33.71$-$0.01 \\AMM\\ P1 $^{n}$ & (-40,-200)    &   5.6 ($\\pm$1.1) & 16.6 ($\\pm$2.3) & 17.2 & 23.0 ($\\pm$ 6.6)  &      &     &     & 52 & 7880  &        & 15.6 \\\\ G33.71$-$0.01 \\AMM\\ P2 $^{n}$ & (60,100)      &   5.3 ($\\pm$0.3) & 17.1 ($\\pm$2.0) & 17.8 & 29.0 ($\\pm$ 3.3)  &      &     &     & 60 & 4928  &        &      \\\\ G79.27$+$0.38 \\AMM\\ P1        & (-20,0)       &   4.8 ($\\pm$0.2) & 12.4 ($\\pm$0.9) & 12.8 & 20.5 ($\\pm$ 2.5)  & 14.2 & 1.4  & 84  & 32 & 155   & 1.8    & 0.4 \\\\ G79.27$+$0.38 \\AMM\\ P2        & (-140,40)     &   6.0 ($\\pm$0.7) & 11.2 ($\\pm$1.4) & 11.5 & 11.3 ($\\pm$ 2.7)  & 12.1 & 0.9  & 55  & 25 & 53    & 1.0    &     \\\\ G79.27$+$0.38 \\AMM\\ P3        & (-250,70)     &   3.3 ($\\pm$2.2) & 11.7 ($\\pm$1.6) & 12.0 & 10.0 ($\\pm$ 25.0) &  6.5 & 1.5  & 19  & 20 & 40    & 2.1    &         \\\\ G79.34$+$0.33 \\AMM\\ P1        & (0,0)         &   5.8 ($\\pm$0.3) & 14.1 ($\\pm$0.8) & 14.6 & 11.9 ($\\pm$ 1.0)  & 17.5 & 0.7  & 143 & 39 & 178   & 1.2    & 0.2 \\\\ G81.50$+$0.14 $^{n}$          & (260,-375)    &   5.6 ($\\pm$1)   & 14.9 ($\\pm$1.3) & 15.4 &  7.6 ($\\pm$ 2.0)  &      &     &     & 78 & 570   &       & 0.1 \\\\ \\hline \\end{tabular} \\label{tab:physical_parameters} \\vfill  Notes: Columns are name, Offset Position from reference, excitation temperature, \\AMM\\ rotational temperature, kinetic temperature, \\AMM\\ column density,  $\\rm H_2$ column density, \\AMM\\ abundance, gas mass estimated from SCUBA data, size of the cores, virial masses from \\AMM\\, the alpha parameter defined as in total cloud mass from \\AMM\\. The formal errors are given in brackets. $^{n}$ are those sources for which the sizes have been derived from \\AMM\\ data. The $\\rm H_2$ column densities are estimated from 850 $\\mu$m SCUBA map \\citep{carey2000:irdc}).} \\end{minipage} \\end{sidewaystable*}    \\begin{table*} \\begin{center} \\caption{Mean Values of \\AMM\\ Core Samples.}\\label{tab:nh3_lowmass} \\vspace*{2mm} \\begin{tabular}{rddddd} \\hline \\hline Sample & \\dlabel{FWHM (\\kms)} & \\dlabel{$T_{\\rm kin}$(K)} & \\dlabel{size(pc)}& \\dlabel{$M_{\\rm NH_3}$ ($M_{\\odot}$)}& \\dlabel{$\\Sigma$} \\\\ \\hline Taurus&0.33  & 10& 0.06 & 1.1 & 0.08 \\\\ Perseus&0.55  &13 & 0.12 &  9 & 0.16 \\\\ Orion&1.12 & 17& 0.15 & 21 & 0.24 \\\\ IRDC cores& 1.7 &15& 0.57& 492 & 0.40 \\\\ \\hline \\end{tabular} \\end{center} Notes: Columns are the source sample, the mean FWHM of the \\AMM\\ (1,1) line, kinetic temperature, size of the core, mass as determined from \\AMM\\ assuming a fractional abundance of \\AMM\\ relative to $\\rm H_2$ and the mean surface density.  The mean values for dense cores in Taurus, Perseus and Orion have been taken from \\citet{ladd1994:nh3_perseus}.  \\end{table*}"
    },
    "0601/astro-ph0601552_arXiv.txt": {
        "abstract": "We report the first measurements of the absolute ionization yield of nuclear recoils in liquid xenon, as a function of energy and electric-field. Independent experiments were carried out with two dual-phase time projection chamber prototypes, developed for the XENON Dark Matter project.  We find that the charge yield increases with decreasing recoil energy, and exhibits only a weak field dependence. These results are a first demonstration of the capability of  dual phase xenon detectors  to discriminate between electron and nuclear recoils, a key requirement for a sensitive dark matter search at recoil energies down to 20~keV. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601287_arXiv.txt": {
        "abstract": "From October 1999 to August 2002, 45 transient X-ray sources were detected in M31 by {\\it Chandra} and {\\it XMM-Newton}.  We have performed spectral analysis of all {\\it XMM-Newton} and {\\it Chandra} ACIS detections of these sources, as well as flux measurements of {\\it Chandra} HRC detections.  The result is absorption-corrected X-ray lightcurves for these sources covering this 2.8 year period, along with spectral parameters for several epochs of the outbursts of most of the transient sources.  We supply a catalog of the locations, outburst dates, peak observed luminosities, decay time estimates, and spectral properties of the transient sources, and we discuss similarities with Galactic X-ray novae.  Duty cycle estimates are possible for 8 of the transients and range from 40\\% to 2\\%; upper limits to the duty cycles are estimated for an additional 15 transients and cover a similar range.  We find 5 transients which have rapid decay times and may be ultra-compact X-ray binaries.  Spectra of three of the transients suggest they may be faint Galactic foreground sources.  If even one is a foreground source, this suggests a surface density of faint transient X-ray sources of $\\gap$1 deg$^{-2}$. ",
        "introduction": "Bright X-ray transient (XRT) events provide brief glimpses into the physical state of accreting compact objects.  Detailed studies of Galactic soft XRTs, also known as X-ray Novae (XRNe), have put constraints on the physics at work in these systems (e.g. \\citealp{chen1997}).  Optical follow-up of 15 Galactic XRNe has revealed that the systems each contain an accreting object more massive than $\\sim 3$ M$_{\\odot}$ \\citep{mcclintock2003}, which cannot be a stable neutron star (NS).  Their high X-ray luminosities ($>10^{38}$\\ergs\\/) and millisecond X-ray variability further support the argument that these binary systems contain black holes (BHs). Therefore, these objects are among the most secure examples of black holes known \\citep{charles1998}.  Galactic XRTs occur as both high-mass X-ray binaries (HMXBs) and low-mass X-ray binaries (LMXBs). The secondary stars of outbursting HMXBs are typically B/Be type, and the primaries are typically neutron stars \\citep{tanaka1996}. These HMXBs tend to exhibit pulsations, hard spectra, and periodically repeating outbursts.  At the same time, outbursting LMXBs have often been found to contain black hole (BH) primaries.  As black holes do not have a physical surface, X-rays from the black hole X-ray novae (BHXNe) do not pulse. In addition, the spectra of BHXNe are typically softer than those of HMXBs (see the review by \\citealp{mcclintock2003}).  Outside the Milky Way, nearby galaxies are observed to have XRTs which are likely to be similar to those in our Galaxy.  However, their greater distance makes it difficult to find optical counterparts and hinders efforts to decipher detailed information about their X-ray properties.  On the other hand, there are some advantages to extragalactic studies of XRNe.  For example, in the case of M31, all of the bulge can be observed in one Chandra or XMM observation, and each such observation often reveals a new XRT source (e.g. \\citealp{williams2004}).  These observations provide accurate positions for followup observations as well as spectral information for constraining the physical nature of the sources during these accretion events.  Because M31 is the nearest neighboring spiral galaxy, transient X-ray sources detected in M31 have been well-reported.  The earliest transient source was discovered by \\citet{trinchieri1991}.  Many others followed, as observations of M31 by {\\it Chandra} and {\\it XMM-Newton} have been combed in order to catalog events that resemble Galactic XRNe.  \\citet{white1995} reported a supersoft repeating transient source. \\citet{atel79} reported two more new transient sources.  Five additional sources were seen by \\citet{trudolyubov2001} and \\citet{osborne2001} in {\\it XMM-Newton} and {\\it Chandra} observations from the year 2000.  \\citet{iauc7659} reported another. \\citet{kong2002} cataloged 10 new X-ray transients in their Chandra-ACIS study of the central region of M31.  \\citet{iauc7798} reported 3 more transient events from {\\it XMM-Newton} data. Recently, \\citet{williams2004} cataloged 8 additional transient sources from a {\\it Chandra}-HRC survey.  Finally, \\citet{distefano2004} have completed a survey for supersoft X-ray sources (SSS) in four regions of M31, finding 12 previously unpublished transient objects and noting 6 objects in the $ROSAT$ catalog that were undetected by $Chandra$ or $XMM$.  Herein, we identify 2 previously-detected objects as transient candidates.  All of these reports add up to 51 X-ray transients, 45 of which were detected by {\\it XMM-Newton} or {\\it Chandra} over a 3 year period.  These 45 are presented in Table~\\ref{trans.tab}.  These transient sources in M31 offer the possibility of studying a statistical sample of sources all at a well-known distance (780 kpc; \\citealp{williams2003sfh}) and covering only a small part of the sky. In addition, while portions of our Galaxy cannot be effectively surveyed in low-energy X-rays due to absorption, the M31 sample does not suffer from this problem.  The M31 sample allows comparative studies with the Galactic sample and adds to the number of transient sources known.  Finally, M31 could harbor X-ray transient sources with uncommon properties not yet discovered in the Galaxy.  Only through detailed studies of the extragalactic siblings of Galactic XRNe will be able to obtain a complete understanding of this class of X-ray sources.   In order to fully exploit the available X-ray data for M31 transients that were detected in 1999--2002, we have measured the fluxes and, where possible, the spectra of every {\\it XMM-Newton} and {\\it Chandra} detection of these 45 transient X-ray sources. Section~\\ref{data} describes the data reduction technique for detecting and extracting the sources, measuring their fluxes and fitting their spectra.  Section~\\ref{results} details our results, including lightcurves, spectral parameters, and classification of the different types of spectra displayed by the transient sources. Section~\\ref{discussion} compares the properties of these events with Galactic XRNe, and finally Section~\\ref{conclusions} summarizes our conclusions. ",
        "conclusions": "We have put together all of the {\\it Chandra} and {\\it XMM-Newton} data sets pertaining to 45 transient X-ray sources in M31 detected from October 1999 to August 2002.  These data have provided a catalog of positions, spectral properties, and lightcurves for transient X-ray sources in M31.  The locations of these transient sources indicate that they most often form near the center of the galaxy, suggesting that most of them are old LMXBs.  The distribution of peak luminosities of these sources range from $\\sim$10$^{36}$ to $\\sim$10$^{39}$ \\ergs , similar to Galactic XRTs.  The estimated $e$-folding decay times are similar to those seen in Galactic XRNe, also suggesting that many of these sources are BH-containing LMXBs.  The transient source lightcurves and positions indicate 5 potential UCXBs, 3 of which are located in the M31 bulge.  Two other sources have possible periodic outbursts and hard spectra consistent with HMXBs, but they have no optical counterparts brighter than M$_V \\approx -1$.  The lightcurves also constrain the transient rate in M31.  If we take only those sources with the best lightcurves (Hi/Lo $>$15, measured decay time, and one at least one clear outburst), there were 26 transient sources seen in the 35 months covered by our sample. Removing the total of $\\sim$9 months during which M31 was not observable during this time period yields a transient rate in M31 of $\\sim$1 per month, consistent with the previous result in \\citet{williams2004}.  The rate of X-ray novae in the Galaxy ($\\sim$2.6 per year; \\citealp{chen1997}), hinting that the rate in M31 could be higher.  However, since the Galactic rate is likely incomplete because sampling the entire Galaxy is much more difficult than sampling M31, the similarity of the two rates is noteworthy.  Furthermore, these data are some of the best available for studying the duty cycles of a large sample of transient sources.  We have used the lightcurves to constrain the duty cycles of the transients, finding about half have duty cycles $<$0.1 and all but 3 sources have duty cycles $<$0.2.  This range is similar to that seen in the Galaxy \\citep{chen1997}, where the transients have duty cycles ranging from 0.0005--0.3 and most duty cycles are $<$0.1.  The spectra of the sources provided more details about their physical properties.  Some source spectra are well-fitted by simple absorbed power-law models with photon indexes ranging from 1--4 and absorption values from 0.1--10 $\\times$ 10$^{21}$ cm$^{-2}$.  Others are only well-fitted by multi-component spectral models, including power-law plus blackbody or disk blackbody models with temperatures 0.03--1.2 keV.  Still other source spectra are too soft to provide an acceptable fit with any of these model combinations.  Three soft sources are well-fit by a neutron-star atmosphere model, but the best-fit effective temperatures indicate that these sources must be in the Galactic foreground.  The surface density of these sources ($\\gap$1 deg$^{-2}$) is consistent with that of magnetic CVs (polars) but higher than expected for neutron stars, suggesting that these sources are foreground polars. Their spectra are also well-fit by multi-component models similar to those of some polars, showing that polar spectra and neutron star atmosphere spectra can be similar.  Clearly the discovery of these XRTs marks only the beginning of their study.  These sources provide the largest extra-galactic sample of XRTs available in one galaxy and will open the door to future comparison studies of outbursting X-ray binary populations in galaxies.  We thank Robin Barnard for suggesting the NSA model for fitting the spectra of some of the soft sources.  Support for this work was provided by NASA through grant GO-9087 from the Space Telescope Science Institute and through grant GO-3103X from the {\\it Chandra} X-Ray Center."
    },
    "0601/astro-ph0601064.txt": {
        "abstract": "We report the discovery of a small cluster of massive stars embedded in a NIR nebula in the %%@ direction of the IRAS 15411-5352 point source, which is related to the alleged planetary nebula %%@ W16-185. The majority of the stars present large NIR excess characteristic of young stellar %%@ objects and have bright counterparts in the Spitzer IRAC images; the most luminous star (IRS1) %%@ is the NIR counterpart of the IRAS source. We found  very strong unresolved Br$\\gamma$ emission %%@ at the IRS1 position and more diluted and extended emission across the continuum nebula. From %%@ the  sizes and  electron volume densities  we concluded that they represent ultracompact and %%@ compact HII regions, respectively. Comparing the Br$\\gamma$ emission with the 7 mm free-free %%@ emission, we estimated that the visual extinction ranges between 14 and 20 mag. We found that %%@ only one star (IRS1) can provide the number of UV photons necessary to ionize the nebula. ",
        "introduction": "W16-185 was classified as a planetary nebula by \\citet{wray66} but because of its proximity to %%@ the IRAS point source IRAS15411-5352, \\citet{acker87} concluded that it is a compact HII region. Further investigation by \\citet{nouma93}, through low-dispersion spectroscopy and narrow band %%@ CCD imaging in the H$\\alpha$ line and the adjacent continuum, showed that the monochromatic %%@ H$\\alpha$ image has its centroid displaced by $\\Delta\\alpha=11^{\\prime\\prime}$ and %%@ $\\Delta\\delta=-13^{\\prime\\prime}$ from the IRAS source  although it lies within its error %%@ ellipse. Even so, since  the object is located in the border between compact HII regions and PNe in the %%@ far infrared color-color diagram \\citep{pott88}, they suggested that it was better described as %%@ a low-excitation planetary nebula.  Although there is no point source in the 5 GHz PMN catalog \\citep{wright94} in the direction of %%@ W16-185, the region can be considered part of the G326.7+0.6 radio source, associated with the %%@ RCW95 complex (Rodgers, Campbell \\& Whiteoak 1960). \\citet{goss70} mapped this region at 5GHz %%@ and found the maximum emission in the direction of the IRAS15408-5356 source and  an unresolved %%@ extension in the direction of W16-185. Strong CS(2-1) and NH$_3$(1,1) line emission were %%@ detected by Bronfman, Nyman \\& May (1996) and  by \\citet{roman02} for both IRAS sources. The %%@ radial velocities of the CS and NH$_3$ lines differ by less than 2 km s$^{-1}$ between the two %%@ regions, suggesting that they are at the same heliocentric distance.  \\citet{roman04}  reported the existence of a rich cluster of OB stars centered at %%@ IRAS15408-5356. In this paper we report the detection of a  newly formed cluster of stars embedded in a small %%@ NIR nebula, associated to the IRAS15411-5352 source. We observed the region in the $J$, $H$ and $K$ filters, as well as in the integrated Br$\\gamma$ %%@ hydrogen line. The continuum subtracted Br$\\gamma$ image shows two distinct regions, an unresolved ultra %%@ compact HII region (UCHII) and a more extended $10^{\\prime\\prime}$ compact HII region (CHII). Using the Itapetinga radiotelescope at 43 GHz, with $2^\\prime$.2 resolution, we were able to %%@ separate the radio source at the position of W16-185 from G326.7+0.6.  In Section 2 we describe the LNA and radio observations as well as the Spitzer/IRAC data. In %%@ Section 3 we identify the stellar cluster, the IRAS source and discuss the continuum and %%@ Br$\\gamma$ emission from the nebula. In Section 4 we summarize our results. ",
        "conclusions": "We found that the ionized region classified as the planetary nebula W16-185 is in reality %%@ powered by a cluster of young and  massive stars, still embedded in the parental molecular %%@ cloud. From the NIR photometry we identified twenty cluster member candidates; the majority of %%@ them presenting large NIR excess, characteristic of very young objects.  Using Spitzer infrared images we found that all of selected sources presenting large $H-K$ %%@ colors have bright Spitzer counterparts. From  those images we noted the presence of at least %%@ six highly absorbed mid infrared sources, southeast of the IRS1 source. Only one of them was %%@ detected in the $K$ band of our survey, implying that the extinction must be very high in that %%@ direction.  The IRS1 source is the most luminous object in the region and also presents large infrared %%@ excess. This source is the NIR counterpart of the  MSX source G326.7238+00.6148, which is the %%@ MIR counterpart of the IRAS15411-5352 source. From its near to far spectral energy distribution %%@ and assuming that all the stellar luminosity is used to heat the dust, we determined a lower %%@ limit for its bolometric luminosity of $L_{Bol}\\geq 7.7\\times 10^4 L_\\odot$, which corresponds %%@ to an O8.5V ZAMS star.  We found  very strong Br$\\gamma$ emission at the IRAS source position and more diluted and %%@ extended emission across the continuum nebula. The first one is very compact and was not %%@ resolved by our observations, which corresponds to a diameter of less than 0.012 pc at the %%@ distance of 2.4 kpc, size characteristic of ultracompact HII regions. The size of the extended %%@ region is $10^{\\prime\\prime}$, corresponding to a linear size of 0.11 pc, typical of  compact %%@ HII regions  (Churchwell 2002).  From the observed Br${\\gamma}$ flux we computed the emission measure $E$ and the expected %%@ free-free emission at radio wavelengths. From the measured 7 mm flux density, we found lower and %%@ upper limits for the extinction $A_V$ of 14 and 20 mag, respectively.  From the absorption corrected emission measures, we obtained limits to the volume densities of %%@ $(7.4-9.8)\\times 10^3$ and $(1.4-1.9)\\times 10^5$ cm$^{-3}$ for the extended and compact %%@ regions, respectively. These density values agree with the assumption of the compact and %%@ ultracompact nature of the sources, as inferred earlier from their sizes. Moreover, the relation %%@ between volume density and size agrees very well with what is found for other compact regions %%@ (eg. Kim \\& Koo, 2001).   We also obtained upper and lower limits for the number of Lyman continuum photons necessary to %%@ produce the Br$\\gamma$ emission in both the compact and extended region, obtaining $2.2 \\times %%@ 10^{48} < N_{Ly} < 3.9 \\times 10^{48}$ and $1.6 \\times 10^{48} < N_{Ly} < 2.8 \\times 10^{48}$ %%@ photons s$^{-1}$, respectively. The spectral type of the single ionizing star in the compact %%@ region agrees with that  calculated from the near to far infrared spectral energy distribution. %%@ However, the  stars embedded in the extended region  do not provide enough UV photons to ionize %%@ it, suggesting that IRS1  is probably the main ionizing source of the whole region. In fact, %%@ from the shape of the observed NIR nebula and from the relative position of IRS1 source, we %%@ suggest that it could be produced by the champagne flow as was found for other ultracompact %%@ cores.   %% The \\notetoeditor{TEXT} command allows the author to communicate %% information to the copy editor.  This information will appear as a %% footnote on the printed copy for the manuscript style file.  Nothing will %% appear on the printed copy if the preprint or %% preprint2 style files are used.  %% The eqnarray environment produces multi-line display math. The end of %% each line is marked with a \\\\. Lines will be numbered unless the \\\\ %% is preceded by a \\nonumber command. %% Alignment points are marked by ampersands (&). There should be two %% ampersands (&) per line.  %% If you wish to include an acknowledgments section in your paper, %% separate it off from the body of the text using the"
    },
    "0601/astro-ph0601002_arXiv.txt": {
        "abstract": "Theoretical studies and current observations of the high-redshift intergalactic medium (IGM) indicate that at least two cosmic transitions occur by the time the universe reaches gas metallicities of about $10^{-3} Z_\\odot$.  These are the cosmological reionization of the IGM, and the transition from a primordial to present-day mode of star formation. We quantify this relation through new calculations of the ionizing radiation produced in association with the elements carbon, oxygen and silicon observed in Galactic metal-poor halo stars, which are likely second-generation objects formed in the wake of primordial supernovae. We demonstrate that sufficient ionizing photons per baryon are created by enrichment levels of [Fe/H] $\\sim$ -3 in the environment of metal-poor halo stars to provide the optical depth in the cosmic microwave background of $\\sim$ 0.1 detected by $WMAP$. We show, on a star by star basis, that a genuine cosmic milestone in IGM ionization and star formation mode occurred at metallicities of $10^{-4}$--$10^{-3} Z_\\odot$ in these halo stars. This provides an important link in the chain of evidence for metal-free first stars having dominated the process of reionization by $z \\sim$ 6.  We conclude that many of the Fe-poor halo stars formed close to the end of or soon after cosmological reionization, making them the ideal probe of the physical conditions under which the transition from first- to second-generation star formation happened in primordial galaxies. ",
        "introduction": "Theoretical studies of the early universe predict the existence of first-generation stars which lack metals and consequently have unique physical structures and properties \\citep{bkl, tsv, sch02}.  Such objects have not, however, been detected to date in either the high-redshift or local universe. Therefore, indirect inferences from modelling the impact of primordial stars and supernovae (SNe) on their environment through radiative and chemical feedback remain important. The reionization of the intergalactic medium (IGM) and the widespread metal enrichment of high-$z$ galaxies, QSOs and the IGM provide strong observational constraints for such theoretical studies. A particularly effective tool in constraining the masses and properties of the first stars has been to combine their ionizing photon and nucleosynthetic signatures \\citep{venktruran, tvs04}.  When we examine the current data and the results of theoretical work, an interesting coincidence becomes apparent: that three important cosmic transitions occur at gas metallicities of about $10^{-5}$--$10^{-3} Z_\\odot$. To begin with, the universe first generates sufficient ionizing photons for hydrogen reionization when the IGM reaches these metallicities \\citep{mirrees97,venktruran}.  Second, theoretical studies of the cooling and fragmentation of gas in early galaxies reveal that the primordial stellar initial mass function (IMF) evolves from its metal-free to present-day form at gas transition metallicities of $Z_{\\rm tr} \\sim 10^{-5}$--$10^{-3} Z_\\odot$ \\citep{brommetal01,schneider02,omukai05, santoroshull}. This transition results from the physics of gas cooling in early galaxies, when the dominant cooling channel switches from molecular hydrogen to metal lines or dust grains, with the latter pathway dominating in systems with $Z \\la 10^{-4} Z_\\odot$ \\citep{omukai05}.  Third, there are indications that the relative abundances of CNO and alpha-elements in metal-poor stars in the Galactic halo have increased scatter and altered ratios when the iron content of these stars falls below $\\sim 10^{-3}$ of solar values (\\citealt{barklem05,cayrel}; \\citealt{beersaraa05}, and references therein). These subsolar mass objects are thought to be second-generation stars that formed from metal-poor fragmenting gas in the shells of the very first SNe in early galaxies \\citep{salvaterra04}.  Such stars contain the fossil imprint of nucleosynthesis from the preceding generations of metal-free stars, reflecting the ``prompt inventory'' of metals \\citep{qianwass02,wadavenk}.   When these three points are taken together, we see that several cosmological milestones occured by metallicities of $\\sim 10^{-3} Z_\\odot$, a relation that must result from the strong feedback between the buildup of ionizing radiation and the formation of second-generation stars.  In this Letter, we quantitatively demonstrate a novel connection in the environment of extremely metal-poor (EMP) halo stars between two of these relationships: that the stellar population that was important for early nucleosynthesis is also responsible for cosmological reionization. We define EMP stars as those with [Fe/H] $\\la$ -3 \\citep{beersaraa05}. We examine the significance of the second transition in relation to the other two in a subsequent work, where we plan to use numerical simulations to investigate the spatially varying feedback between gas cooling and the propagation of ionizing radiation in the first galaxies. ",
        "conclusions": "For our calculations here, we take the values of $\\eta_{\\rm Lyc}$ for carbon and oxygen computed in \\citet{venktruran}, and compute them independently for silicon. We find that for the 1--100 $M_\\odot$ Salpeter-slope IMF, $\\eta_{\\rm Lyc}$ has values of 0.48, 0.2 and 1.34 for C, O and Si respectively; for a VMS IMF, the corresponding values are 0.098, 0.01, and 0.057. Using these numbers, we show the results of our calculations in Figure 2, where the number of ionizing photons created per baryon in EMP stars, $N_{\\gamma}/N_{\\rm b}$, is shown as a function of [Fe/H] for the elements C, O and Si in three separate panels. In each case, the ionizing photon contribution from a present-day and a VMS IMF are plotted.  A number of trends seen in Figure 2 are of relevance to the first-stars field. First, VMS make a low but persistent contribution to the overall ionizing photon budget in association with the metals locked in EMP stars, typically an order of magnitude less than the 1--100 $M_\\odot$ IMF. This is because VMSs are in general less efficient at generating total ionizing radiation per unit metal yield relative to other metal-free stellar populations, owing to the large metal production from VMSs.  Second, the slow buildup of metals and the associated ionizing radiation is apparent, for each element and for each IMF. This trend appears most obvious for oxygen, although there is a significant scatter in the correlation between $N_{\\gamma}/N_{\\rm b}$ and [Fe/H] at any given [Fe/H]. As we stated earlier, the notable exceptions are the two stars at the lowest values of [Fe/H].  Third, we show in each panel the approximate criterion to generate an optical depth in the cosmic microwave background (CMB) of about 0.1, consistent with data from the {\\it Wilkinson Microwave Anisotropy Probe} ($WMAP$; \\citealt{spergel06}), of about $N_\\gamma/N_{\\rm b} \\sim$ 34,000 \\citep{tum06,tvs04}.  This number is derived from interpolating between cases studied by \\citet{tvs04} where the lifetime-integrated ionizing photon production from various stellar populations and IMFs were calculated and used as inputs in detailed cosmological reionization models. Interestingly, the EMP data indicates that this line is crossed at least once {\\it and perhaps twice} between -4 $\\leq$ [Fe/H] $\\leq$ -3 by a present-day IMF. This is best seen in the panel corresponding to Si, where a clear dip in $N_\\gamma/N_{\\rm b}$ occurs between -3.6 $\\leq$ [Fe/H] $\\leq$ -3. The VMS IMF appears to never meet this criterion over the metallicity ranges considered here, consistent with the role of VMSs in IGM reionization predicted by \\citet{venktruran}.  \\\\  \\centerline{\\epsfxsize=1.0\\hsize{\\epsfbox{f2.eps}}}  \\vspace{0.1in} \\figcaption{The minimum number of ionizing photons per baryon in stars that is created in association with specific metals observed in MP halo stars. Top, middle and bottom panels correspond to the elements C, O, and Si. The thick dashed black line in each panel corresponds to the value required to roughly match the $WMAP$ data for an optical depth of 10\\% in the CMB. All cases are for metal-free stars; each panel displays values for the 1--100 $M_\\odot$ (diamond points) and VMS (triangle points) IMFs considered in this work. See text for discussion.} \\vspace{0.2in}   Conducting a metal census of the universe from EMP stars and from the IGM is important in both environments, despite their differing systematics. The IGM contains nearly all of the baryons at high redshifts but the former case permits a detailed star-by-star analysis of the cosmic enrichment history.  By working with nucleosynthetic data in EMP stars rather than the IGM, we have not assigned a global fixed metallicity with the appropriately scaled individual element values according to their solar ratios. Rather, we have derived parallel constraints for reionization from a number of elements independently, which is a considerably stronger approach. In addition, observations of EMP stars relative to IGM absorbers provide significantly larger element coverage in the well-constrained environment of stellar atmospheres, a situation that will greatly improve in the near future.   Although we naively apply the same factor for ionization efficiency for [Fe/H] $\\ga$ -3, $\\eta_{\\rm Lyc}$ is likely to be closer to the values associated with metal-poor stars, typically a factor of a few to 10 lower, owing to their increased metal yield. Depending on the individual metal abundance, this could lead to a less steep rise or flattening in the behavior of $N_{\\gamma}/N_{\\rm b}$ at [Fe/H] $\\ga$ -3. This prediction is somewhat more speculative, as the integrated contribution from many subsequent stellar generations of varying metallicities will be difficult to distinguish in this environment.   The notable exceptions to most of the trends noted above are the two most Fe-poor stars at [Fe/H] $\\sim$ -5.5. At such low metallicities, one would expect that these stars reflect enrichment from a single SN \\citep{tum06}. These stars are, however, highly over-enriched in several elements including C, O and Si, to the degree that they are actually not metal-poor, but rather at metallicities of nearly a tenth solar. How these stars came to acquire such skewed element abundances relative to solar ratios remains puzzling at present.  Some possibilities include pollution from a binary companion, from a subluminous hypernova with fine-tuned ratios of mixing and fallback \\citep{iwamoto05}, or the effects of stellar rotation and ensuing mass loss which could dramatically boost the production of CNO elements \\citep{chiappini06}. The ratio of these elements to iron or relative to each other is sometimes used as a chronometer in other astrophysical environments. For the most Fe-poor stars, however, it is clear that ratios such as [Fe/O] or [C/O] are not easily related to time. These should normally increase with age of a stellar population as the stars of progressively lower mass evolve off the main sequence, but as we saw in Figure 1 for EMP stars, [C/O] achieves its highest values at the lowest [Fe/H]. Perhaps these stars were created in atypical environments as the IGM approached reionization, as seen from the generation of ionizing radiation in Figure 2 -- the most Fe-poor stars appear to have formed at epochs when, for a metal-free present-day IMF, $N_{\\gamma}/N_{\\rm b}$ nearly equals or exceeds the criterion from $WMAP$, at least for C. One interesting explanation of this, in addition to those above, could be that the most Fe-poor halo stars are also true second-generation stars: ones that formed through the dust grain cooling pathway ($Z_{\\rm tr} \\sim 10^{-5} Z_\\odot$) rather than through metal line cooling using C, O or Si ($Z_{\\rm tr} \\sim 10^{-3} Z_\\odot$). The unusual element trends in these stars could then be explained through the element segregation predicted in a scenario involving dust transport in early galaxies. Here, conditions in primordial SN remnants result in the radiative transport of dust in the hot radiation field from metal-free stars, leading to selective element enhancement from compounds such as graphites and silicates in the cooling gas in SN shells \\citep{vns06}, which are likely the sites of EMP star formation.   What do the results in this paper imply for cosmology? Those EMP stars which have observational age estimates are about 12--13 Gyr old \\citep{snedenetal03,beersaraa05}. In the context of our results, it appears that most of the EMP stars formed close to the end of or soon after cosmological reionization.  That reionization occurs in the first Gyr of the universe (by $z \\sim$ 6) is no surprise, as this is already clear from Gunn-Peterson studies of the high-$z$ IGM \\citep{becker} and from current CMB data. A more interesting result of translating the EMP star abundances in C, O and Si into the accompanying production of ionizing photons per baryon in stars is the clear buildup of ionizing radiation until reionization is accomplished at [Fe/H] $\\sim$ -3, [Si/H] $\\sim$ -3.5 and [O/H] $\\sim$ -2.5 ([C/H] ranges from -4 to -2 at these Fe-metallicities). At this point, the initial stellar population (presumably metal-free) must turn off, as seen from the altered element abundance patterns at higher [Fe/H], leading to a new pattern of nucleosynthesis from different SNe. \\citet{cayrel} argue for the presence of a primordial SN population up to [Fe/H] or [Mg/H] values of about -3, with a combination of primordial and nonprimordial SNe at higher metallicities. Such a truncation in the mode of star formation is most likely attributed to the critical metallicity being reached in the gas in which EMP stars formed, signalling an end to the first-stars epoch. Given certain conditions of density and temperature, $Z_{\\rm tr}$ can be higher in iron than in carbon, oxygen, or silicon. This is not consistent, however, with the overall trend of enhanced C, O and Si at {\\it low} [Fe/H].  A naive projection of the above metallicities in Fe, Si, O and C at which reionization ``occurs'' in Figure 2 implies that EMP stars should have formed in gas of number densities $\\ga 10^3$ cm$^{-3}$ \\citep{santoroshull}.   A more puzzling coincidence is why the end of the first-stars phase should occur at about the epochs corresponding to cosmological reionization. This points towards the possibility of strong feedback between intrahalo metal reincorporation timescales and halo-IGM transport of ionizing radiation. We return to this issue in a subsequent work, where we will quantitatively relate the reionization epoch to the transition metallicity in early galaxies. Clearly, EMP stars provide an important window to study the transition from first- to second-generation star formation and the physical conditions at which this occurs in primordial galaxies. The novel result derived here is that the abundances in such stars show a strong relation to cosmological reionization. Targetted observations are required with more uniform element coverage of EMP stars, particularly the C-rich stars, in the critical range of -6 $\\la$ [Fe/H] $\\la$ -4. This range currently has very little data, and will be useful in closing the loop in our understanding of these important cosmic events at high redshifts."
    },
    "0601/astro-ph0601234_arXiv.txt": {
        "abstract": "{  We present results from a comparative study of XMM-Newton observations of four classical T Tauri stars (CTTS), namely \\object{BP Tau}, \\object{CR Cha}, \\object{SU Aur} and \\object{TW Hya}. In these objects coronal, i.e. magnetic, activity and as recently shown, magnetically funneled accretion are the processes likely to be responsible for the generation of X-ray emission. Variable X-ray emission with luminosities in the order of $10^{30}$\\,erg/s is observed for all targets. We investigate light curves as well as medium and high-resolution X-ray spectra to determine the plasma properties of the sample CTTS and to study the origin of their X-ray emission and its variability. The emission measure distributions and observed temperatures differ significantly and the targets are dominated either by plasma at high densities as produced by accretion shocks or by predominantly hotter plasma of coronal origin. Likewise the variability of the X-ray luminosity is found to be generated by both mechanisms. Cool plasma at high densities is found in all stars with detected \\ion{O}{vii} triplet emission, prevented only for SU~Aur due to strong absorption. A general trend is present in the abundance pattern, with neon being at solar value or enhanced while oxygen, iron and most other metals are depleted, pointing to the presence of the inverse FIP effect in active coronae and possibly grain formation in evolved disks. We find that both accretion shocks and coronal activity contribute to the observed X-ray emission of the targets. While coronal activity is the dominant source of X-ray activity in the majority of the CTTS, the fraction for each process differs significantly between the individual objects. ",
        "introduction": "\\label{intro}  Young, late-type pre main-sequence stars, known as T~Tauri stars, are copiously found in or near star forming regions.  Historically, they are classified according to their H$\\alpha$ equivalent width as classical T~Tauri stars (CTTS) with EW$>$10\\AA\\,\\,or otherwise as weak line T~Tauri stars (WTTS). The principal physical difference between these two classes is that CTTS are thought to be in an earlier evolutionary stage, i.e. they still possess a disk containing dust and gas and are accreting matter. WTTS are thought to be more evolved, to have mainly lost their disks and are approaching the main-sequence. While CTTS are generally younger with ages of a few up to $\\sim$\\,10\\,Myr and the overall fraction of CTTS in a given star-forming region decreases with age, both types of T~Tauri stars are commonly found in the same star-forming regions, indicating individual evolutionary time scales. Since the H$\\alpha$-line is often time-variable, a classification solely based on its strength is sometimes misleading and other or additional spectral features have been used to identify and classify CTTS more reliably. The existence of a disk containing a significant amount of matter and ongoing accretion onto the host star impose major differences between the two types of TTS, leading to different spectral properties which can be observed at different wavelengths. The effects of the warm and probably structured disk are for instance reflected in the different infrared designations of CTTS (Class II) and WTTS (Class III) due to their additional near-infrared emission, which is an important indicator for CTTS. Consequently, the different criteria and underlying data used for the identification of CTTS may introduce selection effect in the respective samples.  Both types of T~Tauri stars are well known strong and variable X-ray emitters and large numbers of these objects were detected with {\\it Einstein} \\citep{fei81,fei89} and ROSAT \\citep{fei93,neu95}. Their X-ray emission and the observed flaring was usually interpreted as a scaled-up version of coronal activity in analogy to cool main-sequence stars. This picture is still valid and the X-ray emission from WTTS is thought to originate from coronal activity. However, T Tauri stars are thought to be at least in their early stages fully convective, therefore the question arises about the underlying processes for the generation of the magnetic field and the influence of accretion on the stellar structure and corona. Further on, the evolutionary stages of young stars and their high energy emission is of major importance for the process of planetary formation via its effects on chemistry, dust and protoplanets. A comprehensive review of high energy processes and underlying physics in young stellar objects, summarising the results of the era prior to {\\it Chandra} and XMM-Newton, is presented by \\cite{fei99}.  The largest sample of pre main-sequence stars are found in the Orion Nebula Cloud. Using {\\it Chandra} ACIS observations, \\cite{fei03} analysed a sample of 500 PMS stars with known basic properties. An even deeper exposure, named the {\\it Chandra} Orion Ultradeep Project (COUP) was presented recently, see e.g. \\cite{pre05}, confirming and extending their results. The X-ray luminosity of CTTS in this sample was found to be correlated with bolometric luminosity. No strong correlation between $L_{\\rm X}$ and rotation as observed for main-sequence stars was found, contrary, $L_{\\rm X}$ correlates with $L_{\\rm bol}$ and stellar mass, indicating the presence of different, e.g. turbulent, dynamos or of supersaturation effects in solar-type dynamos. The question, whether the presence of an accretion disk influences the stellar X-ray emission is still under debate. While \\cite{fei03} found no intrinsic differences of X-ray activity related to the presence of a circumstellar disk, the results found by \\cite{pre05} indicate that accretion diminishes  $L_{\\rm X}$ on average and weaken the $L_{\\rm X}$/$L_{\\rm bol}$ correlation. A similar influence was found by \\cite{ste01} in the X-ray properties of young stars in the Taurus-Auriga complex observed with ROSAT PSPC. The COUP sources are rather consistent with the bulk of their X-ray emission being produced in typical coronae. In general, magnetic activity in CTTS may also involve the circumstellar material, e.g. via star-disk interaction, but the topology of their magnetic fields in virtually unknown; \\cite{fav05} present evidence for magnetic star-disk interaction in the analysis of larger flares on several COUP objects.  Moreover, recent high-resolution spectroscopy of CTTS with {\\it Chandra} and XMM-Newton, indicated accretion shocks at least as an additional mechanism for the generation of X-ray emission in CTTS. In this interpretation X-ray emission is generated by a magnetically funneled accretion stream falling onto small areas of the stellar surface, producing shocked high density plasma \\citep{shu94,cal98}. Since in magnetospheric accretion the matter falls from several stellar radii at nearly free-fall velocity, strong shocks and consequently hot plasma emitting in the X-ray regime are formed. This contribution is expected to be at energies below 0.5\\,keV and is distinguishable from cool coronal plasma by its density; therefore it can be recognised only through high resolution X-ray spectroscopy. It is thus not contradictory that the COUP sample based on ACIS data with medium resolution spectroscopy and poorer sensitivity at low energies finds accretion to play no significant role.  TW~Hya was the first example of a CTTS, where observational evidence for an accretion scenario could be found in X-ray data. Analysis using density sensitive lines of He-like triplets (e.g. \\ion{O}{vii}, \\ion{Ne}{ix}) indicated very high densities, exceeding by far the densities previously found in coronal plasma from any other star \\citep{twc,twx}. These low line ratios were also found in the spectra of BP~Tau \\citep{bptau}, again strongly supporting an accretion scenario. However, accretion is only capable to produce plasma with temperatures of a few MK and therefore contributes nearly exclusively to the low energy X-ray emission. The observed flaring and additional hard emission attributed to plasma with temperatures of several tens of MK in BP~Tau and other CTTS require the presence of an additional X-ray generating mechanism such as magnetic activity. In this work we use the term `coronal activity', but more complex phenomena like star-disk interactions cannot be ruled out. While TW~Hya like objects where accretion is thought to be the dominant source of the X-ray emission are probably rare, the BP Tau observation revealed that accretion as an additional X-ray generating mechanism may be more common. Another X-ray generating mechanism, shocks in strong Herbig-Haro outflows that may also produce a soft X-ray excess, was proposed for jet-driving CTTS (which resemble Class I protostars) like DG Tau \\citep{gue05}. We note that this model cannot explain the soft excess in the case of TW~Hya or the other sample CTTS.  We utilise XMM-Newton data from four X-ray bright CTTS to study the origin of their X-ray emission in accretion shocks vs. coronal activity. We specifically use medium and high resolution X-ray spectra and spectrally resolved X-ray light curves to investigate the physical properties of our target stars BP~Tau, CR~Cha, SU~Aur and TW~Hya and to check the relevance of these mechanisms for the generation of the X-ray emission and variability in these young stars. This work is the first comparative X-ray study of CTTS using high and medium resolution spectra so far, therefore it is complementary to large sample studies of a specific star forming region such as COUP. Moreover, the data of CR~Cha and partly of SU~Aur and BP~Tau is analyzed and presented here for the first time. Results derived from RGS spectra of BP~Tau and the simultaneous obtained UV data were presented in our previous letter \\citep{bptau}, while the data on TW~Hya first presented by \\cite{twx} is re-analysed for purposes of comparison in a manner identical to the analysis of all other sources.  The plan of our paper is as follows: In Sect.\\,\\ref{tar} we describe the individual targets and previous X-ray results, in Sect.\\,\\ref{obsana} the observations and the methods used for data analysis and in Sect.\\,\\ref{results} we present and discuss the results subdivided into different physical topics, followed by a summary and our conclusions in Sect.\\,\\ref{summ}. ",
        "conclusions": "\\label{summ}  We have presented the first comparative study using high and medium resolution X-ray spectra of classical T~Tauri stars observed with the new generation X-ray telescopes so far. The results derived from these data are complementary to the ones obtained from the large sample of COUP sources, since high-resolution spectroscopy and low-energy sensitivity additionally permits the investigation of the cool plasma components.  Using the XMM-Newton observations of BP~Tau, CR~Cha, SU~Aur and TW~Hya we determined X-ray properties of accreting young pre-main sequence stars. Two different mechanisms are likely to contribute to the production of X-ray emission in those objects, first coronal, i.e. magnetic activity, and second, magnetically funneled accretion. We investigate variability, global spectra and density sensitive lines to check for the presence and relative contribution of these two mechanisms for the individual objects. For this purpose we utilize the emission measure distribution and its changes, specific abundance pattern and density diagnostics. Interpreting very high densities as indicator for accretion shocks and a high temperature component as indicator for coronal activity, we derive the following conclusions which are likewise supported by the analysis of abundance analysis and spectral variability. In all targets where the \\ion{O}{vii}-triplet is observable density analysis leads to higher densities than observed in any pure coronal source. Additionally, a high temperature component is present in all targets. We therefore argue that both X-ray generating processes are present in our sample CTTS, but at very different levels and importance in individual objects.  The X-ray emission of BP~Tau is overall dominated by coronal activity and a stronger flare of again probably coronal origin is observed. Even in quasi-quiescent phases the emission measure distribution is dominated by medium and hot temperature plasma whereof significant amounts are present at temperatures around 20\\,MK. However, the analysis of density sensitive lines confirms that accretion is also present and actually dominates the cool plasma with temperatures of a few MK. CR~Cha is similar to BP~Tau but more strongly absorbed and the coronal plasma temperatures are slightly lower. It also appears to be dominated by the coronal contribution, likewise there is evidence for high density plasma at cool temperatures. SU~Aur is the by far brightest and with temperatures of at least up to 50\\,MK hottest X-ray source in our sample. A powerful and active corona is the essential contributor to its X-ray spectrum and frequent flaring is observed. Unfortunately the very strong absorption prevents definite conclusions about the cooler plasma component. TW~Hya is the other hand strongly dominated by accretion, but an additional coronal component is clearly detected. It is the prototype and still the outstanding and only example of an accretion dominated star. Most of its plasma is at cool temperatures typical for accretion spots and it exhibits by far the lowest \\ion{O}{vii} f/i ratio. Also the highest neon to oxygen ratio, which may be interpreted via depletion of grain forming element, is found for TW~Hya. This is plausible since TW~Hya is the oldest and most evolved object in our sample. In the other sample CTTS neon is also enhanced, but the abundance pattern is more reminiscent of the inverse FIP effect found in active stars.  Considering the global picture of stellar evolution towards the main sequence, we find that magnetic processes play a major role in high energy phenomena in all our sample CTTS. While the specific stellar evolution and its timescale might depend on e.g. spectral type, the generation of X-rays via accretion and coronal activity is apparently a common feature of CTTS in general."
    },
    "0601/astro-ph0601144_arXiv.txt": {
        "abstract": "Radiatively-driven transfer flow perpendicular to a luminous disk is examined in the subrelativistic regime of $(v/c)^1$, taking into account the gravity of the central object. The flow is assumed to be vertical, and the gas pressure is ignored, while internal heating is assumed to be proportional to the gas density. The basic equations were numerically solved as a function of the optical depth, and the flow velocity, the height, the radiative flux, and the radiation pressure were obtained for a given radius, an initial optical depth, and initial conditions at the flow base (disk ``inside''), whereas the mass-loss rate was determined as an eigenvalue of the boundary condition at the flow top (disk ``surface''). For sufficiently luminous cases, the flow resembles the case without gravity. For less-luminous cases, however, the flow velocity decreases, and the flow would be impossible due to the existence of gravity in the case that the radiative flux is sufficiently small. Application to a supercritical accretion disk with mass loss is briefly discussed. ",
        "introduction": "Mass outflow from a luminous disk is a clue to the formation mechanism of astrophysical jets and winds in the active objects. In particular, in a supercritical accretion disk, the mass-accretion rate highly exceeds the critical rate, the disk local luminosity exceeds the Eddington one, and mass loss from a disk surface driven by radiation pressure would take place (see Kato et al. 1998 for a review of accretion disks).  So far, radiatively driven outflow from a luminous disk has been extensively studied by many researchers (Bisnovatyi-Kogan, Blinnikov 1977; Katz 1980; Icke 1980; Melia, K\\\"onigl 1989; Misra, Melia 1993; Tajima, Fukue 1996, 1998; Watarai, Fukue 1999; Hirai, Fukue 2001; Fukue et al. 2001; Orihara, Fukue 2003), as on-axis jets (Icke 1989; Sikora et al. 1996; Renaud, Henri 1998; Luo, Protheroe 1999; Fukue 2005a), as outflows confined by a gaseous torus (Lynden-Bell 1978; Davidson, McCray 1980; Sikora, Wilson 1981; Fukue 1982), or as jets confined by the outer flow or corona (Sol et al. 1989; Marcowith et al. 1995; Fukue 1999), and by numerical simulations (Eggum et al. 1985, 1988). Line-driven outflows were also extensively examined (Drew, Verbunt 1985; Vitello, Shlosman 1988; Proga et al. 1998; Murray, Chiang 1995; Proga et al. 2000; Proga, Kallman 2004). In almost all of these studies, however, the luminous disk was treated as an external radiation source, and the radiation transfer in the flow was not solved.   Radiation transfer in the disk, on the other hand, was investigated in relation to the structure of a static disk atmosphere and the spectral energy density from the disk surface (e.g., Meyer, Meyer-Hofmeister 1982; Cannizzo, Wheeler 1984; Shaviv, Wehrse 1986; Adam et al. 1988; Hubeny 1990; Ross et al. 1992; Artemova et al. 1996; Hubeny, Hubeny 1997, 1998; Hubeny et al. 2000, 2001; Davis et al. 2005; Hui et al. 2005; see also Mineshige, Wood 1990). In these studies, however, the vertical movement and the mass loss were not considered.  Recently, mass outflow as well as radiation transfer have been examined for the first time in the subrelativistic and fully relativistic cases (Fukue 2005b, c). In these studies, the gravity of the central object and the gas pressure were both ignored as a first step. In this paper, we thus consider the radiatively driven vertical outflow -- {\\it moving photosphere} -- in a luminous flat disk within the framework of radiation transfer in the subrelativistic regime of $(v/c)^1$, while taking into account the gravity of the central object, although the gas pressure is ignored.  In the next section we describe the basic equations in the vertical direction. In section 3 we show our numerical examination of the radiative flow. In section 4 we briefly apply the present model to the case of a supercritial accretion disk. The final section is devoted to concluding remarks. ",
        "conclusions": "In this paper we have examined the radiative transfer flow in a luminous disk in the subrelativistic regime of $(v/c)^1$, while taking account of gravity of the central object. The flow is assumed to be vertical, and the gas pressure is ignored for simplicity, while internal heating is assumed to be proportional to the density. The basic equations are numerically solved as a function of the optical depth $\\tau$, and the flow velocity $v$, the height $z$, the radiative flux $F$, and the radiation pressure $P$ are obtained for a given radius $r$, the initial optical depth $\\tau_0$, and the initial coditions at the flow base (disk ``inside''), whereas the mass-loss rate $J$ is determined as an eigenvalue of the boundary condition at the flow top (disk ``surface''). For sufficiently luminous cases, the flow resembles the case without gravity. For less-luminous cases, however, the flow velocity decreases, and flow would be impossible due to the existence of gravity.  The radiative flow investigaed in the present paper must be a quite {\\it fundamental problem} for accretion-disk physics and astrophysical jet formation, although the present paper is only the second step (cf. Fukue 2005b), and there are many simplifications at the present stage.  For example, we have ignored the gas pressure. In general cases, where the gas pressure is considered, there usually appears sonic points (e.g., Fukue 2002), and the flow is accelerated from subsonic to supersonic. In this paper we consider a purely vertical flow, and the cross section of the flow is constant. If the cross section of the flow increases along the flow, the flow properties such as a transonic nature would be influenced.   Fully relativistic effects are also important. In such a fully relativistic case, where the flow speed is of the order of the speed of light, the boundary condition at the flow top should be carefully treated (cf. Fukue 2005c). The radiation drag becomes much more important.  There remain many problems to be solved.     \\vspace*{1pc}  This work has been supported in part by a Grant-in-Aid for the Scientific Research (15540235 J.F.) of the Ministry of Education, Culture, Sports, Science and Technology."
    },
    "0601/astro-ph0601372_arXiv.txt": {
        "abstract": "{This paper presents the results of a pilot redshift survey of 18 candidate compact groups from the distant DPOSS survey that extends the available surveys of compact groups of galaxies to $z \\sim 0.2$ in redshift, mainly Hickson Compact Groups and Southern Compact Groups.} {The goal of our survey was to confirm group membership via redshift information and to measure the characteristic parameters of a representative, albeit small, sample of DPOSS survey groups.} {Of the 18 candidates observed, seven are found to be indeed isolated compact groups, i.e. groups with 3 or more concordant members and with no neighbouring known cluster, while 7 are chance projection configurations on the sky. Three remaining candidates, despite having 3 or more concordant member galaxies, are located in the neighbourhood of known clusters, while another candidate turned out to be a dense sub-condensation within Abell 0952.} {The median redshift of our 7 confirmed groups is $z \\sim 0.12$, to be compared with a median redshift of 0.03 for the local sample of compact groups by Hickson.  The typical group size is $\\sim$ 50 kpc, and the median radial velocity dispersion is 167 \\kms, while typical crossing times range from 0.005 H$_{0}^{-1}$ to 0.03 H$_{0}^{-1}$ with a median value of 0.018 H$_{0}^{-1}$, all similar to the values usually found in the literature for such structures in the local universe. The average mass-to-light ratio for our groups, M/L$_{B}$, is 92{\\it h}, higher than the value found for nearby Hickson compact groups but lower than that found for loose groups.  Our results suggest that, once full redshift information for its members becomes available, the DPOSS sample will provide a reference sample to study the properties of compact groups beyond the local universe. }{} ",
        "introduction": "Compact groups (hereafter CGs) are well known systems, ever since the discovery by Edouard Stephan of the first one in 1877, at a time when we did not even know about the expansion of the universe and about the existence of other galaxies outside our own. These galactic systems drew attention due to their small angular scale; their sizes are comparable to the mean distance between their member galaxies. Physically, we naively classify these systems as having low mass, high projected density, and low velocity dispersion. Galaxy-galaxy interactions (e.g. close tidal encounters) and mergers are therefore likely to dominate their evolution. Until now, however, signs of interactions in CGs selected from the Hickson catalogue, the most widely studied catalogue of such objects (Hickson 1982), were found to be uncommon and traces of mergers to be rare (Zepf 1993). When present, these features are mainly found in spiral-dominated groups.\\newline Subsequent studies by Mendes de Oliveira et al (1994) showed that, while interactions (i.e. encounters which do not disrupt the galaxies) between galaxies in compact groups are quite common, mergers remain rare, to the level of 6$\\%$ of the group galaxies. Addtional studies by Proctor et al. (2004) showed that the stellar population of the early-type member galaxies of nearby, z $\\le$ 0.03, compact groups is quite old (Proctor et al, 2004). This led to doubt the existence of compact groups as physically bound system, but detection of diffuse intra-group matter in 75$\\%$ of HCGs (Ponman et al., 1996) confirmed that these are indeed bound, self-gravitating structures.  The question then arises about the origin and evolution of compact groups: how can they survive for so long?\\newline Studies of the environment of HCGs (de Carvalho et al. 1997; Ribeiro et al. 1998) have shown that compact groups can be divided into three categories, namely: real compact groups, systems composed of a core+halo structure, and sparse groups. Objects belonging to different categories have different surface density profiles, and we observe a propensity for higher activity level in lower velocity dispersion groups. These differences are also reflected in their X-ray emission: groups detected in X-ray have higher galaxy density than groups without detectable X-ray emission. Moreover early type galaxies are more centrally concentrated in X-ray emitting groups than in the non-emitting ones; and finally groups dominated by late type spirals do not show X-ray emission (Zabludoff \\& Mulchaey, 1998).  The different properties of the group categories identified by de Carvalho et al. (1994, 1997) might be interpreted as an evolutionary scheme in which the groups form in a looser concentration of galaxies, followed by a period of strong evolution with merging episodes. They then settle into a more quiescent phase and finally end up as isolated field ellipticals (Coziol et al., 2004). However, recent studies by Proctor et al. (2004) and Mendes de Oliveira et al. (2005) find that early type galaxies in compact groups are older than field galaxies of the same type and similar to cluster galaxies. This means that formation of early type galaxies in groups by galaxy-galaxy mergers must have happened a long time ago (of the order of a few Gyrs). On the other hand, compact groups have a very short crossing time, of the order of a few percent of the Hubble time, making it unclear how groups dominated by early type galaxies are observable today, because they should have disappeared a long time ago. A different possibility might be that compact groups are quite young configurations and that we are observing different stages of an ongoing merging process at the same time. If this is true, then we might wonder what a search for compact groups at increasing redshift would produce: would we find an higher number of interacting groups than at present time? Would we find an increase in the activity of the member galaxies? Would we find changes in the group's physical characteristics, like velocity dispersion, mass, radius, and crossing time?  To achieve these goals requires a complete sample of compact groups whose observational and statistical biases are well understood. A new sample of 459 compact group candidates with a median expected redshift of $\\sim$ 0.12, i.e. 10 times higher than the typical crossing time of present-day compact groups, has been selected with an automatic algorithm applied to the Digitized Second Palomar Observatory Sky Survey II (herafter DPOSS II) galaxy catalogues (Iovino et al. 2003 and de Carvalho et al.  2005). We refer the reader to these cited papers for detailed information on the sample, however, we report here the selection criteria for the sake of completeness: \\begin{itemize} \\item{} {\\it richness}: n$_{member}$ $\\ge$ 4 in the magnitude interval $\\Delta$mag$_{comp}$ = m$_{faintest}$-m$_{brightest}$, with the constraint $\\Delta$mag$_{comp}$ $\\le$ 2. Here n$_{member}$ is the number of member galaxies and m$_{brightest}$, m$_{faintest}$ are the magnitude of the brighest and the faintest group members respectively. \\item{} {\\it isolation}: R$_{isol}$ $\\ge$ 3R$_{gr}$, where R$_{isol}$ is the distance from the centre of the smallest circle encompassing all group galaxies to the nearest non-member galaxy within 0.5 magnitudes of the faintest group member. R$_{gr}$ is the radius of the smallest circle emcompassing all the galaxies of the group. This criterion avoids finding small aggregates of galaxies within a larger structure, e.g. a cluster. \\item{} {\\it compactness}: $\\mu_{gr}$ $<$ $\\mu_{limit}$, where $\\mu_{gr}$ is the mean surface brightness within the circle of radius R$_{gr}$ and $\\mu_{limit}$ = 24 mag arcsec$^{-2}$ in r band. For comparison, Hickson uses a $\\mu_{limit}$ = 26 mag arcsec$^{-2}$, which would have increased the contamination of the sample to 80$\\%$, against the current estimate of 27$\\%$. \\end{itemize} The magnitude difference criterion is considerably stricter than Hickson's ($\\Delta$mag$_{comp}$ $\\le$ 3), meaning that we have a lower contamination rate, 10$\\%$ against 27$\\%$, but also reduced completeness.\\newline  A first step in exploiting such a sample is to define via spectroscopic follow-up how many groups are indeed bound, selecting subsamples of groups with three/four galaxies sharing the same recession velocity.  In addition, the spectroscopic data allow us to estimate the dynamic characteristics of the groups and assess the level of activity in their galaxy population. Here we present the first results from a pilot study carried out at the 3.58m New Technology Telescope (NTT) at La Silla Observatory. In the next section we describe the observation and data reduction, in Sect. 3 we present our results and in Sect. 4 our findings. \\newline  Throughout the paper, a $\\Lambda$CDM cosmology ($\\Omega_{M}$=0.3; $\\Omega_{\\Lambda}$=0.7) and H$_{0}$=67 \\kms Mpc$^{-1}$ have been used.\\newline ",
        "conclusions": "Our pilot study of compact groups at intermediate redshift has shown that the confirmed candidates have very similar properties to those observed for Hickson Compact Groups and that no significant evolution can be detected up to z $\\sim$ 0.12. This finding agrees with models predicting an early formation of a massive, common halo, within which the individual galaxies form and evolve. Such models, however, are still unable to explain the low velocity dispersion, high activity level groups found in the nearby universe.\\newline Our results suggest that the DPOSS sample, once full redshift information for its members becomes available, will provide a reference sample for studying the properties of compact groups beyond the local universe."
    },
    "0601/astro-ph0601691_arXiv.txt": {
        "abstract": "% Current catalogues of open clusters are rather heterogeneous and incomplete list of clusters than true catalogues. Before there has been no attempts of automatic search for open clusters in huge photometric catalogues using homogeneous all-sky approach.  We have developed such a method based on extraction from catalogue and successive analysis of stellar densities using the colour-magnitude diagrams. Our algorithm finds density peaks and then verifies the significance of these peaks exploiting the fact that in real clusters only stars lying on the isochrone must show density peaks, in contrast with stars lying far from the isochrone on the CMD. In addition such procedure allows to determine the physical parameters: age, distance and color excess of open clusters.   Preliminary study of 150 sq. degrees in Galaxy anticenter region yielded several dozens of new clusters. The software developed will allow to build first homogeneous all-sky open cluster catalogue (it will include essentially refined data for known clusters as well) based on 2MASS data. It will be possible to apply automated pipeline on any multicolor catalogue, where the VO-compliant way of access (e.g. ADQL + SkyNode) is provided (SDSS, etc.). ",
        "introduction": "Here we present here several new methods which allow us to involve the data from existing huge stellar catalogues into search and analysis of star clusters. We also show some preliminary but very promising results of these methods.  Currently there are several compiled catalogues of star clusters, but they all have several significant disadvantages:  \\begin{itemize} \\item Tens of percents of their clusters do not have reasonable CMD measurements, so they are only unverified groups of stars \\item Some of them do not even have precise celestial coordinates measurements. \\item Even some rich clusters do not have important physical parameters measured (age, distance, color excess) \\item In general, all existing catalogues of open clusters are very heterogeneous and are just the compilations of data from different sources. Nobody has ever performed fully automatic homogeneous search of clusters in full-sky regime.  \\end{itemize} ",
        "conclusions": ""
    },
    "0601/astro-ph0601685_arXiv.txt": {
        "abstract": "We study whether hierarchical galaxy formation in a concordance $\\Lambda$CDM universe can produce enough massive and red galaxies compared to the observations. We implement a semi-analytical model in which the central black holes gain their mass during major mergers of galaxies and the energy feedback from active galaxy nuclei (AGN) suppresses the gas cooling in their host halos. The energy feedback from AGN acts effectively only in massive galaxies when supermassive black holes have been formed in the central bulges. Compared with previous models without black hole formation, our model predicts more massive and luminous galaxies at high redshift, agreeing with the observations of K20 up to $z\\sim 3$. Also the predicted stellar mass density from massive galaxies agrees with the observations of GDDS.  Because of the energy feedback from AGN, the formation of new stars is stopped in massive galaxies with the termination of gas cooling and these galaxies soon become red with color $R-K>$5 (Vega magnitude) , comparable to the Extremely Red Objects (EROs) observed at redshift $z\\sim$1-2. Still the predicted number density of very EROs is lower than observed at $z\\sim 2$, and it may be related to inadequate descriptions of dust extinction, star formation history and AGN feedback in those luminous galaxies. ",
        "introduction": "There are many recent observations of high-redshift galaxies that probe the star formation history of the Universe. The finding of many massive galaxies, especially massive Extreme Red Objects (EROs), at high redshift is particularly interesting. These observations show that some EROs are passive ellipticals, and were already in place at redshift z$\\sim 2$.  It is usually argued that in a Cold Dark Matter (CDM) universe, structures form via a hierarchical formation process in which small galaxies form first at early times, and massive galaxies form later through the continuous mergers of the smaller systems. With representative semi-analytical models (SAMs; Kauffmann et al. 1999, Somerville \\& Primack 1999, Cole et al. 2000), it was found that in the concordance $\\Lambda$CDM universe, it is difficult to produce enough massive and red galaxies that look like those observed(e.g. Cimatti et al. 2002a, Glazebrook et al. 2004). On the other hand, the existence of the observed massive galaxies at high redshift is not necessarily in conflict with the concordance $\\Lambda$CDM model, because the conversion of just ten percent of baryons in dark matter halos of mass $M >10^{13}M_{\\odot}$ to stars is sufficient to produce the number of observed massive galaxies (Somerville 2004a).  Many authors have studied the formation of these massive, red objects using SAMs or Smoothed Particle Hydrodynamics (SPH) simulations. It was shown that the SAMs (Kauffmann et al. 1999, Somerville \\& Primack 1999, Cole et al. 2000) cannot produce enough massive/red objects at redshift $z>1$ (e.g. Firth et al 2002, Somerville et al. 2004b, Daddi et al. 2005). The SPH simulations (e.g. Nagamine et al. 2004, 2005) have succeeded in producing massive and red galaxies at high redshift, but at the cost of introducing more uncertainties. First, it is unknown if these SPH simulations can produce the local galaxy luminosity function. It seems that these simulations produce too many bright galaxies at $z=0$ (Nagamine et al. 2004). Secondly, Nagamine et al. (2005) used a high dust extinction for the entire galaxy population, but the observations show that some EROs are passive ellipticals with little dust extinction (Cimatti et al. 2002b).  The main reason that the SAMs fail to produce enough massive and luminous galaxies at high redshift is that the gas cooling and star formation in early massive halos is over-suppressed. In previous SAMs, the gas cooling in massive halos is switched off in order not to produce more luminous central galaxies than observed at redshift $z=0$. The suppression of gas cooling is also motivated by the X-ray observations that massive cooling flows are not observed in groups and clusters (e.g. Peterson et al. 2003). But as the consequence, the gas cooling may be over-suppressed at high redshift if a simplified prescription is used for the cooling cutoff.  For example, in the Munich group model and also in Kang et al. (2005), the gas cooling is shut off by hand in halos with the virial velocity greater than $350km/s$. Since the halo mass is much lower at high redshift than at the present for a given virial velocity, the gas cooling is suppressed in this model for halos with the virial mass greater than 2.5$\\times 10^{12}M_{\\odot}$ at z $=$ 3. This artificial cooling switch-off seems to be the main reason that these models do not produce as many massive galaxies as observed.  In this paper, we implement a new model in which the energy from AGN is used to suppress the cooling of hot gas in halos.  Following Kauffmann \\& Haehnelt (2000) we use a simple model wherein black holes gain most of their mass during major mergers. Our implementation of the feedback from AGN is very similar to that used recently by Croton et al. (2006) and Bower et al. (2005), and resembles a combination of their models. In our model, the total energy from the AGN is proportional to the Eddington luminosity of the central black hole and the efficiency of reheating the gas is proportional to a power of the virial velocity of the galaxy. Then the energy compensates for the radiative energy of the cooling gas, and the actual cooling rate is determined by the ratio between the two energies. The cooling is totally suppressed if the energy from AGN is larger than the energy radiated by the cooling gas.  Compared with the previous model used by Kang et al. (2005) with an artificial cut-off of the gas cooling in the halos with the virial velocity larger than $350km/s$, the gas cooling and AGN feedback in the new model are treated in a more self-consistent way.  The $M_{\\rm bh}$-$\\sigma$ relation of black hole mass $M_{\\rm bh}$ and the bulge velocity dispersion $\\sigma$ implies that massive black holes are present only in massive spheroids. In our present model, the energy feedback from AGN indeed is efficient in galaxies with a massive spheroid. We also require that the star formation rate in quiescent disks is reduced at high redshift as motivated by the observed evolution of cosmological cold gas content with redshift (Keres et al. 2005); thus the gas-rich mergers result in earlier formation of supermassive black holes in massive central bulges. Once the energy feedback is enough to suppress the gas cooling, the termination of new star formation will soon make the galaxies red. We will compare the model prediction of the number density of luminous galaxies with the K20 survey, and find that good agreement holds up to z$\\sim$3, beyond which there is little observational data. Compared with previous SAMs, our present model can also produce some very red ($R-K>5$, magnitudes are given in the Vega system unless otherwise stated) passive ellipticals which are observed by the Great Observatories Origins Deep Survey (GOODS) at z$\\sim 1-2$.  We arrange our paper as follows. In section 2, we briefly introduce our new model with AGN feedback and compare our model predictions with the local galaxy population. In section 3, we give the model predictions and compare them with the observations at high redshift. Finally, we discuss our results and conclude our work in section 4. ",
        "conclusions": "Here we have implemented a new semi-analytical model in which the energy from AGN suppresses hot gas cooling in massive halos. The growth of black holes and bulges, and the gas cooling, are determined in a self-consistent way. In our description, the AGN feedback becomes efficient in massive galaxies after a massive black hole is formed in the galaxy center. The AGN feedback model has drawn much recent attentions. The main motivation is that in massive groups and clusters cooling flows are not observed. There should be some physical process to reheat the cooling region, and the energy from AGN has been proposed as an effective source (e.g. B$\\ddot{\\rm o}$hringer et al. 2002, Begelmen et al. 2002, Sijacki \\& Springel 2006). At the same time, the AGN feedback models have also been incorporated into the SAMs recently and it has been shown that AGN feedback can produce a break of the luminosity function at the bright end and produce the color-magnitude relation observed in SDSS (Croton et al. 2006, Bower et al. 2005). Our model of AGN feedback is very similar to theirs in spirit, but the detailed prescription is different. In this paper we use this model to address some issues about the number distribution and color distribution of galaxies at high redshift. We compare the model predictions with the K20 and GOODS surveys. Our conclusions are as follows. \\begin{itemize} \\item The predicted number distribution of $K<20$ galaxies matches well with that of the GOODS and K20 galaxies up to a redshift of z $\\sim$ 3; \\item The predicted color distribution is similar to that observed in the surveys and many extremely red galaxies ($R-K_{AB}>4$) are produced, which has  not been seen in previous models (Somerville et al. 2004b).  At $z > 1.5$ the galaxy population already displays a bimodal color distribution; \\item The predicted stellar mass density can marginally agree with the GDDS observation even with the uncertainties in the IMFs; \\end{itemize} These results demonstrate that it is not difficult to produce massive and red galaxies at z $\\sim$ 1-2 in the concordance CDM universe.  The stellar mass in galaxy centers continues to grow until the energy from central AGN is high enough to suppress the gas cooling. In our model the black holes acquire most of their mass during major mergers, so the AGN energy feedback is expected to be effective after the last major merger which led to massive bulge formation at galactic centers. In our model we can produce some of those passive ellipticals at z$\\sim 1-2$ with extremely red colors $(R-K)_{AB}>4$.  Many observations have shown that the star formation rate was higher in massive galaxies at high redshift and these support the \"downsizing\" formation scenario (Cowie et al. 1996). It is often argued that hierarchical galaxy formation cannot reproduce the downsizing formation process. But recent works (de Lucia et al. 2005, Bower et al. 2005, Scannapieco et al. 2005) have shown that models with AGN feedback in the hierarchical universe can reproduce the downsizing process in which the massive galaxies forms earlier. In this paper, we also find that the predicted luminous and massive galaxies are increased to the degree that is in agreement with the observations, though the predicted number of red galaxies may still be fewer than observed. Once more observations are available on the dust extinction in these galaxies, the number density and evolution of red passive ellipticals will put more stringent constraints on the galaxy formation models. It is also possible that a new ingredient is needed, such as the star formation induced by AGN feedback prior to disruption of the cold gas supply (Silk 2005), in order to make bulge formation more efficient and to account for the chemical evolution of massive early-type galaxies."
    },
    "0601/astro-ph0601366_arXiv.txt": {
        "abstract": "{ The mass loss rates, expansion velocities and dust-to-gas density ratios from millimetric observations of 119 carbon-rich giants are compared, as functions of stellar parameters, to the predictions of recent hydrodynamical models. Distances and luminosities previously estimated from HIPPARCOS data, masses from pulsations and C/O abundance ratios from spectroscopy, and effective temperatures from a new homogeneous scale, are used. Predicted and observed mass loss rates agree fairly well, as functions of effective temperature. The signature of the mass range $\\rm{M\\le 4\\,M_{\\sun}}$ of most carbon-rich AGB stars is seen as a flat portion in the diagram of mass loss rate {\\it vs.} effective temperature. It is flanked by two regions of mass loss rates increasing with decreasing effective temperature at nearly constant stellar mass. Four stars with detached shells, i.e. episodic strong mass loss, and five cool infrared carbon-rich stars with optically-thick dust shells, have mass loss rates much larger than predicted values. The latter (including CW Leo) could be stars of smaller masses ($\\rm{M\\simeq 1.5-2.5\\,M_{\\sun}}$) while $\\rm{M\\simeq 4\\,M_{\\sun}}$ is indicated for most of the coolest objects. Among the carbon stars with detached shells, R Scl returned to a predicted level (16 times lower) according to recent measurements of the central source. The observed expansion velocities are in agreement with the predicted velocities at infinity in a diagram of velocities {\\it vs.} effective temperature, provided the carbon to oxygen abundance ratio is $\\rm{1\\le \\epsilon _{C}/\\epsilon _{O}\\le 2},$ i.e. the range deduced from spectra and model atmospheres of those cool variables. Five stars with detached shells display expansion velocities about twice that predicted at their effective temperature. Miras and non-Miras do populate the same locus in both diagrams at the present accuracy. The predicted dust-to-gas density ratios are however about 2.2 times smaller than the values estimated from observations. Recent drift models can contribute to minimize the discrepancy since they include more dust. Simple approximate formulae are proposed. ",
        "introduction": "The observations of CO and HCN lines in the millimetric range together with infrared fluxes, especially from the IRAS catalogue (\\cite{IRAS88}), provide estimates of parameters of the circumstellar envelopes of carbon-rich giants. They are essentially TP-AGB stars, i.e. late stages of stellar evolution of low and intermediate mass stars, experiencing substantial mass loss. This latter phenomenon may terminate their evolution on the AGB before rapid nucleosynthesis forces them to leave that branch. Mass loss also contributes to the replenishment of the interstellar medium in carbonaceous material (dust and gas), mostly from extreme (``infrared'') carbon stars. Models and simplifying assumptions are required to derive the estimates of the mass loss rate $\\rm{\\dot{M}}$ in solar mass per year, the expansion velocity $\\rm{v_{e}}$ in kilometer per second, and the dust-to-gas mass density ratio. For instance a model for the photodissociation of the circumstellar CO by Mamon et al. (\\cite{mamon88}) is usually applied to radial brightness distributions with fair success, but clear deviations were noted in a few cases (e.g. Sch\\\"{o}ier \\& Olofsson \\cite{schoie01}). Predominant species, like hydrogen or helium, are not traced out and the adopted abundances influence the estimates of the total mass loss rate. The estimated distances are also of paramount importance since the rate calculated from observations increases as the squared distance. A full discussion is outside the scope of the present paper and we refer the reader to the papers cited in Sect. 2.  It has become clear that the interaction of pulsation and dust formation plays a key role in the mass loss phenomenon (e.g. Wallerstein \\& Knapp \\cite{waller98} for a review). The effects of stellar pulsation are simulated in dynamical models by applying a piston at the inner boundary. Consistent models of circumstellar dust shells around carbon-rich long period variables (LPV) include time-dependent hydrodynamics, a detailed treatment of the processes of formation, growth and evaporation of dust grains, a carbon-rich chemistry, and radiative transfer (e. g. Fleischer et al. \\cite{fleisc92}, Fleischer \\cite{fleisc94}, Winters et al. \\cite{winter94}, \\cite{winter95}, H\\\"{o}fner et al. \\cite{hofner95}, \\cite{hofner96}). The empirical Reimers relation for RGB (Reimers \\cite{reimer75}) proved its inadequacy for AGB. Alternative relations were derived from observational results and/or theoretical models of mass loss (e.g. Volk \\& Kwok \\cite{volk88}, Vassiliadis \\& Wood \\cite{vassi93}, Bl\\\"{o}cker \\cite{blocke95}). The mass loss relation of Dominik et al. (\\cite{domini90}) was the first one based on detailed calculations of dust formation and growth (Gail \\& Sedlmayr \\cite{gail88}), with rather extreme stellar parameters adapted to the very late stages of AGB evolution. Approximate formulae for mass loss rate, expansion velocity and dust to gas density ratio were then calculated by Arndt et al. (\\cite{arndt97}). They considered six stellar parameters as independent ones, namely the temperature, the luminosity, the mass, the abundance ratio of carbon to oxygen, the pulsational period and the velocity amplitude of the pulsation. They then derived approximate equations on the basis of 48 dynamical models, exploring various ranges for the six parameters. They also compared their results to those of 22 analogous models from H\\\"{o}fner \\& Dorfi (\\cite{hofner97}) who proved to be well approximated by their set of equations. \\begin{table*} \\caption {Mean distances (pc) of carbon-rich giants and dispersions as obtained from data of various authors: O93 stands for Olofsson et al. (\\cite{olofs93a}), G99 for Groenewegen et al. (\\cite{groene99}), SO1 for Sch\\\"{o}ier \\& Olofsson (\\cite{schoie01}), GO2 for Groenewegen et al. (\\cite{groen02b}), and L93 for Loup et al. (\\cite{loup93}). They are compared with astrometric data from Bergeat et al. (\\cite{berge02a}, \\cite{berge02b}=BKR) for (n) stars in common. Ratios to BKR means are then quoted in the last line (See text for a discussion).} \\begin{flushleft} {\\scriptsize{ \\begin{tabular}{cccccccccccccccccccc} \\hline\\noalign{\\smallskip} n & O93 & BKR &   & n & G99 & BKR &  & n & O93 & SO1 & BKR &  & n & G02 & BKR &   & n & L93 & BKR\\\\ \\noalign{\\smallskip}\\hline \\noalign{\\smallskip} 104 & 707 & 708 &   & 11 & 1141 & 1161 &  & 54 & 567 & 427 & 614 &  & 25 & 1535 & 748 &  & 21 & 721 & 536\\\\ & $\\pm$ 218 & $\\pm$ 258 &   &    & $\\pm$ 459 & $\\pm$ 452 &  &    & $\\pm$ 171 & $\\pm$ 166 & $\\pm$ 218 & &  & $\\pm$ 927 & $\\pm$ 357 &  &   & $\\pm$ 424 & $\\pm$ 182\\\\ & 1.00 &    &    &    & 0.98 &  &   &    & 0.92 & 0.70 &  &   &    & 2.05 &    &   &   & 1.35\\\\ \\noalign{\\smallskip}\\hline \\end{tabular} }} \\label{mass} \\end{flushleft} \\end{table*}  The predictions of Arndt et al. are compared to the data deduced from observations as published by various authors. We use the effective temperatures of Bergeat et al. (\\cite{bergea01}) as stellar temperatures. Using a new homogeneous scale constructed from model atmospheres, observed angular radii and spectral energy distributions, they have proved efficient in solving various problems. We also adopt the distances and luminosities derived by Bergeat et al. (\\cite{berge02b}) from the HIPPARCOS data (\\cite{ESA97}) by taking into account three distinct biases (their Sect. 2.3; see also Knapik et al. \\cite{knapik98} and Bergeat et al. \\cite{berge02a}). Pulsation masses were derived for carbon-rich LPVs by Bergeat et al. (\\cite{berge02c}). The carbon to oxygen (C/O) ratios were taken from Lambert et al. (\\cite{lamber86}), Olofsson et al. (\\cite{olofs93b}) and Abia \\& Isern (\\cite{abia96}). The periods of pulsation used are the photometric periods taken from the General Catalogue of Variable Stars (GCVS; Kholopov et al. \\cite{GCVS}) or from the HIPPARCOS data (\\cite{ESA97}). The values are those of fundamental mode pulsators from the study of Bergeat et al. (\\cite{berge02c}). Considering the complexities of those pulsating atmospheres and the uncertainties on transfer in their models, we adopt the reference value $\\rm{\\Delta u\\,=\\,2\\,km\\,s^{-1}}$ of Arndt et al. (\\cite{arndt97}) for the velocity amplitude of the pulsation. We shall see that its influence is quite limited.  We re-scaled the mass loss rates from the literature to the distances of Bergeat et al. (\\cite{berge02b}) and a f-correction was applied to convert them to the scale of Sch\\\"{o}ier \\& Olofsson (\\cite{schoie01}) who properly dealt with radiative transfer (Sect. 2). In addition, the expansion velocities and dust-to-gas density ratios were compiled for a sample of carbon-rich giants (Sect. 2). The available predictions from hydrodynamical models are summarized in Sect. 3 and equations fitted on a grid from the literature are adopted. Values of stellar parameters from Bergeat et al. (\\cite{bergea01}, \\cite{berge02b}, \\cite{berge02c}) were compiled for stars in common with millimetric studies and mean values collected for eight photometric groups (Sect. 4) to be replaced in the formulae of Arndt et al. (\\cite{arndt97}). Then, the comparison of observed and predicted mass loss rates is performed in a diagram against effective temperature (Sect. 5). In a similar approach, the observed and predicted expansion velocities are compared by plotting {\\it vs.} effective temperature (Sect. 6). The discrepancy noted between observed and predicted dust-to-gas density ratios, is estimated and discussed (Sect. 7). A few conclusions are given and discussed (Sect. 8) and simple approximate formulas are proposed in Appendix A. ",
        "conclusions": "We have compared the results of millimetric (and infrared) observations, namely gas mass loss rate, expansion (outflow) velocity and dust-to-gas density ratio, to the predictions from hydrodynamical time-dependent models of dust-driven winds for carbon-rich chemistry. Corrections were applied to the published ``observed'' values of the first (distance and radiative transfer) and the third quantity (luminosity and radiative transfer) as described in Sect. 2. The prediction of models are applied to long period variables whose pulsation is simulated by a piston of given amplitude velocity and period acting at the basis of the atmosphere. We made use of the approximate formulae of Arndt et al. (\\cite{arndt97}) based on 48 models exploring ranges in the six parameters they include. Among them, the pulsation period and the piston amplitude velocity have only little influence (Sect. 3). We found no marked difference in the observed data between Miras and non-Miras (here semi-regulars with available periods). Four parameters remain whose ranges may influence the results, namely effective temperature taken as the stellar temperature used in the models, luminosity, mass and carbon to oxygen abundance ratio. The observational data was summarized in Table 2 for 119 carbon-rich giants formerly studied by Bergeat et al. (\\cite{bergea01}, \\cite{berge02b} and \\cite{berge02c}), for effective temperatures, luminosities, stellar masses and C/O abundance ratio. The formulae (4), (5) and (6) as taken from Arndt et al. lead for the mean values of Bergeat et al. to the estimated means of Table 4 for the samples corresponding to the photometric groups HC5 and CV1 to CV7 with a range of effective temperatures from 3470 K down to 1950 K. This latter parameter is considered as the most accurate one at present and it is used for reference. We also hold fixed to reference values the stellar mass and the C/O abundance ratio keeping unchanged the run of effective temperatures and luminosities of Table 4, namely the TP-AGB of Galactic carbon-rich giants in the HR diagram, as obtained by Bergeat et al. (\\cite{berge02c}).  As described in Sect. 5, we obtained a good agreement between predictions and observations in the diagram of mass loss rate {\\it vs.} effective temperature (Fig.$~\\ref{pmvste}$). For about 90\\% of our sample, the mechanism involving gas lifted by pulsation and radiation pressure on condensed dust is confirmed. Three equations (7) and thus (8), (9) and (10) were obtained in the three ranges of effective temperatures distinguished. The usual statement (van Loon et al. \\cite{vanloo03}) that $\\rm{\\dot{M}\\propto T_{eff}^{-8}}$ is approximately true only above 2900 K (Eq. (9)) and below 2400 K (Eq. (10)), with different proportionality coefficients however. There is practically no variation over the 2400-2900 K range in Fig.$~\\ref{pmvste}$ (including about 55\\% of our sample). {\\it Predictions and observations are in agreement provided the variations of luminosity on the AGB in the HR diagram (Bergeat et al. \\cite{berge02b}) and the mean masses ranging from about $\\rm{1M_{\\sun}}$ around 3000 K to about $\\rm{4M_{\\sun}}$ around 2400 K (see Table 4), as deduced from theoretical tracks in the HR diagram and from pulsations masses (Bergeat et al. \\cite{berge02c}).\\/} Here mass is the third important parameter, with influence much larger than that of the C/O abundance ratio. At low effective temperatures, Eq. (10) corresponds to a constant mass of about $\\rm{4M_{\\sun}}.$ At high temperatures, Eq. (9) describes carbon variables of about $\\rm{1M_{\\sun}}$ or less, but both observations and predictions are uncertain there. Nine giants in our sample (at least) do have mass loss rates larger than predicted from Eqs. (7) to (10). Four cases of detached shells are found in the 2600-2800 K range, i.e. objects having experienced highly episodic mass loss possibly caused by a thermal pulse. One of them (R Scl) has returned to a moderate mass loss rate (Wong et al. \\cite{wong04}), and this is certainly true for the others (e.g. Olofsson et al. \\cite{olofss00} for TT Cyg). This is a strong argument in favor of a specific mass loss mechanism at work. The other five stars are very cool objects with $\\rm{T_{eff}\\le 2500 K}$ and optically-thick dust shells shown by strong infrared excesses, a consequence of very high mass loss rates, like RW LMi (CIT6) or CW Leo (IRC+10216). Scho\\\"{\\i}er et al. (\\cite{schoie02}) have shown that the gas mass loss rate of those extreme objects did not vary much over the last thousand of years. As argued in Sect. 5, they may be stars with masses of about $\\rm{1.5-2.5M_{\\sun}},$ in contrast with the $\\rm{4.2M_{\\sun}}$ favored for the cool objects with lower rates as given by Eq. (10). If true, no specific mass loss mechanism would be required here, contrary to the case of detached shells.  A good agreement is also obtained between predictions and observations in the diagram of expansion (outflow) velocity {\\it vs.} effective temperature (Fig.$~\\ref{vevsste}$), as described in Sect. 6. The expansion (outflow) velocity increases with decreasing effective temperature, with about 86 \\% of the objects within the range of predictions as quoted in Table 4 from Eq. (6). To within error bars, the observations are located between the two lines obtained for $\\rm{\\epsilon _{C}/\\epsilon _{O}=1}$ and $\\rm{\\epsilon _{C}/\\epsilon _{O}=2}$ respectively, the other parameters being kept equal to their values in Table 4 (including the mean mass distribution successfully used for mass loss rates and described above). This is the range deduced from spectroscopy and model atmospheres (Lambert et al. \\cite{lamber86}; see also Fig. 7 of Bergeat et al. \\cite{berge02c}). In Fig.$~\\ref{vevsste},$ it is remarkable that for effective temperatures higher than about 2900 K, the stars concentrate towards the $\\rm{\\epsilon _{C}/\\epsilon _{O}=1}$ line, below 10-15 km/s. The maximum observed velocity rapidly decreases with increasing temperature, leaving empty a large triangular area in Fig.$~\\ref{vevsste}.$ The result of Bergeat et al. (\\cite{bergea01}, Fig. 7 in \\cite{berge02b}) is confirmed that above 2900 K, the C/O ratio remains close to unity with little dispersion. The $\\rm{\\epsilon _{C}/\\epsilon _{O}}$ abundance ratio is confirmed here as the third major parameter with effective temperature and luminosity, while stellar mass has little influence in Fig.$~\\ref{vevsste}.$ Cases of superwind however blur the information from that figure. {\\it We conclude that both mass loss rates and outflow velocities are correctly modeled from recent hydrodynamical calculations of dust-driven winds. Both the levels and the dependence on effective temperature are fairly well reproduced. \\/}  Both the predicted and observed ratios of dust-to-gas mass loss rates show little variation with effective temperature if any. As found in Sect. 7, the average observed value of $\\rm{<\\dot{M}_{d}/\\dot{M}_{g}>\\simeq \\left(1.3\\pm 1.1\\right)\\,10^{-3}}$ is a factor 2.2 larger than the mean predicted value $\\rm{<\\rho _{d/g}>=<\\rho _{d}/\\rho _{g}>\\simeq \\left(5.9\\pm 0.7\\right)\\,10^{-4}}$ from Eq. (5). We have discussed in Sect. 7 the influence of the stellar mass and C/O abundance ratio. It should be noted that the uncertainties are very large here. Among the possible explanations is a moderate underestimate of abundance ratios (say 20\\%) from model atmospheres, the overal agreement of Fig.$~\\ref{vevsste}$ being preserved. Another explanation of the above-mentioned disagreement might be some phenomenon in the micro-physics of dust grains, especially the non-equilibrium dust formation process (Andersen et al. \\cite{anders03}, Sandin \\& H\\\"{o}fner \\cite{sandin04} and references therein). Allowing drift in the time-dependent (grain growth) wind models alters their structure. In some cases there is several times more dust in drift models than in non-drift ones, even at low drift velocity (a few km/s). It is also found that drift plays an active role in accumulating dust to certain narrow regions (Sandin \\& H\\\"{o}fner \\cite{sandin04}), which renders the above comparison even more difficult. It was however seen in Sect. 7 that element abundances from model atmospheres raise constraints limiting the magnitude of such effects. Further investigations are needed before we can determine which estimate is closer to truth.  We finally give in Appendix A a few simple approximate formulae in a parametric form. They reproduce the results of the present study with fairly good accuracy. They can be useful for readers wishing not to enter into the full details of the present analysis."
    },
    "0601/astro-ph0601199_arXiv.txt": {
        "abstract": "We present a new algorithm --- Eclipsing Binary Automated Solver (EBAS), to analyse lightcurves of eclipsing binaries. The algorithm is designed to analyse large numbers of lightcurves, and is therefore based on the relatively fast EBOP code.  To facilitate the search for the best solution, EBAS uses two parameter transformations. Instead of the radii of the two stellar components, EBAS uses the sum of radii and their ratio, while the inclination is transformed into the impact parameter.  To replace human visual assessment, we introduce a new 'alarm' goodness-of-fit statistic that takes into account correlation between neighbouring residuals.  We perform extensive tests and simulations that show that our algorithm converges well, finds a good set of parameters and provides reasonable error estimation. ",
        "introduction": "The advent of large CCDs for the use of astronomical studies has driven a number of photometric surveys that have produced unprecedentedly large sets of lightcurves of eclipsing binaries \\citep[e.g.,][]{alcocketal97}. The commonly used interactive way of finding the set of parameters that best fit an eclipsing binary lightcurve utilizes human guess for the starting point of the iteration, and further human decisions along the converging iteration \\citep[e.g.,][]{ribas2000}. Such a process is not always repeatable, and is impractical when it comes to the large set of lightcurves at hand.  The OGLE project, for example, yielded a huge photometric dataset of the SMC \\citep{udalski1998} and the LMC \\citep{udalski2000}, which includes a few thousand eclipsing binary lightcurves \\citep{lukas2003}. This dataset allows for the first time a statistical analysis of the population of short-period binaries in another galaxy. A first effort in this direction was performed by \\citet[][hereafter NZ03]{NZ03}, who derived the orbital elements and stellar parameters of 153 eclipsing binaries in the SMC in order to study the statistical dependence of the eccentricity of the binaries on their separation. In a following study, \\citet[][hereafter NZ04]{NZ04} analyzed another sample of 509 lightcurves selected from the 2580 eclipsing binaries discovered in the LMC by the OGLE team \\citep[][]{lukas2003}. However, the OGLE LMC data contain many more eclipsing binary lightcurves. An automated algorithm would have made an analysis of the whole sample possible.  To meet the need for an algorithm that can handle a large number of lightcurves we developed EBAS --- Eclipsing Binary Automated Solver, which is a completely automatic scheme that derives the orbital parameters of eclipsing binaries. Such an algorithm can be of use for the OGLE lightcurves, as we do in the next paper, and for the data of the many other large photometric surveys that came out in the last few years (e.g., EROS, MACHO, DIRECT, MOA). EBAS is specifically designed to quickly solve large numbers of lightcurves with $S/N$ typical of such surveys.  \\citet[][hereafter WW1]{WW1} have already developed an automatic scheme to analyze the OGLE lightcurves detected in the SMC, in order to find eclipsing binaries suitable for distance measurements. However, whereas WW1 used the Wilson-Devinney (=WD) code, EBAS uses the EBOP code, which is admittedly less accurate than the WD code, but much simpler and faster. We used the EBOP \\citep[][]{PE81,etzel81} subroutines that generate an eclipsing binary lightcurve for a given set of orbital elements and stellar parameters, and rewrote a {\\it fully automated} iterative code that finds the best parameters to fit the observed lightcurve.  As EBAS uses extensively the lightcurve generator for each system, we preferred EBOP over the WD code.  At the last stages of writing this paper another study with an automated lightcurve fitter --- Detached Eclipsing Binary Lightcurve (DEBiL), was published \\citep{devor2005}.  DEBiL was constructed to be quick and simple, and therefore has its own lightcurve generator, which does not account for stellar deformation and reflection effects. This makes it particularly suitable for detached binaries. The complexity of the EBOP lightcurve generator is in between DEBiL and the automated WD code of WW1.  To facilitate the search for the global minimum in the convolved parameter space, EBAS performs two parameter transformations. Instead of the radii of the two stellar components of the binary system, measured in terms of the binary separation, EBAS uses two other parameters, the sum of radii (the sum of the two relative radii), and their ratio. Instead of the inclination we use the impact parameter --- the projected distance between the centres of the two stars during the primary eclipse, measured in terms of the sum of radii.  During the development of EBAS we found that some solutions with low $\\chi^2$ could easily be classified as flawed by visual inspection that revealed correlation between neighbouring residuals.  We have therefore developed a new 'alarm' statistic, $\\mathcal{A}$, to replace human inspection of the residuals. EBAS uses this statistic to decide automatically whether a solution is satisfactory.  The EBAS strategy consists of three stages. First, EBAS finds a good initial guess by a combination of grid searches, gradient descents and geometrical analysis of the lightcurve. Next, EBAS searches for the global minimum by a simulated annealing algorithm. Finally, we asses the quality of the solution with the new 'alarm' statistic, and if necessary, perform further minimum searches.  To check our new algorithm we ran many simulations which demonstrated that the automated code does find the correct values of the orbital parameters. We also used simulations to estimate the error induced by two of our simplifying assumptions, namely mass ratio of unity and negligible third light. We then checked the code against the results of NZ04, and found that our code performed as well as their interactive scheme, except for very few systems. Finally, we checked our code against four LMC eclipsing binaries that were solved by \\citet{gonzalez05} using photometry and radial-velocity data.   Section~\\ref{sec:parameters} presents the EBAS parameters and compares them with the EBOP ones.  Section~\\ref{sec:search} details how the algorithm finds the global minimum of the $\\chi^2$ function, and Section~\\ref{sec:alarm} describes our new alarm statistic. In Section~\\ref{sec:simulation} and Section~\\ref{sec:discussion} we check and discuss the performance of EBAS. ",
        "conclusions": "% \\label{sec:discussion}  We have shown that it is possible to solve lightcurves of eclipsing binaries with a fully automated algorithm which is based on the EBOP code. Our simulations have shown that the results of EBAS are close to the ``real'' ones and that EBAS results for most cases have a quality which is better than is achieved with human interaction.  Although the EBOP code does not include the sophistications offered by, e.g., the widely used Wilson-Devinney programme, it has the advantage of being simple and of producing parameters closely related to the real information content of the lightcurve (e.g., surface brightness instead of effective temperature). In addition, it does take into account not only reflection effects, but also tidal deformation of components (even though in a primitive way), so that it remains useful for systems with moderate proximity effects. Comparison with the recent work of GOMoM05 who used the WD code to analyse three colour photometry, radial-velocity and spectroscopic data of four systems shows that the present version of EBAS recovers quite well the sum of radii, the inclination and the ratio of radii.  It is interesting to compare the speed of the present version of EBAS with the very fast automated DEBiL algorithm \\citep{devor2005}, which was used to derive the elements of almost 10,000 eclipsing binaries in the Galactic bulge.  On the average, it took EBAS 50 seconds CPU time to solve one orbit on an AMD Opteron, 250 2.46GHz, 64-bit machine, while error estimation took another 100 seconds. This is about 3 times longer than it took Devor to solve an orbit with his DEBiL with a SUN UltraSPARC5 333MHz. Applying EBAS to 10,000 systems is therefore feasible.  EBAS uses two redundant techniques to ensure the finding of the global minimum --- a search for the minimum with simulated annealing and consultation with the new alarm $\\mathcal A$. The simulated annealing technique causes heavy computation load on EBAS, and if the number of lightcurves is too large, the demanding parameters of the annealing can be slightly relaxed, since we can rely on the alarm to warn us if the global minimum is not reached.  In the next papers we plan to apply EBAS to the sample of the LMC (Mazeh, Tamuz \\& North 2005, Paper II) and SMC OGLE data. Obviously, when EBAS is applied to real data, one should carefully examine the implication of the specific choices done for the nonvariable parameters, the values of the mass ratio, the fractional light of a possible third star and the values for the limb and gravity darkening. However, the goal of applying EBAS to such large datasets is not to derive the exact parameters of a particular system. Instead, the aim is to study statistical features of the short-period binaries, like their frequency and period distribution. In that sense, the set of data points we use includes the photometry obtained for all the eclipsing binaries found in the sample. For the OGLE LMC data, this is about 300 points for more than 2000 systems, which adds up to about 0.6 millions points, admittedly with low $S/N$ ratio. Such a huge dataset should allow us to study some statistical features of the short-period binaries.  An obvious extension of EBAS would be to allow for automated derivation of the mass ratio, the light of a third star, and even the limb and gravity darkening. This can not be done with OGLE data of the LMC, but would be possible for systems with better data and more than one colour photometry. For close binaries with strong proximity effects we plan to allow EBAS to use the WD code. In principle, the approach should be the same, with the same procedure to find the best initial guess, the same simulated annealing search and the same error estimation. The development of these capacities of EBAS are deferred to a later paper.  Finally, we plan to construct an automated algorithm to derive the masses of the two stars in each eclipsing binary in the LMC, in a similar approach to the one presented by J. Devor in the ``Close Binaries in the 21st Century: New Opportunities and Challenges'' meeting. Our approach relies on the fact that we know the distance to all the binaries in our neighbouring galaxy, up to a few percent, and therefore know the absolute magnitude of the LMC OGLE systems. OGLE data includes some measurements in the $V$ band for each star in the LMC, and therefore the available absolute magnitude information includes two colours. Furthermore, MACHO data \\citep{alcocketal97} is also available for most of these systems. This should suffice to derive a crude estimate of the masses and ages of all the eclipsing binaries in the OGLE LMC dataset."
    },
    "0601/astro-ph0601016_arXiv.txt": {
        "abstract": "The Pierre Auger Observatory is designed to unveil the nature and the origins of the highest energy cosmic rays. The large and geographically dispersed collaboration of physicists and the wide-ranging collection of simulation and reconstruction tasks pose some special challenges for the offline analysis software. We have designed and implemented a general purpose framework which allows collaborators to contribute algorithms and sequencing instructions to build up the variety of applications they require.  The framework includes machinery to manage these user codes, to organize the abundance of user-contributed configuration files, to facilitate multi-format file handling, and to provide access to event and time-dependent detector information which can reside in various data sources.  A number of utilities are also provided, including a novel geometry package which allows manipulation of abstract geometrical objects independent of coordinate system choice. The framework is implemented in C++, and takes advantage of object oriented design and common open source tools, while keeping the user side simple enough for C++ novices to learn in a reasonable time. The distribution system incorporates unit and acceptance testing in order to support rapid development of both the core framework and contributed user code. ",
        "introduction": "\\PARstart{T}{he} offline software framework of the Pierre Auger Observatory~\\cite{Abraham:2004dt} provides an infrastructure to support a variety of distinct computational tasks necessary to analyze data gathered by the observatory. The observatory itself is designed to measure the extensive air showers produced by the highest energy cosmic rays ($> 10^{19}$~eV) with the goal of discovering their origins and composition. Two different techniques are used to detect air showers.  Firstly, a collection of telescopes is used to sense the fluorescence light produced by excitation of nitrogen induced by the cascade of particles in the atmosphere. This method can be used only when the sky is moonless and dark, and thus has roughly a 10\\% duty cycle. The second method uses an array of detectors on the ground to sample particle densities as the air shower arrives at the Earth's surface. Each surface detector consists of a tank containing 12 tons of purified water instrumented with photomultiplier tubes to detect the Cherenkov light produced when particles pass through it.  The surface detector has a 100\\% duty cycle. A subsample of air showers detected by both instruments, dubbed hybrid events, are very precisely measured and provide an invaluable energy calibration tool. In order to cover the full sky, the observatory will consist of two sites, one in the southern hemisphere and one in the north.  The southern site is located in Mendoza, Argentina, and construction there is nearing completion, at which time the observatory will comprise 24 fluorescence telescopes overlooking 1600 surface detectors spaced 1.5~km apart on a hexagonal grid.  Colorado has been selected as the location for the northern site.  The requirements of this project place rather strong demands on the software framework underlying data analysis.  Most importantly, the framework must be flexible and robust enough to support the collaborative effort of a large number of physicists developing a variety of applications over the projected 20 year lifetime of the experiment. Specifically, the software must support simulation and reconstruction of events using surface, fluorescence and hybrid methods, as well as simulation of calibration techniques and other ancillary tasks such as data preprocessing. Further, as the experimental run will be long, it is desirable that the software be extensible in case of future upgrades to the observatory instrumentation.  The offline framework must also handle a number of data formats in order to deal with event and monitoring information from a variety of instruments, as well as the output of air shower simulation codes. Additionally, it is desirable that all physics code be ``exposed'' in the sense that any collaboration member must be able to replace existing algorithms with his own in a straightforward manner.  Finally, while the underlying framework itself may exploit the full power of C++ and object-oriented design, the portions of the code directly used by physicists should not assume a particularly detailed knowledge of these topics.  The offline framework was designed with these principles in mind.  Implementation has taken place over the last three years, and the first physics results based upon this code were recently presented at the $29^{\\rm th}$ International Cosmic Ray Conference~\\cite{icrc2005}. ",
        "conclusions": "We have implemented an offline software framework for the Pierre Auger Observatory. It provides machinery to help collaborators work together on data analysis problems, compare results, and carry out production runs of large quantities of simulated or real data. The framework is configurable enough to adapt to a diverse set of applications, while the user side remains simple enough for C++ non-experts to learn in a reasonable time.  The modular design allows straightforward swapping of algorithms for quick comparisons of different approaches to a problem.  The interfaces to detector and event information free the users from having to deal individually with multiple data formats and data sources.  This software, while still undergoing vigorous development and improvement, has been used in production of the first physics results from the observatory."
    },
    "0601/astro-ph0601546_arXiv.txt": {
        "abstract": "{} {We present new medium-resolution high-S/N optical spectra of the ultracompact low-mass X-ray binaries 4U\\,0614+091 and 4U\\,1626-67, taken with the ESO Very Large Telescope. They are pure emission line spectra and the lines are identified as due to \\ion{C}{ii-iv} and \\ion{O}{ii-iii}.} {Line identification is corroborated by first results from modeling the disk spectra with detailed non-LTE radiation transfer calculations. Hydrogen and helium lines are lacking in the observed spectra.} {Our models confirm the deficiency of H and He in the disks. The lack of neon lines suggests an Ne abundance of less than about 10 percent (by mass), however, this result is uncertain due to possible shortcomings in the model atom. These findings suggest that the donor stars are eroded cores of C/O white dwarfs with no excessive neon overabundance. This would contradict earlier claims of Ne enrichment concluded from X-ray observations of circumbinary material, which was explained by crystallization and fractionation of the white dwarf core.} {} ",
        "introduction": "Low-mass X-ray binaries (LMXBs) consist of a neutron star or black-hole accretor and a low-mass donor star (M\\,$\\lappr$\\,1\\,M$_\\odot$). Of particular interest are those systems with orbital periods P$_{\\rm orb}$\\,$\\lappr$\\,80\\,min, which is the minimum period for LMXBs with hydrogen-rich main sequence donors. In these ultracompact binaries (UCBs) the mass donor must be a non-degenerate hydrogen-deficient star or a white dwarf (e.g.\\ Verbunt \\& van den Heuvel 1995). Currently eight such systems with measured orbital periods (P$_{\\rm orb}$=11-50\\,min) are known (Ritter \\& Kolb 2003).  \\begin{figure*}[tbp] \\includegraphics[width=\\textwidth]{fig1.ps} \\caption[]{Blue spectra of 4U\\,0614+091 and 4U\\,1626-67. H and He lines are lacking. The emission lines are identified as \\ion{O}{ii-iii} and \\ion{C}{ii-iii} (see also Table\\,\\ref{lines_tab}). All absorption features are of interstellar (i.s.) origin: \\ion{Ca}{ii} 3934/3968~\\AA, the 4300~\\AA\\ CH band, and two diffuse interstellar bands (DIB). Plotted over the two observed spectra is a synthetic accretion disk spectrum reddened with E(B-V)=0.55 and E(B-V)=0.40, respectively.  } \\label{fig1} \\end{figure*}  \\begin{figure*}[tbp] \\includegraphics[width=\\textwidth]{fig2.ps} \\caption[]{Red spectra; 4U\\,0614+091 shows emission lines from \\ion{C}{ii-iv} and \\ion{O}{ii-iii}, while the 4U\\,1626-67 spectrum is virtually featureless. All absorption features are either telluric (marked by horizontal bars) or of interstellar origin (\\ion{Na}{i} 5890/5896~\\AA\\ and a DIB at 5780~\\AA). The features at 5575~\\AA\\ are artifacts (``x'') due to an \\ion{O}{i} sky emission line. Overplotted is the synthetic accretion disk spectrum like in Fig.\\,\\ref{fig1}. } \\label{fig2} \\end{figure*}  \\begin{figure*}[tbp] \\includegraphics[width=\\textwidth]{fig3.ps} \\caption[]{Complete VLT spectrum of 4U\\,1626-67 compared to a spectrum taken with HST. The VLT spectrum, scaled to the HST spectrum at 5600~\\AA, is flatter possibly due to problems in absolute flux calibration. The HST spectrum has a poorer resolution and S/N-ratio, although some emission features at $\\lambda<5000$~\\AA\\ can be recognized in both datasets. Artificial emission spikes are marked by ``x''. Overplotted is the synthetic accretion disk spectrum, like in previous figures, but now reddened with E(B-V)=0.20 to fit the continuum shape of the HST spectrum. } \\label{fig3} \\end{figure*}  \\begin{figure*}[tbp] \\includegraphics[width=\\textwidth]{fig4.ps} \\caption[]{Detail from three single spectra of 4U\\,0614+091. Thick line: Spectrum taken on Nov. 03, 2003. Two thin lines: Spectra taken consecutively on Dec. 15, 2003 (see Table~\\ref{obs_tab}). The strength of the two strongest emission features, the \\ion{C}{ii} multiplets at 6580~\\AA\\ and 7235~\\AA, show an increase over a time interval of several weeks. } \\label{fig4} \\end{figure*}   In the recent past, the existence of a group of five ultracompact systems with neon-rich white dwarf donors has been claimed based on X-ray spectral properties (Schulz \\etal 2001, Juett \\etal 2001, Juett \\& Chakrabarty 2003). The donors are then C/O or even O/Ne/Mg white dwarfs that have transferred a significant fraction of their mass to the neutron star, in contrast to the usual wisdom that the donors are the remains of He white dwarfs. This has caused new explorations of the formation of these systems (e.g.\\ Yungelson \\etal 2002). Our motivation for studying these systems is that the stripped donor stars offer the possibility to probe the interior composition of white dwarfs, which depends on the interplay of gravitational settling and crystallization of chemical elements.  Three of these five Ne-rich systems belong to the above-mentioned eight UCBs with measured orbital periods, while two of them are believed to be UCBs because of their similar optical and X-ray properties. The class of Ne-rich UCBs consists of: \\begin{tabbing} (1) 4U\\,1626-67  \\quad  \\= (P$_{\\rm orb}$=41\\,min)\\\\ (2) 4U\\,0614+091        \\>                        \\\\ (3) 2S\\,0918-549        \\>                        \\\\ (4) 4U\\,1543-624        \\> (P$_{\\rm orb}$=18\\,min, Wang \\& Chakrabarty 2004)\\\\ (5) 4U\\,1850-087        \\> (P$_{\\rm orb}$=20\\,min) \\end{tabbing} And 4U\\,1626-67 may be regarded as the prototype of this class. The donor's Ne-rich C/O-WD nature is derived from X-ray and ultraviolet spectra that exhibit double-peaked emission lines that obviously stem from the accretion disk (Schulz \\etal 2001, Homer \\etal 2002).  The close relation of the other four objects (2)\\,--\\,(5) to 4U\\,1626-67 was based on extraordinary high Ne/O abundance ratios (when compared to the ISM value) measured from ASCA spectra of (3)--(5) and a Chandra  spectrum of (2). The spectra exhibit photoelectric absorption edges of neutral O and Ne in the interstellar medium (ISM) along the line-of-sight, which is suspected to originate from expelled material close to the binary systems. New X-ray spectroscopic observations of (3) and (4) obtained with Chandra and XMM-Newton confirmed the ASCA results, although different values for the Ne/O enrichment have been obtained from a Chandra and a XMM-Newton spectrum of (4) (Juett \\& Chakrabarty 2003). In contrast to the ASCA measurement of (5), recent XMM-Newton and Chandra observations found no evidence of an unusual Ne/O ratio (Sidoli \\etal 2005, Juett \\& Chakrabarty 2005). These apparently contradictory results can be attributed to a variable \\ion{Ne}{i}/\\ion{O}{i} ratio due to changes in the ionisation structure in the measured absorption columns that, however, are not understood (Juett \\& Chakrabarty 2005). Hence, this means that the measured Ne/O ratio does not reflect the donor composition. Although the orbital periods of these systems are below 80\\,min (or at least believed to be so small), this does not necessarily mean that they contain C/O donors. For example, the X-ray burst properties of the 4U\\,1820-30 (P$_{\\rm orb}$=11\\,min) suggest an He-WD donor in that ultracompact system (e.g.\\ Strohmayer \\& Brown 2002). Another example are the relatively short X-ray bursts observed from (3) that even suggest H-burning of material accreted onto the neutron star (NS) (Jonker \\etal 2001), while optical spectra suggest a C/O WD donor (Nelemans \\etal 2004 and this work). A possible explanation is that the heavy elements (C, O, Ne) undergo spallation during accretion leaving H and He nuclei for thermonuclear burning on the NS (Bildsten \\etal 1992). In a recent paper In't Zand \\etal (2005) conclude that the Ne/O ratio and the X-ray burst properties are all best explained with an He-rich donor.  In conclusion, the only way to confirm that the four systems (2)--(5) -- besides the prototype (1) -- indeed contain C/O-rich donors, perhaps enriched with Ne, is by UV and/or optical spectroscopy. Nelemans \\etal (2004) show that the optical spectra of (2)--(4) are devoid of hydrogen and helium emission lines, which are usually seen in H- or He-rich accreting systems. Their spectra exhibit low-ionisation C and O emission lines, most prominent in the brightest of these objects, 4U\\,0614+091.  In this paper we present new optical spectra of this system that cover a larger wavelength interval. They confirm the earlier conclusion by Nelemans \\etal (2004) that the emission lines arise from the C/O-dominated material that is probably located in the accretion disk and not in the X-ray heated face of the white-dwarf donor. We also present a first detailed optical spectrum of the prototype 4U\\,1626-67, which mainly shows weak oxygen emission lines and proves the H- and He- deficiency in this system, too. Nelemans \\& Jonker (2005) also performed VLT observations of this system (also in spring 2004) with a similar setup. They show a section of their spectrum and emphasize the similarity with 4U\\,0614+91.  In addition we present results from first attempts to model the observed spectra with non-LTE accretion-disk models. We derive upper limits for the H and He abundances and investigate the formation of neon lines. Detailed C and O line-formation calculations can already qualitatively explain the observed emission lines. The ultimate goal is to derive detailed abundances and other disk parameters from the observed line profiles.  The paper is organised as follows. We describe our observations in the following section. We then present our line identifications in Sect.~3. Section~4 contains a description of our disk model calculations and we summarise and conclude in Sect.~5. ",
        "conclusions": "We have presented new high-quality optical spectra of the ultracompact low-mass X-ray binaries 4U\\,0614+091 and 4U\\,1626-67. They are pure emission line spectra and most probably stem from the accretion disk. The spectral lines are identified as due to \\ion{C}{ii-iv} and \\ion{O}{ii-iii}. Line identifications are corroborated by first results from modeling the disk spectra with detailed non-LTE radiation transfer calculations. Hydrogen and helium lines are lacking and our models confirm the deficiency of H and He in the disk. Hence, the donor stars in these systems are in fact the eroded cores of C/O white dwarfs. There are indications that the O/C ratio in 4U\\,0614+091 is higher than in 4U\\,1626-67. This could suggest that the stripping process of the WD in 4U\\,0614+091 is more advanced.  It is hard to estimate the systematic error of the derived upper limit for the neon abundance so that the following conclusions are at the moment rather uncertain. From the lack of \\ion{Ne}{ii} lines we find that the Ne abundance cannot exceed $\\approx 10\\%$ or, in other words, the Ne/O ratio is at most of the order 0.2. The much higher Ne/O ratios ($\\approx 0.7$) derived from the ISM X-ray absorption edges of neutral Ne and O (Schulz \\etal 2001, Juett \\etal 2001) would produce detectable \\ion{Ne}{ii} lines in the disk spectra. This confirms the conclusion of Juett \\& Chakrabarty (2005), that the determined ISM abundances of Ne and O are affected by ionisation effects and, hence, do not reflect the abundances of the donor stars.  For an initial solar metallicity the $^{22}$Ne abundance in an Ne-enriched crystallized and fractionated WD core can be estimated from theoretical models to $\\approx 0.07$, but this value is very uncertain (Yungelson \\etal 2002). It could be even higher by a factor of 3 (Isern \\etal 1991).  Such an Ne-rich core would have a mass of $\\approx 0.06$~M$_\\odot$ (Yungelson \\etal 2002). The mass of the WD donor in 4U\\,1626-67 is much smaller (0.01~M$_\\odot$; Yungelson \\etal 2002). If we accept an upper limit of Ne=0.1 for the observed neon abundance, then we may draw the following conclusion. Either the WD core has crystallized and fractionated, then our observation favors a relatively small Ne-enrichment as a result of the crystallization process. Or, if one accepts that fractionation would result in a high Ne abundance of 0.2 (unobserved), then the WD core in 4U\\,1626-67 had no time to crystallize which, depending on various details, lasts several Gyr (e.g.\\ Hernanz \\etal 1994).  Future work will concentrate on the disk modeling for the prototype 4U\\,1626-67, for which the observational database is the best of all such systems. Ultimately, the spectral properties (flux distribution and emission line strengths) over the complete wavelength range comprising spectral observations with Chandra, HST, and VLT, must be explained by a unique disk model."
    },
    "0601/astro-ph0601636_arXiv.txt": {
        "abstract": "In the first spectroscopic campaign for a PG\\,1716 variable (or long-period pulsating subdwarf B star), we succeeded in detecting velocity variations due to g-mode pulsations at a level of 1.0--1.5~km~s$^{-1}$, just above our detection limit.  The observations were obtained during 40 nights on 2~m class telescopes in Arizona, South Africa, and Australia.  The target, PG\\,1627+017, is one of the brightest (V = 12.9) and largest amplitude ($\\sim$0.03 mag) stars in its class.  It is also the visible component of a post-common envelope binary.  Our final radial velocity data set includes 84 hours of time-series spectroscopy over a time baseline of 53 days, with typical errors of 5--6~km~s$^{-1}$ per spectrum.  We combined the velocities with previously existing data to derive improved orbital parameters.  Unexpectedly, the velocity power spectrum clearly shows an additional component at twice the orbital frequency of PG\\,1627+017, supporting Edelmann et al.'s recent results for several other short-period subdwarf B stars, which they claim to be evidence for slightly elliptical orbits.  Our derived radial velocity amplitude spectrum, after subtracting the orbital motion, shows three potential pulsational modes 3--4$\\sigma$ above the mean noise level of 0.365~km~s$^{-1}$, at 7201.0~s (138.87~$\\mu$Hz), 7014.6~s (142.56 $\\mu$Hz) and 7037.3~s (142.10~$\\mu$Hz).  Only one of the features is statistically likely to be real, but all three are tantalizingly close to, or a one day alias of, the three strongest periodicities found in the concurrent photometric campaign.  We further attempted to detect pulsational variations in the Balmer line amplitudes.  The single detected periodicity of 7209~s, although weak, is consistent with theoretical expectations as a function of wavelength.  Furthermore, it allows us to rule out a degree index of $l$ = 3 or $l$ = 5 for that mode.  Given the extreme weakness of g-mode pulsations in these stars, we conclude that anything beyond simply detecting their presence will require larger telescopes, higher efficiency spectral monitoring over longer time baselines, improved longitude coverage, and increased radial velocity precision. ",
        "introduction": "Subdwarf B (sdB) stars are evolved hot compact stars (22,000~K $<$ T$_{\\rm eff} <$ 40,000K, 5.0 $<$ log~{\\it g} $<$ 6.2) that are commonly found in the disk of our Galaxy (Saffer et al.\\ 1994). They are core helium-burning stars with extremely thin hydrogen envelopes ($<$~0.01$M_{\\sun}$), believed to have masses of about 0.5$M_{\\sun}$ (Saffer et al.\\ 1994; Heber 1986).  There is considerable interest in their evolutionary history, in part because a large fraction of sdB stars occur in post-common envelope binaries ({\\it e.g.}\\ Allard et al.\\ 1994; Green, Liebert \\& Saffer 1997; Maxted et al.\\ 2001, Morales-Rueda et al.\\ 2003).  Their structural details provide an important independent test of stellar evolution theory because the vast majority of low mass stars are expected to develop nearly identical helium cores during their red giant phase.  Within the last decade, two different types of multimode pulsators have been found among sdB stars: shorter period variables exhibiting mainly pressure (p-)modes, and longer period variables whose pulsations must be due to gravity (g-)modes.  These discoveries have opened up the possibility of using asteroseismology to investigate the structure of sdB stars, and, by extension, the helium-burning cores of most red giants.  The first short-period pulsating sdB stars were discovered by astronomers from the South African Astronomical Observatory (Kilkenny et al.\\ 1997), and are commonly called EC\\,14026 stars\\footnote{Now formally named the V361 Hya stars}, after the prototype. Independently and nearly simultaneously, their existence was predicted by Charpinet et al.\\ (1996, 1997), based on a $\\kappa$-mechanism associated with the radiative levitation of iron in the thin diffusion-dominated envelopes.  The resulting low-order and low-degree ($l$=0,1,2,3) radial and nonradial p-modes produce velocity variations primarily along the star's radial axis.  Typical EC\\,14026 stars have pulsation periods of 100--200~s with amplitudes of a few hundredths of a magnitude or less.  As predicted by Charpinet et al., they are found mostly among the hotter sdB stars, clustering around T$_{\\rm eff}$ $\\sim$ 33,500~K and log~{\\it g} $\\sim$ 5.8, although a few cooler, lower gravity examples exist with somewhat longer periods (up to about 10 minutes).  The current tally of published EC\\,14026 stars is 34 (Kilkenny 2002a, and references therein; Silvotti et al.\\ 2002; Dreizler et al.\\ 2002; Bonanno et al.\\ 2003; Oreiro et al.\\ 2004; Solheim et al.\\ 2004; Baran et al.\\ 2005), although more than half are relatively faint (V~$<$~14).  The second class of sdB pulsators was serendipitously discovered by Green et al.\\ (2003) during a light curve survey originally intended to search for binary eclipses and reflection effects.  PG\\,1716 stars (Reed et al.\\ 2004), sometimes called long-period sdB variables, pulsate on timescales of approximately an hour.  They have extremely small photometric amplitudes, of the order of millimags (Randall et al.\\ 2004, 2005bc).  In contrast to the EC\\,14026 stars, long-period sdB pulsators are found only among cooler sdB stars, clustering around T$_{\\rm eff}$ $\\sim$ 27,000~K and log~{\\it g} $\\sim$ 5.4. Another difference is that the slower pulsators are much more common. Thirty of these variables, roughly 80\\% of sdB stars with temperatures between 24000~K and 29500~K, have been confirmed since 2001, all but two brighter than V~=~14 (Green et al., in preparation).  Fontaine et al.\\ (2003) showed that PG\\,1716 stars can be excited by the same $\\kappa$-mechanism proposed by Charpinet et al., if their g-modes are due to higher-degree ($l$=3,4,5...)  nonradial pulsations.  Thus, although the analogy is not perfect, sdB p-mode and g-mode pulsators bear a strong resemblance to the main sequence $\\beta$ Cephei and Slowly Pulsating B stars, respectively.  Pulsational modes are most easily detected photometrically.  Since the first discovery, several groups have conducted surveys to search for new EC\\,14026 stars (Kilkenny 2002b, and references therein; Bill\\`eres et al.\\ 2002; Solheim et al.\\ 2004) and PG\\,1716 stars (ongoing Steward Observatory survey).  There have now been a number of photometric campaigns on selected short-period pulsators (e.g., Kilkenny et al.\\ 2002b), and more recently, long-period pulsators (Reed et al.\\ 2004; Randall et al.\\ 2004, 2005bc).  The common goal is to detect a sufficient number of pulsational modes for asteroseismology.  Unfortunately, experience has shown that photometry alone, especially if carried out on small telescopes, is usually insufficient to identify these modes unambiguously with a single theoretical model.  Unique identifications have been achieved for less than a handful of sdB pulsators (all observed on a 4 m class telescope), most notably PG\\,0014+067 (Brassard et al.\\ 2001), PG\\,1219+534 (Charpinet et al.\\ 2005b), and Feige\\,48 (Charpinet et al.\\ 2005a), but also including PG\\,1047+003 (Charpinet, Fontaine, \\& Brassard 2003).  Charpinet et al.\\ (2000) pointed out that spectroscopic detection of radial velocity (RV) modes should contain complementary information that might help to solve the mode identification problem.  The brightest EC\\,14026 star, PG\\,1605+072, was the ideal first target from an observational point of view, since it has the largest photometric amplitude variations (0.2 mag) of any sdB pulsator, and a very rich pulsational spectrum (Kilkenny et al.\\ 1999). Due to its unusually low surface gravity, it also has the longest periods ($\\sim$200--550s) of any EC\\,14026 variable.  O'Toole et al.\\ (2000) were the first to detect radial velocity shifts in PG\\,1605+072 using medium-resolution time-resolved spectroscopy.  Additional time-series spectroscopy has since been carried out at the 4~m Anglo Australian Telescope and the William Herschel Telescope (Woolf, Jeffery, \\& Pollacco 2002a) at higher spectral ($<$ 1\\AA) and temporal resolution, and also at the Danish 1.5~m and the Mt.\\ Stromlo 1.9~m by O'Toole et al.\\ (2002). Woolf et al.\\ resolved three pulsation modes, while O'Toole et al.\\ identified five modes, all of which were present in Kilkenny et al.'s (1999) photometry.  In 2002, an ambitious world-wide campaign, the Multi-Site Spectroscopic Telescope (MSST), was organized to resolve closely spaced periods in PG\\,1605+072 and to obtain more precise velocity amplitudes for weaker modes (Heber at al.\\ 2003; O'Toole et al.\\ 2004).  This resulted in 151 hours of spectroscopy, leading to at least 20 detected periodicities in PG\\,1605+072.  Woolf, Jeffery, \\& Pollacco (2002b) also investigated KPD\\,1930+2752, a fainter EC\\,14026 variable (V~=~13.8), but found no significant peaks in the RV power spectrum above their $\\sim$~4~km~s$^{-1}$ detection limit that matched known photometric modes.  Telting \\& \\O stensen (2004) succeeded in recovering the one dominant 138~s pulsation mode in PG\\,1325+101 (also V~=~13.8) at an amplitude of nearly 12~km~s$^{-1}$, using spectra from three nights on the Nordic Optical Telescope.  However, it is apparent that it will be quite difficult to obtain high quality spectroscopic detections of a sufficiently large number of frequencies to clarify mode identification in significantly fainter, shorter period, and lower amplitude p-mode sdB pulsators, given the short exposures required and the difficulties of obtaining large blocks of time on bigger telescopes.  During the planning stages of the first multisite photometric campaign on a long-period sdB pulsator, PG\\,1627+017 (Randall et al.\\ 2004), which took place during April--June 2003, it was natural to consider the feasibility of a complementary spectroscopic campaign to detect g-mode pulsations.  Near-simultaneous spectroscopy and photometry are essential, because the amplitude of individual pulsation modes can change significantly over time (e.g.\\ Kilkenny et al.\\ 1999; O'Toole et al.\\ 2002).  Furthermore, assessing the reality of marginal spectroscopic detections is much simpler if the same frequencies are also present in the photometry.  One advantage of spectroscopy of long-period PG\\,1716 variables is that they are more common than their short-period EC\\,14026 counterparts, and thus there exist more bright targets for future investigations if spectroscopically observed g-modes prove helpful for mode identification.  A primary disadvantage is that their pulsational amplitudes are significantly smaller.  The fact that typical PG\\,1716 pulsational periods are roughly a factor of thirty longer than the periods in EC\\,14026 stars is both a blessing and a curse.  Periods of an hour or longer can be more easily sampled spectroscopically using smaller telescopes, but the effective time baseline and the number of telescope nights required for comparable frequency resolution must also be about thirty times larger.  From the ground, multisite observations to reduce daily aliases are much more critical for the long-period sdB variables, since only 4--10 cycles can be observed per night from any one site.  With these factors in mind, we organized a spectroscopic collaboration between Steward Observatory (SO) in Arizona, the South African Astronomical Observatory (SAAO), and Siding Spring Observatory (SSO) in New South Wales, Australia, which would potentially allow nearly 24~hour coverage of PG\\,1627+017 with 2~m class telescopes.  Ideally, we hoped to get as much overlapping time as possible during an approximately 7--10 day period in May, with additional nights at two and four week intervals before and after the joint time in May, to improve the time baseline.  As will be seen in \\S 4, the telescope allocation committees were as generous as we could have wished, but the weather was considerably less so.  In the end, we were able to acquire (barely) sufficient data for a productive first look at g-modes in long-period sdB stars.  One additional point deserves consideration.  In addition to RV variations, the relative pulsational amplitudes as a function of wavelength provide another way to constrain mode identifications. Theory predicts that pulsational amplitudes should be much greater at the central wavelengths of the Balmer lines than in the nearby continuum.  O'Toole et al. (2003) found a significant increase in their line index flux for higher order Balmer lines in PG\\,1605+072. The continuum amplitudes also increase slowly from redder wavelengths to bluer ones, with a particularly rapid increase from blue to ultraviolet.  The ratio of UV/red flux is a strong function of the degree index {\\it l} (Randall et al.\\ 2004, 2005a).  To help investigate the amplitude dependence on wavelength in our PG\\,1627+017 data, we obtained a number of additional spectra for the similar but apparently nonvariable sdB, PG\\,1432+004, in order to calculate the atmospheric extinction correction with wavelength.  In the remainder of this paper, we describe our multisite time-series spectroscopy for PG\\,1627+017 and present the results. In \\S 2 and \\S 3, we discuss the observed characteristics of our target and the feasibility of detecting g-modes in this star.  \\S 4 and \\S 5 describe the observations obtained at each site and the reduction methods. A description of the radial velocity cross-correlations and the derived orbital parameters is given in \\S 6.  We discuss the power spectrum of the residual velocities in \\S 7, and consider the wavelength dependence of the pulsational amplitudes in \\S 8.   \\S 9 contains the summary and conclusions. ",
        "conclusions": "We were allocated 40 nights during March--July 2003 on 2~m class telescopes in Arizona, South Africa, and Australia to attempt the first spectroscopic detection of gravity modes in a long-period (g-mode) sdB pulsator.  Our equatorial target, PG\\,1627+017, is one of the brightest PG\\,1716 variables, with moderately large amplitudes for its type.  Previous photometry suggested pulsational periods of the order of 1--2 hours and amplitudes $\\lta$1\\%.  For ground-based observations, minimum time baselines of many weeks are required to adequately resolve the modes, as well as good longitude coverage to eliminate aliasing.  Order-of-magnitude calculations suggested that we would need a detection limit $\\lta$1~km~s$^{-1}$ and a velocity precision about 4--5~km~s$^{-1}$, to set useful upper limits on possible g-mode amplitudes.  Ideally, we wanted to obtain time-series spectroscopy on 2~m class telescopes for a week of overlapping nights at all three observatories during the middle of the corresponding photometric campaign, plus additional blocks of 4--5 nights at a time at Steward Observatory, every few weeks for at least a month or so before and after the main effort.  Although such spotty coverage can't provide very good resolution, it took a huge effort merely to obtain this much 2~m telescope time and carry out the observations.  It would be extremely difficult to get a higher duty cycle without dedicated telescopes.  Unfortunately, bad weather limited the number of our useful nights to a bare minimum of 18: a total of 11 nearly overlapping May nights from the three sites, plus one additional April night and five June nights in Arizona.  The final velocity data set included 84 hours of time-series spectroscopy over a formal time baseline of 53 days.  We combined our new velocities for PG\\,1627+017 with published velocities from Morales-Ruelas et al.\\ (2003), and with newly reported MMT velocities (Table~3), to derive improved orbital parameters (Table~4).  Surprisingly, the radial velocity power spectrum shows a significant peak at 27.910~$\\mu$Hz, which, to within the errors, is exactly half of the orbital period (27.916~$\\mu$Hz), corroborating Edelmann et al.'s (2005) discovery of similar small velocity distortions in several other short-period sdB stars.  Edelmann et al.\\ conclude that slightly elliptical orbits are the only possible explanation.  We are not completely convinced that this is the case, but we could not find any alternative explanations.  The simultaneous photometry and spectroscopy campaigns for PG\\,1627+017 constrain the possibilities even more tightly, because the mechanism responsible for the half-orbital velocity term produced no detectable brightness variations.  In spite of an overall efficiency of only 6.6\\% and a somewhat higher detection limit than we had hoped, we appear to have reached the level where g-modes can be observed in a sdB star.  The power spectrum of the residual velocities, after subtracting the orbital motion, shows excess power in the same 120--170~$\\mu$Hz frequency range as the observed photometric pulsations.  Three peaks are found at 7201.0~s (138.87~$\\mu$Hz), 7014.6~s (142.56~$\\mu$Hz) and 7037.3~s (142.10~$\\mu$Hz), with amplitudes that are 3--4$\\sigma$ above the mean noise level of 0.365~km~s$^{-1}$.  At least one of the peaks is likely to be real based on noise statistics alone, but all three are good candidates because each one is tantalizingly close to, or a one day alias of, one of the three strongest photometric peaks found by Randall et al.\\ (2004, 2005b).  As this is most unlikely to happen by chance, we conclude that we have indeed detected g-mode pulsations in PG\\,1627+017 at the 1.0--1.5~km~s$^{-1}$ level.  This indicates that the strongest velocity mode in this cool, low gravity, g-mode pulsator, relative to its $\\sim$0.03 mag photometric variations, is only about a factor of two weaker than the strongest velocity mode (14~km~s$^{-1}$) observed in the correspondingly cool, low gravity p-mode pulsator, PG\\,1605+072, relative to its 0.2 mag photometric variations.  We further attempted to detect pulsational variations in the Balmer line amplitudes.  The higher order Balmer lines, in particular, showed evidence for amplitude variations with a periodicity of 7209~s, once again quite close to the dominant photometric frequencies.  The trend of the pulsational amplitudes with wavelength is consistent with with theoretical expectations, and we were able to rule out the $l$ = 3 and $l$ = 5 possibilities for that mode.  We conclude that spectroscopic detection of gravity modes in sdB stars is (just) within the realm of possibility for the strongest PG\\,1716 pulsators.  Truly useful results for velocity modes would require much higher efficiency spectral monitoring over at least a comparable time baseline (6--8 weeks), with a radial velocity accuracy closer to 1--2~km~s$^{-1}$.  Achieving this level of efficiency and precision with ground-based facilities is very unlikely at present, since it would require long blocks of telescope time on multiple telescopes larger than 2~m, with good longitude coverage.  Since the necessary RV accuracy would also be a problem for satellite observations, it appears that spectroscopic mode detection in the PG\\,1716 stars, as well as in many EC\\,14026 stars, has a way to go before it will be truly practical."
    },
    "0601/astro-ph0601400_arXiv.txt": {
        "abstract": "Several recent observations suggest that gas-poor (dissipationless) mergers of elliptical galaxies contribute significantly to the build-up of the massive end of the red sequence at $z \\la 1$.  We perform a series of major merger simulations to investigate the spatial and velocity structure of the remnants of such mergers. Regardless of orbital energy or angular momentum, we find that the stellar remnants lie on the \\fpp\\ defined by their progenitors, a result of virial equilibrium with a small tilt due to an increasing central dark matter fraction.  However, the locations of merger remnants in the projections of the \\fpp\\ -- the Faber-Jackson and $\\re-\\mstar$ relations -- depend strongly on the merger orbit, and the relations steepen significantly from the canonical scalings ($L\\propto \\sige^4$ and $\\re\\propto \\mstar^{0.6}$) for mergers on radial orbits.  This steepening arises because stellar bulges on orbits with lower angular momentum lose less energy via dynamical friction on the dark matter halos than do bulges on orbits with substantial angular momentum.  This results in a less tightly bound remnant bulge with a smaller velocity dispersion and a larger effective radius.  Our results imply that the projections of the \\fpp\\ -- but not necessarily the plane itself -- provide a powerful way of investigating the assembly history of massive elliptical galaxies, including the brightest cluster galaxies at or near the centers of galaxy clusters.  We argue that most massive ellipticals are formed by anisotropic merging and that their \\fpp\\ projections should thus differ noticeably from those of lower mass ellipticals even though they should lie on the same \\fpp.  Current observations are consistent with this conclusion.  The steepening in the $L-\\sige$ relation for luminous ellipticals may also be reflected in a corresponding steepening in the $\\mbh-\\sige$ relation for massive black holes. ",
        "introduction": "Several independent lines of evidence point to the importance of gas-poor (dissipationless) mergers in the assembly of massive elliptical galaxies.  For example, the decreasing rotational support and transition from disky to boxy isophotal shapes with increasing stellar mass in elliptical galaxies suggests an increasing fraction of dissipationless mergers in the growth of the most massive elliptical galaxies (e.g., \\citealt{bender1992, kormendy1996, faber1997, naab2006}).  The lack of correlation between metallicity \\citep{gallazzi2005} or metallicity gradients \\citep{carollo1993} and galaxy mass for massive ellipticals, while both are positively correlated with galaxy mass for lower-mass ellipticals, provides additional support for this picture.  There is also indirect evidence from the evolution of galaxy luminosity functions \\citep{bell2004,faber2005,blanton2005a} and semi-analytic models of galaxy formation \\citep{kauffmann2000,khochfar2003, de-lucia2006} that a significant fraction of low-redshift ellipticals were assembled via dissipationless mergers.  Moreover, there is direct observational evidence for a substantial number of ``red'' mergers between early-type galaxies in an intermediate redshift galaxy cluster \\citep{van-dokkum1999, tran2005} as well as in more local galaxies \\citep{bell2005, van-dokkum2005}.  In spite of the importance of gas-poor mergers, however, the properties of these mergers and their remnants are not fully understood.  If elliptical galaxies gain substantial mass via dissipationless merging at $z\\la 1$, this process could have a significant impact on their properties.  The purpose of this paper is to study how such mergers affect the global dynamical and kinematic structure of the remnants, as ellipticals show a very tight correlation in the space spanned by their half-light radii ($\\re$), central velocity dispersions ($\\sige$), and central surface brightnesses ($I_e$) known as the \\fpp\\ \\citep{dressler1987, djorgovski1987}.  Relationships between two of the \\fpp\\ variables, such as the Faber-Jackson relation \\citep{faber1976} and the stellar mass-size relation, may also serve as important constraints for galaxy formation models (e.g., \\citealt{mcintosh2005}).  This is particularly true for the most luminous elliptical galaxies, including brightest cluster galaxies (BCGs) that populate the exponential tails of the galaxy luminosity, mass, and velocity dispersion functions and contain the most massive black holes in the universe.  These galaxies are likely to have undergone multiple generations of gas-poor mergers, as repeated merging during cluster formation is a well-motivated mechanism for BCG formation \\citep{merritt1985,dubinski1998}.  In this model, anisotropic infall along cosmological filaments introduces a preferential direction in the merging process \\citep{west1994}, in contrast to mergers of galaxies in less overdense environments.  As discussed below, we suspect that this special assembly history has implications for the scaling relations of massive ellipticals: the preferentially radial orbits of mergers that form the most luminous ellipticals may lead to different Faber-Jackson and stellar mass-size relations for these galaxies than for normal ellipticals.  There are, in fact, tantalizing observational hints of deviations from the canonical elliptical galaxy scaling relations for BCGs and other luminous ellipticals.  For example, \\citet{oegerle1991} find that while BCGs lie on the same \\fpp\\ as lower mass ellipticals, their central velocity dispersions change very little with increasing luminosity and their effective radii increase steeply with luminosity, significantly more so than for normal ellipticals.  A similar effect is present in the early-type galaxy sample of the Nuker team (see fig.~4 of \\citealt{faber1997} and Lauer et al. 2006, in prep.) as well as in the \\citet{postman1995} BCG sample \\citep{graham1996a}, while \\citet{wyithe2006} has claimed to detect an analogous change in the $\\mbh-\\sige$ relation for the most massive black holes.  In this paper, we present the results of a large number of simulations of major mergers of elliptical galaxies.  Our galaxy models contain dark matter haloes and stellar bulges, both simulated with sufficient force and mass resolution to reliably investigate the remnants down to scales of $\\sim 0.1 \\re$ (over $10^6$ total particles per simulation).  We consider a wide variety of energies and angular momenta for the merger orbits and quantify the effects of orbital parameters on the merger remnants.  This work is a significant extension of our earlier related study \\citep[hereafter BMQ]{boylan-kolchin2005}, where we considered only a few possible orbits for the merging galaxies.  The simulations presented in this paper allow us to study in detail the effects of dissipationless merging and merger orbits on the \\fpp\\ and its projections. Section~\\ref{sec:methods} describes the initial galaxy models and simulation parameters.  Section~\\ref{sec:results} contains the results of our simulations, including the locations of merger remnants in the \\fpp\\ (Section~\\ref{sec:fp}) and its projections (Section~\\ref{sec:projections}). In Section~\\ref{sec:discuss} we discuss our results and their implications for the assembly of massive elliptical galaxies and their central black holes. ",
        "conclusions": "\\label{sec:discuss} \\subsection{The Fundamental Plane and its Projections}  The results of Section~\\ref{sec:fp} show that the \\fpp\\ is preserved during dissipationless mergers for a wide range of orbits, a result of virial equilibrium with a slight tilt due to an increasing central dark matter fraction.  On the other hand, the results of Section~\\ref{sec:projections} show that the projections of the \\fpp\\ are quite sensitive to the amount of energy transferred from the bulges to the dark matter haloes during the merger, which is determined in large part by the merger orbit.  If gas-poor mergers help build up the massive end of the red sequence, as several current observations and models suggest, then our results indicate that the projections of the \\fpp\\ -- but not necessarily the plane itself -- provide a powerful way of investigating the assembly history of massive elliptical galaxies.  While a sample of 43 BCGs has been reported to lie on the \\fpp\\ \\citep{oegerle1991}, they appear to have different \\fpp\\ projections from those of lower-mass elliptical galaxies.  The correlation between effective radius and luminosity is steeper for these BCGs ($\\re \\propto L^{1.25}$) than for all early-type galaxies ($\\re \\propto L^{0.6}$) and the BCG velocity dispersions lie substantially lower than the extrapolation of the Faber-Jackson relation to higher luminosity \\citep{oegerle1991}.  Other properties of BCGs are notably different from those of normal ellipticals as well, such as their high incidence of secondary nuclei \\citep{schneider1983}, their more prolate shapes \\citep{ryden1993}, the small dispersion in their metric luminosities \\citep{postman1995}, and the large and relatively constant (out to $z \\sim 0.4$) luminosity gap between BCGs and second-ranked cluster galaxies \\citep{loh2006}.  The results presented in this paper suggest that these differences are all signatures of how massive elliptical galaxies were assembled.  In particular, dissipationless merging of elliptical galaxies provides a natural mechanism for steepening both the Faber-Jackson and $R-L$ relations \\emph{and} preserving the overall \\fpp, \\emph{provided that merger orbits become preferentially more radial for the most massive galaxies} (Fig.~\\ref{fig:rperi-slopes}).  There is, in fact, support for this possibility in the literature. Observationally, it is well-established that BCGs are aligned with both their host clusters and the surrounding large scale structure (e.g. \\citealt{binggeli1982,fuller1999,west2000}).  Numerical simulations show that anisotropic merging is natural in CDM models \\citep{dubinski1998,zentner2005} and preferentially radial merging along filaments can produce many of the observed properties of clusters and BCGs \\citep{west1994,west1995,dubinski1998}.  These lines of reasoning favour radial mergers during BCG formation and point to a strong connection between the characteristics of the mergers that form BCGs and the origin of observed BCG properties (and the differences between BCGs and normal ellipticals).  In addition, the most luminous ellipticals in the Virgo cluster \\emph{all} seem to lie on a filament \\citep{west2000}, which likely means that anisotropic merging is also important for luminous ellipticals that are not BCGs.  According to our calculations, projections of the \\fpp\\ provide independent information about the assembly history of massive ellipticals.  Analysis of a large spectroscopic sample of massive galaxies would provide significantly stronger constraints on changes in the \\fpp\\ projections at the highest masses; we expect that the indications of deviations seen in current samples will become more apparent with additional data.   \\subsection{Supermassive black holes}  Any departure from the canonical Faber-Jackson relation in massive galaxies would also have significant implications for the demography of supermassive black holes.  In particular, under the assumption that black holes coalesce following a gas-poor merger (which is not guaranteed; see \\citealt{merritt2004c} for a review), the black hole mass should increase roughly in proportion to the galaxy mass since the amount of mass-energy carried away in gravitational waves is likely to be quite small (e.g., \\citealt{baker2004}).  We therefore expect that dissipationless mergers will maintain a linear $\\mbh-M_*$ relation.  Any steepening of the $L-\\sigma_c$ relation for luminous ellipticals, however, should then be reflected in a corresponding steepening in the $\\mbh-\\sigma$ relation for massive black holes. \\citet{wyithe2006} has argued that just such a steepening is present in the current sample of black hole masses, but the statistics are rather poor and a larger sample is needed in order to make definitive statements.  If correct, black hole masses for very luminous galaxies based on the canonical $\\mbh-\\sige$ relation \\citep{gebhardt2000, ferrarese2000, tremaine2002} would systematically under-predict the true black hole mass (see also Lauer et al. 2006, in prep.).  \\subsection{Further discussion}  Several issues can be investigated in more detail to test our suggestion that preferentially radial merging is responsible for many of the morphological and kinematical properties of BCGs.  One promising possibility is to study the shapes of BCGs.  \\citet{porter1991} find that the outer regions of BCGs are generally prolate, with the inner regions being rounder.  \\citet{ryden1993} find that the distribution of BCG ellipticities has a similar mean value to that of all ellipticals but with a smaller dispersion and a tendency to be slightly more prolate.  For comparison, we find that the remnants of 1:1 mergers on low angular momentum orbits are prolate with a tendency toward triaxiality: generally $a$:$b$:$c \\sim$ 1.8:1:1 - 1.7:1.1:1.  By contrast, the runs with significant angular momentum have $a$:$b$:$c \\sim$ 1.7:1.3:1.  However, the main difference between our equal mass and 0.33:1 mergers is that the 0.33:1 remnants are notably rounder than their 1:1 counterparts, an effect also seen by \\citet{villumsen1982}.  Unequal mass major mergers thus lead to the same overall trends we see in the \\fpp\\ and its projections while producing more spherical remnants.  Using shapes to quantitatively test the predictions of our models will thus require detailed information about the orbits and mass ratios of the mergers that formed BCGs.  The above considerations highlight the need to move beyond non-cosmological major merger simulations of spherical and isotropic galaxies that have been the focus of this paper as well as many earlier papers.  Indeed, the relaxation and virialization of a cluster generally involves several massive galaxies with a range of masses \\citep{dubinski1998}.  The hierarchical build-up of clusters and BCGs also likely depends on the cluster mass (e.g.  \\citealt{brough2005, loh2006}).  In addition, the shapes of cluster-mass dark matter haloes in dark matter-only cosmological simulations are typically quite prolate at formation but then become rounder (but still triaxial) as time goes on \\citep{hopkins2005, allgood2005} due to a transition from filamentary to spherical accretion \\citep{allgood2005}.  Clearly, high resolution two-component (stars and dark matter) simulations in a cosmological context with full formation histories are necessary for making detailed predictions about the properties of massive ellipticals.  We are currently working on such simulations.  Although we expect that our basic conclusions about the \\fpp\\ and its projections will remain unchanged, repeated major and minor merging on realistic cosmological orbits will significantly influence the shapes and kinematics of the remnant galaxies.  As one final implication of the non-spherical nature of our merger remnants, we note that the measured velocity dispersions can change by more than 15\\% with viewing angle for a given remnant.  For the most massive galaxies, this could be significant: a galaxy that has $\\sige \\approx 350$ km/s from most viewing angles (typical for a BCG) could have $\\sige > 400$ km/s in certain rare configurations (we find that the observed velocity dispersion is $\\approx \\, 15 \\%$ higher than the mode $\\approx 3 \\%$ of the time for the remnants of relatively radial mergers).  These objects, which are prolate merger remnants viewed down the major axis, also have smaller effective radii than most viewing angles of remnants with the same stellar mass.  \\citet{bernardi2005} have recently reported the detection of a population of luminous galaxies with these properties.  A simple test of our explanation for these galaxies is to look at their projected shapes: our calculations predict these galaxies should appear statistically rounder than other galaxies of comparable luminosity.  This viewing angle effect would also be a source of scatter in the $\\mbh-\\sige$ and Faber-Jackson relations.   \\vspace{4mm} We thank Sandy Faber, Alister Graham, Zoltan Haiman, and Tod Lauer for useful discussions and Volker Springel for making {\\tt GADGET} publicly available.  This work used resources from NERSC, which is supported by the US Department of Energy.  C-PM is supported in part by NSF grant AST 0407351 and NASA grant NAG5-12173.  EQ is supported in part by NSF grant AST 0206006, NASA grants NAG5-12043 and ATP05-54, an Alfred P. Sloan Fellowship, and the David and Lucile Packard Foundation."
    },
    "0601/hep-th0601037_arXiv.txt": {
        "abstract": "We study inflation arising from the motion of a Bogomol'nyi-Prasad-Sommerfield (BPS) D3-brane in the background of a stack of $k$ parallel D5-branes. There are two scalar fields in this set up-- (i) the radion field $R$, a real scalar field, and (ii) a complex tachyonic scalar field $\\chi$ living on the world volume of the open string stretched between the D3 and D5 branes. We find that inflation is realized by the potential of the radion field, which satisfies observational constraints coming from the Cosmic Microwave Background. After the radion becomes of order the string length scale $l_s$, the dynamics is governed by the potential of the complex scalar field. Since this field has a standard kinematic term, reheating can be successfully realized by the mechanism of tachyonic preheating with spontaneous symmetry breaking. ",
        "introduction": "There has been a resurgence of interest in the time-dependent dynamics of extended objects found in the spectrum of string theory, inspired in part by Sen's construction of a boundary state description of open string tachyon condensation. See, for example, Ref.~\\cite{senrev} for review. This description has been supplemented well by an effective theory described by a Dirac Born Infeld (DBI) type action for the tachyon field \\cite{DBI}. More recent work has focused on the dynamics of a probe BPS D-brane in a variety of gravitational backgrounds inspired by the observation that there exists a similarity between the late time dynamics of the probe D-branes and the condensation of the open string tachyon on the world-volume of non-BPS brane in flat space. The latter dynamics is also described by the DBI action \\cite{kutasov}, see also Refs.~\\cite{Sahakyan,pani,TW}. Both systems describe rolling matter fields which have a vanishing pressure at late times. As a result we can, through an appropriate field transformation, investigate the physics of gravitational backgrounds in terms of non-trivial fields on a brane in flat space using the DBI effective action. This has led to the interesting proposal that the open string tachyon may be geometrical in nature.  Many of the backgrounds that have been probed in this manner have been supergravity (SUGRA) brane solutions of type II string theory. By ensuring that the number of background branes is large we can trust our SUGRA solutions. Moreover we can neglect any back reaction of the probe upon the background geometry. This allows us to use the DBI action to effectively determine the relativistic motion of extended objects in a given background. Quantum corrections can also be calculated in those backgrounds that have an exact Conformal Field Theory (CFT) description \\cite{Sahakyan}. The dynamics of branes in various backgrounds is expected to be relevant for string theory inspired cosmology, just as in the case of open string tachyon matter \\cite{tachinfl} since the field (radion) which parameterizes  the distance between the probe brane and the static background branes is a scalar and may be a potential candidate for being the inflaton.  One of the most important theoretical advances in modern cosmology has been the inflationary paradigm, which relies on a scalar field to solve the horizon and flatness problems in the early universe (see Refs.~\\cite{inflation} for review). Recent observations from WMAP \\cite{WMAP}, SDSS \\cite{SDSS} and 2dF \\cite{2dF} impose tight restrictions on the possible mechanisms that can satisfy the paradigm \\cite{general}, and hence provide the interesting possibility for us to test string theoretic inflation models. The observations of Supernova Ia \\cite{SNIa} also suggest that our universe is currently undergoing a period of accelerated expansion, which is attributed to dark energy. It still remains a fundamental problem to describe dark energy in a purely stringy context, although there has been several recent developments \\cite{darkenergy}.  There have been many attempts to embed inflation within string theory. The most popular approach has been to invoke the use of the open string tachyon living on a non-BPS brane as a candidate for the inflaton \\cite{tachinfl} (see Refs.~\\cite{cospep} for a number of cosmological aspects of tachyon). Unfortunately it has been shown that this cannot be implemented in a consistent manner, at least in the simplest scenarios \\cite{KL02,Frolov}. The other common approach is so-called D-brane inflation in which the separation between branes plays the role of the inflation \\cite{Dvali,braneantibrane,Garcia,Hirano}. In particular this is well accommodated in a form of hybrid inflation where tachyonic open string fluctuations are the fields which end inflation, and another field is chosen to be the inflaton. These open string fluctuations arise in the context of all D-brane cosmological models once the branes are within a string length of one another. A concrete example of this occurs in brane/anti-brane inflation \\cite{braneantibrane} and recently in the context of more phenomenological warped compactifications \\cite{kklmmt} (see also Refs.~\\cite{dbranepapers}). It should be noted that most of the work done in this direction assumes that the dimensionalities of the brane and anti-brane are the same apart from D3/D7 brane inflation models studied in Refs.~\\cite{D3D7} which does not include the open string tachyon dynamics at late times. On the contrary, in our model a probe D3 brane is used to lead to inflation in the presence of static D5 branes (see Refs.~\\cite{TWinf,PST} for related works) and the open string tachyon dynamics naturally comes in. In any event there has been very little work done on trying to understand the relationship between inflation and the current dark energy phase which we observe.  A potential solution for both inflation and dark energy, in this context, can be obtained as a mixture of these two scenarios. We require a mechanism which drives inflation independently of the open string tachyon, but then falls into the tachyonic state at late times. This can be achieved by considering the motion of a D3-brane in a type IIB background. By switching to our holographic picture of a non-trivial field on a non-BPS brane \\cite{kutasov} we will find that the radion field naturally exits from inflation once it reaches a critical velocity. If this occurs at a distance larger than the string length, we can then use the open string tachyon, which sets in at a distance equal to or less than the string length, to explain the dark energy content of the universe.  In this paper we aim to explore the motion of a probe D3-brane in the background of $k$ coincident, static D5-branes. For simplicity we will neglect any closed string radiation which would be emitted from the probe brane as it travels down the throat generated by the background branes. We will also neglect any gauge fields which may exist on the D3-brane world-volume. Note that this is S-dual to the solution considered in Ref.~\\cite{kutasov}. In order to make contact with four dimensional physics we must consider the dual picture of a non-trivial field on a non-BPS brane in flat space, where we also toroidally compactify the remaining six dimensions\\footnote{This is not necessary if we consider holographic cosmology as in Ref.~\\cite{mirage}.}. We will assume that there is some mechanism which freezes the various moduli of the compactification manifold so that they do not appear in the effective action. The resulting theory should represent the leading order contribution which would arise from compactifying the full type IIB background. At distances large compared to the string scale, the DBI description is known to be valid, however once the probe brane approaches small distances (order of string length scale) we must switch to the open string analysis. Open strings will stretch from the D3-brane to the D5-branes, and their fluctuation spectrum contains a tachyonic mode. Thus when the separation is order of the string length, the DBI description will no longer be valid and we must resort to a purely open string analysis. We expect that inflation will occur in the large field (radion) regime and it ends before the separation comes closer to string length and that as the branes get closer, the open string tachyon reheats the universe.  This paper is organized as follows. In the next section we describe the dynamics of a single probe D3-brane in the presence of a large number of static background D5-branes. Because of the dimensionalities of the branes we expect to find an open string tachyonic mode once we begin to probe distances approaching the string length \\cite{branonium}. In section III, we present the inflation dynamics and observational constraints on the various parameters of our model. In section IV, we discuss the role of the open string tachyon after the inflationary phase and the possibility of reheating in our model and a brief discussion on dark energy. In the last section, we present some of our conclusions and future outlook. ",
        "conclusions": "In this paper we studied the motion of a BPS D3-brane in the presence of a stack of $k$ parallel D5-branes in type II string theory. Inflation is realized by the potential energy of a radion field $R$ which characterizes the distance of D3 and D5 branes. This potential is not in general written explicitly, but is approximately given by (\\ref{pot2}) for $R \\ll \\sqrt{kg_s}l_s$ and (\\ref{pot}) for $R \\gg \\sqrt{kg_s}l_s$. We evaluated the spectral index of scalar metric perturbations and the tensor-to-scalar ratio together with the number of $e$-foldings under the condition that inflation ends in the region $R \\ll \\sqrt{kg_s}l_s$. This model satisfies observational constraints coming from CMB, SDSS and 2dF independently of the value of $s$ defined by Eq.~(\\ref{sdef}). We also note that this result does not change even when inflation ends in the region $R \\gg \\sqrt{kg_s}l_s$.  The only strong constraint coming from CMB is the COBE normalization, i.e., $kg_{s}(l_{s}M_{p})^2 \\simeq 10^{10}$ for $s \\gtrsim 1$. If we demand that the inflationary period is over before the radion reaches the point $R=l_s$, this gives the constraint on the number of D5-branes; see Eq.~(\\ref{nseq3}). Combining this with the condition of the COBE normalization, the string mass scale is constrained to be $M_s/M_p \\ll 0.1$ for the validity of the weak-coupling approximation ($g_s \\ll 1$).  When the radion field enters the region $R \\lesssim l_s$, the description of closed string background is no longer valid. Instead the dynamics should be studied using a complex scalar field $\\chi$ living on the world volume of open strings stretched between the probe D3-brane and the D5-branes. We assumed that the radion field is frozen in the region $R \\lesssim  l_{s}$, which gives rise to a positive cosmological constant. The potential of the field $\\chi$ is given by Eq.~(\\ref{chipot}), which has a negative energy at the potential minimum. If this energy is cancelled by the positive cosmological constant, we obtain the double-well potential given by Eq.~(\\ref{dwell}).  We found that the absolute value of the mass of this double-well potential at $\\chi=0$ is larger than the Hubble parameter provided that $M_s/M_p>4 \\times 10^{-8}$. Hence in this case the second stage of inflation does not occur and the evolution of the field $\\chi$ is described by a fast roll. Since the action of the complex field has a standard kinematic term, the problem of reheating present in DBI tachyon models is absent in our scenario. Reheating in our model is described by tachyonic preheating in which quantum fluctuations grow exponentially by a negative instability. The symmetry breaking would end after one oscillation of the field distribution as the field evolves toward the potential minimum.  It is also possible to explain the origin of dark energy if a positive cosmological constant does not exactly cancel the negative potential energy of the field $\\chi$. Although this requires a fine tuning, it is intriguing that our model provides a number of promising ways to provide viable cosmological evolution.  One of the potential problems with our model is that the compactification is not necessarily realistic. Although we can encode the physics of the gravity background as a non-trivial scalar field on a flat brane, we are treating this latter object as being fundamental. Thus by compactifying this on a $T^6$ we will be missing higher order terms coming from the full compactification of the D5-solution. These terms may play a more important role in the cosmological theory on the D3-brane. It may be useful to compare the results obtained in this paper with a full string compactification by smearing the SUGRA harmonic function on a $T^4$ and compactifying the remaining directions on the two-cycles of a torus. The resultant analysis is complicated since the DBI may not be valid, however this is beyond the scope of the current endeavor."
    },
    "0601/astro-ph0601120_arXiv.txt": {
        "abstract": "We study the dependence of the galaxy contents within halos on the halo formation time using two galaxy formation models, one being a semianalytic model utilizing the halo assembly history from a high resolution {\\it N}-body simulation and the other being a smoothed particle hydrodynamics simulation including radiative cooling, star formation, and energy feedback from galactic winds. We confirm the finding by Gao et al. that at fixed mass, the clustering of halos depends on the halo formation time, especially for low-mass halos. This age dependence of halo clustering makes it desirable to study the correlation between the occupation of galaxies within halos and the halo age. We find that, in halos of fixed mass, the number of satellite galaxies has a strong dependence on halo age, with fewer satellites in older halos. The youngest one-third of the halos can have an order of magnitude more satellites than the oldest one-third. For central galaxies, in halos that form earlier, they tend to have more stars and thus appear to be more luminous, and the dependence of their luminosity on halo age is not as strong as that of stellar mass. The results can be understood through the star formation history in halos and the merging of satellites onto central galaxies. The age dependence of the galaxy contents within halos would constitute an important ingredient in a more accurate halo-based model of galaxy clustering. ",
        "introduction": "Two main simulations are used in our study. The one that the SAM is based on is a $\\pppm$ $N$-body simulation in \\citet{Jing02} with $512^3$ particles in a cubic box of 100 $\\mpc$. A spatially-flat $\\Lambda$CDM model is adopted with density parameters $\\Omega_m=0.3$ and $\\Omega_\\Lambda=0.7$. We adopt the cold dark matter power spectrum of \\citet{Bardeen86} with a shape parameter $\\Gamma=\\Omega_mh=0.2$ and the amplitude $\\sigma_8=0.9$. The softening length is $10 \\kpc$, and the particle mass is $M_p \\sim 6.2 \\times 10^{8}\\msun$. We denote this simulation as SAM-$\\Lambda$CDM. The galaxies of the SAM sample are generated by the model of \\citet{Kang05} that explicitly follows the assembly history of each halo in the simulation. The model successfully explains a wide variety of observational facts, including luminosity functions in different bands and the bimodal color distribution of galaxies.  The other simulation, denoted as HYD-Gadget2, is performed as smoothed particle hydrodynamics with the TREE-PM code Gadget2 \\citep{Springel05,Springel01}, following the evolution of $512^3$ dark matter particles and $512^3$ gas particles in a cubic box of 100 $\\mpc$.  It includes the physical processes of radiative cooling, star formation, and energy feedback from galactic winds \\citep{sh03}. The cosmological and initial density fluctuation parameters of the model are ($\\Omega_m,\\Omega_\\Lambda,\\Omega_b,\\sigma_8,n,h$)=($0.268, 0.732, 0.044,0.85,1,0.71$), and the transfer function is computed from CMBFAST \\citep{Seljak96}. The softening length is also $10 \\kpc$, and the particle mass (dark matter+gas) is $M_p\\sim 5.54 \\times 10^{8} \\msun$.  Galaxies are identified as gravitationally bound groups of stars and cold gas particles that are associated with a common local maximum in the baryon density. We note that the effect of dust extinction, which is included in the SAM model, is not taken into account in this model.  In addition to the two main simulations, we also use a low-resolution simulation of $512^3$ particles in a box of 300$\\mpc$ only for the purpose of probing the age dependence of halo clustering at high halo mass. This simulation (denoted as 300-$\\Lambda$CDM) has the same cosmological parameters as SAM-$\\Lambda$CDM, and the particle mass is $M_p \\sim 1.67 \\times 10^{10} \\msun$.  In all the simulations, halos are identified with the friends-of-friends algorithm \\citep{Davis85} with a linking length $0.2$ times the mean particle separation. ",
        "conclusions": "In this Letter, we investigate the dependence of the galaxy contents on the halo formation time using two galaxy formation models, one with a semi analytic method utilizing the realistic halo assembly history from $N$-body simulations and the other with a smoothed particle hydrodynamics simulation including radiative cooling, star formation, and energy feedback from galactic winds. Even though our techniques for building merger trees as well as our simulation methods are independent, we confirm the dependence of halo clustering on the formation time reported by \\citet{Gao05}. We study the CLF and CSMF of galaxies as a function of the halo formation time and find that galaxies inside halos have a strong correlation with the halo formation time. In halos that form earlier, there are fewer satellite galaxies, and their central galaxies appear to be more luminous, which can be understood through the early onset of star formation and the merging of satellites onto central galaxies. The dependence of the halo occupation number of satellite galaxies on the formation time in this study is similar to that of subhalos presented in \\citet{Wechsler05} (also see \\citealt{Zentner05}), which is not surprising since satellite galaxies reside in subhalos.  The HOD or CLF model assumes that the {\\it statistical distribution} of galaxies inside halos depends only on halo mass. However, we note that this assumption does not mean that at fixed halo mass, the occupation number is independent of halo age. The correlation between the occupation number and the halo formation time exists even if the halo assembly history is based on the excursion set theory (e.g., \\citealt{Zentner05}), and the distribution of the halo age at fixed mass could provide the scatter that goes into the probability distribution of the occupation number at fixed halo mass [e.g., $P(N|M)$ in the HOD model]. It is the dependence of halo clustering on halo age that makes the above correlation potentially important in interpreting galaxy clustering data. Given the age-dependent halo clustering, the age dependence of HOD/CLF constitutes a key ingredient for a more accurate halo-based model of galaxy clustering. In view of this, a detailed study of the correlation between galaxy contents and halo formation time, like that performed in this Letter, is desired.  Depending on the application, the current version of the HOD/CLF model without taking into account the non-Markovian nature of halo formation, can still give accurate descriptions (e.g., \\citealt{Yoo05}). We reserve a detailed investigation for future work, to see to what extent it is adequate in different applications."
    },
    "0601/astro-ph0601316_arXiv.txt": {
        "abstract": "The \\oii\\ \\lamb3727 emission line is often used as an indicator of star formation rate in extragalactic surveys, and it can be an equally effective tracer of star formation in systems containing luminous active galactic nuclei (AGNs). In order to investigate the ongoing star formation rate of the host galaxies of AGNs, we measured the strength of \\oii\\ and other optical emission lines from a large sample ($\\sim 3600$) of broad-line (Type 1) AGNs selected from the Sloan Digital Sky Survey.  We performed a set of photoionization calculations to help evaluate the relative contribution of stellar and nonstellar photoionization to the observed strength of \\oii. Consistent with the recent study of Ho (2005), we find that the observed \\oii\\ emission can be explained entirely by photoionization from the AGN itself, with little or no additional contribution from \\hii\\ regions.  This indicates that the host galaxies of Type 1 AGNs experience very modest star formation concurrent with the optically active phase of the nucleus.  By contrast, we show that the sample of ``Type 2'' quasars selected from the Sloan Digital Sky Survey does exhibit substantially stronger \\oii\\ emission consistent with an elevated level of star formation, a result that presents a challenge to the simplest form of the AGN unification model. ",
        "introduction": "It is now almost universally accepted that massive black holes are common and that they play an important role in many facets of galaxy evolution (see reviews in Ho 2004).  Numerous theoretical models have been proposed to explain the strong observed correlations between black hole mass and bulge properties (Magorrian et al. 1998; Gebhardt et al. 2000; Ferrarese \\& Merritt 2000).  While everyone agrees that the growth of black holes must be closely linked with galaxy formation (e.g., Silk \\& Rees 1998; Kauffmann \\& Haehnelt 2000; Begelman \\& Nath 2005; Robertson et al.  2006), there is no consensus as to exactly how the two very different processes involved---accretion and star formation---are really coupled.  Are they well synchronized, or does one process precede the other, and if so, what is the time lag?  When the black hole is actively growing, does the feedback from the active galactic nucleus (AGN) actually trigger or inhibit star formation in the host galaxy? These important issues are unlikely to be settled through theoretical speculations or numerical simulations alone.  Some empirical guidance from observations would be welcomed.  In recent years, there has been mounting evidence that AGN activity and starburst activity are often intermixed.  The most commonly cited example occurs in ultraluminous infrared galaxies, whose dominant energy source has been much debated (e.g., Genzel et al.  1998).  Detailed studies of individual active galaxies have revealed a significant population of young stars in a number of instances, either through stellar absorption features (e.g., Boisson et al. 2000; Canalizo \\& Stockton 2000; Cid Fernandes et al. 2001) or colors (e.g., Jahnke et al. 2004).  Additional hints for a close connection between star formation and black hole accretion comes from the high degree of chemical enrichment in quasars as deduced from analysis of their spectra (Hamann et al. 2004 and references therein). Lastly, the host galaxies of quasars, at least in the nearby Universe, possess a rich supply of both molecular gas (Scoville et al. 2003) and dust (Haas et al. 2003), and thus have the potential for sustaining ongoing or future star formation.  While there is little doubt that the above observations indicate that AGN activity and star formation do often coincide within the same galaxy, one of the main challenges in interpreting this evidence in a physical context lies in the difficulty of establishing a direct, {\\it causal}\\ connection between the two phenomena.  An important step forward was recently achieved.  From an analysis of a large sample of narrow-line (Type 2) AGNs selected from the Sloan Digital Sky Survey (SDSS; York et al. 2000), Kauffmann et al. (2003) showed that, unlike low-luminosity sources (Ho et al. 2003), more powerful AGNs often exhibit stellar absorption features indicative of young to intermediate-age stars.  Although more difficult to ascertain due to the strong contamination from the bright nucleus, Kauffmann et al. show that the spectral signatures for young stars seem to persist even among broad-line (Type 1) AGNs.\\footnote{We will use the terms ``Type 1'' and ``Type 2'' AGNs to refer to broad-line and narrow-line objects, respectively.  Whenever possible, we will refrain from using the terms ``Seyferts'' or ``quasars,'' which, for historical reasons, carry artificial luminosity criteria.} Most crucially, they find, for their primary sample of Type 2 objects, that the fraction of young stars in the central region of the host galaxies increases with increasing AGN luminosity, providing, for the first time, tantalizing evidence for a causal link between star formation and accretion.  It is important to recognize, however, that the SDSS results constrain only the {\\it post-starburst}\\ population, with ages $\\sim 10^8-10^9$ yr.  There is no information on the presence of younger, ionizing stars (ages \\lax$10^7$ yr). Since the Type 2 AGN sample was identified using standard optical emission-line ratios designed to isolate accretion-dominated sources, {\\it by selection}\\ it excludes any sources with a sizable contribution from \\hii\\ regions. On the other hand, it is clearly of significance to constrain the ongoing star formation rate (SFR) in AGNs, in order to assess the extent to which the two processes are coeval.  AGN lifetimes are uncertain, but current best estimates range from $10^6$ to $10^8$ yr (Martini 2004), precisely in the unexplored regime of interest.  Estimating SFRs in AGNs presents significant complications, since the AGN, by definition, dominates the integrated output of the source.  Nearly all of the standard measures of SFRs employed for inactive galaxies (e.g., Kennicutt 1998) are rendered useless by emission from the AGN itself.  There is, however, one important exception: the \\oii\\ \\lamb3727 line.  As discussed by Ho (2005), \\oii, an emission line commonly used in spectroscopic surveys to estimate SFRs in galaxies at redshifts $z$ \\gax\\ 0.4, should be an equally effective tracer of \\hii\\ regions in Seyferts and quasars.  In such high-ionization active galaxies, the \\oii\\ emission intrinsic to the narrow-line region (NLR) of the AGN is both observed and predicted to be low (e.g., Ferland \\& Osterbrock 1986; Ho et al. 1993b).  If high-ionization AGNs experience substantial levels of ongoing star formation, the integrated contribution from \\hii\\ regions will boost the strength of the \\oii\\ line (compared to, say, \\oiii\\ \\lamb 5007, which can be largely ascribed to the AGN itself).  From a survey of the existing literature on \\oii\\ measurements on quasars, Ho (2005) concludes that their SFRs are surprisingly modest. More intriguingly, for a subset of low-redshift quasars with available molecular gas measurements, Ho finds that their star formation efficiencies (SFR per unit gas mass) appear to be abnormally low compared to either normal or infrared-bright galaxies.  He suggests that this unusual behavior may be a manifestation of AGN feedback suppressing star formation.  The results of Ho (2005) regarding SFRs in quasars were based largely on \\oii\\ strengths deduced from composite spectra (ensemble averages) constructed from quasar surveys.  While the conclusions are statistically robust, it would be desirable to revisit this problem from actual measurements of a large, well-defined sample of individual objects.  This is accomplished in this paper, where we make use of an extensive sample of Type 1 AGNs selected from the SDSS to investigate the properties not only of \\oii, but also a number of other diagnostically important narrow optical emission lines.  We perform a new set of photoionization calculations to demonstrate that the observed strength of \\oii\\ emission is entirely consistent with standard AGN photoionization.  These results support and strengthen the conclusion of Ho (2005) that Type 1 AGNs experience very limited ongoing star formation.  By contrast, we show that Type 2 quasars exhibit markedly enhanced \\oii\\ emission, and hence presumably elevated SFRs.  This result seems to violate the basic orientation-dependent unification model for AGNs.  We adopt the following cosmological parameters: $H_0 = 100\\,h = 71$ \\kms~Mpc$^{-1}$, $\\Omega_{\\rm m} = 0.27$, and $\\Omega_{\\Lambda} = 0.75$ (Spergel et al. 2003). ",
        "conclusions": "\\subsection{AGN Feedback}  Ho (2005; see also Ho 2006) recently proposed that the \\oii\\ \\lamb3727 emission line can effectively trace ongoing star formation in the host galaxies of luminous, high-ionization AGNs such as Seyferts and quasars.  The \\oii\\ line has long been used as a SFR indicator in spectroscopic surveys of distant galaxies, but previously it had not been applied to systems containing active nuclei.  From a review of the extant literature on \\oii\\ measurements in quasars, Ho came to the conclusion that the SFRs in their host galaxies are quite modest, being no greater than $\\sim 1-10$ \\solmass\\ \\peryr.  This result was somewhat unexpected considering the recent evidence for there being a substantial post-starburst population in AGNs (Kauffmann et al. 2003).  Perhaps even more surprising, Ho found, from a more limited sample of low-redshift quasars with available CO measurements, that not only are the SFRs low but that the star formation efficiencies (SFR per unit gas mass) are also low, suggesting that the AGN somehow {\\it inhibits}\\ star formation.  Apart from the small sample of low-redshift objects that had individual spectroscopic data, Ho's more general conclusions regarding the quasar population as a whole was based on published measurements of \\oii\\ strengths derived from composite, or statistically averaged, spectra.   To place these results on a more secure footing, we have undertaken a major task to homogeneously measure the \\oii\\ line in a large sample of Type 1 AGNs.  Apart from \\oii, we also simultaneously measured a number of other diagnostically important narrow emission lines, and we have generated a new set of photoionization models to interpret them.  Our analysis indicates that the principal narrow emission lines of our sample, including \\oii, can be entirely explained by AGN photoionization assuming fairly standard parameters.  There is no evidence for any excess \\oii\\ emission arising from a separate origin, such as \\hii\\ regions.  Thus, the observed \\oii\\ strengths can be used to set a limit on the ongoing SFR in these systems.   However, as discussed by Ho, a number of uncertainties complicate the SFR estimates based on the \\oii\\ line.  These include the unknown magnitude of dust extinction, the possible influence of metallicity corrections, and precisely how much of the observed \\oii\\ strength to attribute to star formation. For illustrative purposes, we make three simple assumptions (see  Ho 2005 for a more detailed discussion): (1) that the amount of dust extinction in AGN host galaxies is comparable to that deduced for moderately actively star-forming galaxies ($A_V \\approx 1$ mag; e.g., Hopkins et al. 2003); (2) that the metallicity is twice solar (e.g., Storchi-Bergmann et al. 1998); and (3) that one-third of the observed \\oii\\ strength comes from \\hii\\ regions.  For a typical \\oii\\ luminosity of $5.2 \\times 10^{40}$ \\lum\\ (Fig. 4{\\it a}), the empirical calibration of Kewley et al. (2004) yields a SFR of merely 0.5 \\solmass\\ \\peryr.  This estimate is consistent with, but even lower than, than that given by Ho (2005).  To put this estimate in perspective, we note that the SFRs of nearby normal spiral galaxies, including the Milky Way, range from $\\sim$1 to 3 \\solmass\\ \\peryr\\ (Solomon \\& Sage 1988). The situation in AGNs is thus truly unusual.  Unlike the low-redshift quasars considered by Ho (2005), our current sample of SDSS AGNs has no available information on its cold gas content.  However, the unbiased CO survey by Scoville et al. (2003) shows that optically selected quasars, not unlike those studied here, are quite gas-rich.   This is not unexpected, considering that the nuclei, by selection, are quite active, and therefore accreting at a fairly high rate.  Thus, although the evidence is certainly less direct, we can reasonably surmise that the objects studied here also may also possess abnormally low star formation efficiencies.  How can this finding be reconciled with Kauffmann et al.'s (2003) result that luminous AGNs show a preponderance of post-starburst activity?  A possible solution is to invoke an evolutionary scenario, not unlike that proposed by Sanders et al. (1988), in which a starburst phase {\\it precedes}\\ the optically revealed quasar phase.  From the age estimates given in Kauffmann et al., the time lag seems to be approximately $10^8$ to $10^9$ years.  The starburst phase terminates not because all of the gas has been transformed into stars, nor because it has been completely expelled by AGN or supernova feedback, since evidently plenty of gas coexists during the quasar phase, but because the active nucleus somehow prevents the gas from forming stars. Through a combination of radiative heating and kinetic energy injection, an AGN is expected to significantly alter the thermodynamical state of the gas in the circumnuclear regions of galaxies (Begelman 1993; Di~Matteo et al. 2005), but precisely how this results in suppression of star formation remains to be elucidated.   \\begin{figure*}[t] \\centerline{\\psfig{file=table1.ps,keepaspectratio=true,angle=0}} \\end{figure*} \\vskip 5mm    \\subsection{Possible Evidence for Starburst Activity in Type 2 Quasars}  Zakamska et al. (2003) have drawn attention to a class of narrow-line AGNs from the SDSS whose \\oiii\\ line strength, when translated into continuum luminosities, place them in the league of quasars.  This, along with the detection in some sources of polarized broad-line emission through spectropolarimetry (Zakamska et al. 2005) and heavy photoelectric absorption through X-ray observations (Ptak et al. 2006), has prompted the identification of these objects with the long-sought ``Type 2 quasars'' anticipated from the the basic orientation-dependent unification model of AGNs.   If the Zakamska et al. objects truly are intrinsically identical to normal, high-luminosity Type 1 AGNs, but simply viewed from a line of sight obscured from the nucleus, then the narrow-line spectrum of the two groups, assumed to be isotropic, should be similar.  We were surprised to discover that this appears not to be the case.  Comparison of the line strengths for our Type 1 objects with those tabulated for the composite SDSS spectrum of Zakamska et al. (2003) reveals that both groups have similar line intensity ratios, with two noticeable exceptions (see Table 1).  First, the Type 2 sources have a higher Balmer decrement, and hence a higher inferred internal extinction for their NLR, compared to the Type 1 sources. Recall from our spectral fitting (\\S3) that, when allowed to vary freely, the majority of our sources yielded an \\hal/\\hb\\ ratio of 3.3, to be compared with \\hal/\\hb\\ = 4.1 for the Type 2 composite.  If interpreted in terms of dust extinction assuming an intrinsic Balmer decrement of 3.1 and a Galactic extinction curve, this translates into an internal extinction of $A_V = 0.2$ mag versus $A_V = 0.9$ mag, respectively. Second, and more pertinent to the central theme of this study, we find that the Type 2 quasars exhibit an anomalously enhanced (extinction corrected) \\oii/\\oiii\\ ratio.  This point is illustrated in Figures 6 and 7, where we highlight the location of the Zakamska composite in relation to the locus of the Type 1 sources.  While the relatively high \\oii/\\oiii\\ ratio can normally be attributed to a somewhat lower ionization parameter, we note that this explanation is in serious conflict with the high luminosity of these sources.  As discussed in \\S 4.2, the \\oii/\\oiii\\ ratio varies strongly with \\oiii\\ luminosity.  For $L_{\\rm [O~{\\sc III}]} \\approx 3\\times 10^{42}$ \\lum, the average for the Zakamska et al. sample, Figure 5 predicts that the \\oii/\\oiii\\ ratio should be $\\sim 0.06-0.1$.  By contrast, the Type 2 quasar composite gives \\oii/\\oiii\\ = 0.75, nearly an order of magnitude larger than the predicted value.  We speculate that the enhanced \\oii\\ emission in the Type 2 quasars comes from star formation.  If the observed \\oii/\\oiii\\ ratio in Type 2 quasars is $\\sim$10 times larger than expected for their \\oiii\\ luminosity, then the excess \\oii\\ luminosity of $\\sim 2\\times 10^{42}$ \\lum, adopting the assumptions specified in \\S5.1, corresponds to an estimated SFR of $\\sim 20$ \\solmass\\ \\peryr.  This level of star formation qualifies as a moderately healthy starburst.  Incidentally, we note that a dusty, gas-rich environment characteristic of starbursts may naturally account for the higher internal extinction mentioned above.  The possible existence of significant ongoing star formation in Type 2 quasars, in contrast to their stark absence in normal (Type 1) quasars, presents an obvious challenge to the conventional orientation-dependent model of AGN unification (Antonucci 1993).  Type 2 objects are not intrinsically the same as Type 1 objects, but instead may represent an earlier evolutionary phase, perhaps similar to the scenario discussed in \\S5.1, wherein the central engine is either not yet fully developed or else lies buried by extended, dusty regions of star formation.  In this context, the optically selected Type 2 quasars closely resemble those ultraluminous infrared galaxies that host both a strong starburst and a highly obscured AGN (e.g., Genzel et al. 1998; Gerssen et al. 2004).  Observations of radio-loud AGNs have suggested that the \\oiii-emitting region, unlike that producing \\oii, may not be entirely isotropic (Jackson \\& Browne 1990; Hes et al. 1996). If this is generically true of all AGNs, then the \\oii/\\oiii\\ ratio is orientation-dependent, and it may account for the higher ratios in Type 2 quasars compared to Type 1 quasars.  This explanation, however, is unsatisfactory because Figure 5 shows no evidence that this pattern holds for lower-luminosity sources."
    },
    "0601/astro-ph0601585_arXiv.txt": {
        "abstract": "{We present a powerful new tool to analyse and disentangle the 3-D geometry and kinematic structure of gaseous nebulae. The method consists in combining commercially available digital animation software to simulate the 3-D structure and expansion pattern of the nebula with a dedicated, purpose built rendering software that produces the final images and long slit spectra which are compared to the real data. We show results for the complex planetary nebulae NGC~6369 and Abell~30 based on long slit spectra obtained at the San Pedro M\\'artir observatory. }  \\resumen{Presentamos una nueva herramienta poderosa para analizar y desenredar la geometr\\'{\\i}a 3-D y la estructura cinem\\'atica de nebulosas gaseosas. El m\\'etodo consiste en combinar software comercial para animaci\\'on digital para simular la estructura 3-D y el modo de expansi\\'on de la nebulosa junto con un software de representaci\\'on gr\\'afica de imagenes y perfiles de l\\'{\\i}nea dise\\~nado especialmente para este prop\\'osito. Las im\\'agenes resultantes se comparan directamente con datos reales. Presentamos resultados para las complejas nebulosas planetarias NGC~6369 y Abell~30 basados en espectros de rendija larga obtenidos en el observatorio de San Pedro M\\'artir.}  \\addkeyword{ISM: modeling} \\addkeyword{Planetary nebulae: individual (NGC6369), individual (Abell 30)} \\addkeyword{Stars: Mass loss}   \\begin{document} ",
        "introduction": "\\label{sec:introduction}  In recent years, the discovery of a variety of complex structures in planetary nebulae has opened many questions regarding the origin and evolution of these objects (e.g. L\\'opez 2000). Deviations from simple expanding shells can include collimated outflows, poly-polar and point-symmetric structures, rings or disks. These observations have lead to a wealth of theoretical research into the effects of stellar magnetic fields, rapid rotation and binarity of the central stars, and their evolutionary path from spherically symmetric to bipolar mass-loss (e.g. Balick \\& Frank, 2002, and references therein). In the absence of spherical symmetry, the tilt of the nebula with respect to the line of sight and the location and position angle of the slit on the nebula can often result in complicated position-velocity (P-V) diagrams that can be difficult to interpret. The correct interpretation of the nebular 3-D geometry and kinematic structure of PNe is key to the understanding of the dynamics ruling their origin and evolution.  Modeling of line emission intensity maps have been used to obtain density distributions over the face of the nebula in order to assess 3-D structures, assuming pure photoionization from the central star (e.g. Monteiro et al. 2004; Morisset, Stasi\\'nska \\& Pe\\~na, 2005) but without incorporating kinematic information or assuming simple velocity laws (e.g. Ragazzoni et al. 2001).  In this paper we present a new interactive 3-D modeling tool called \\shape which combines the versatility of commercial 3-D modeling software with a rendering module specifically developed for application in astrophysical research. The application of this method yields a 3-D emissivity and velocity distribution for the object. Furthermore, different velocity laws can be applied to different sections of the nebula to reproduce the complex velocity patterns often observed in PNe. We exemplify the power of this new method with models of the particularly complex planetary nebulae Abell 30 and NGC 6369. ",
        "conclusions": "In this paper we have presented the upgraded version of \\shape, with a novel approach to the determination of the three-dimensional structure and kinematics of gaseous nebula. \\shape combines the capabilities of commercial 3-D modeling software with a purpose-built rendering software and graphical control interface for its application to astrophysical nebulae. \\shape produces images and P-V diagrams for highly complex nebulae. The results can be compared directly with observations and help to understand structure, kinematics and orientation of the objects.  The kinematics as observed from long-slit spectra often provides sufficient information about the 3-D structure and topology of an object for the modeling with \\shape to provide a self-consistent 3-D structure. In addition the \\shape models can be applied directly as input density distributions for photo-ionization codes. In \\shape one can define and combine different velocity laws that fit complex structures. Solutions are expected to be most reliable if the object shows evidence for a significant degree of symmetry, then the full 3-D structure and kinematics can be deduced unambiguously. In other cases the object may be devided into regions or subsystens which allows them to be solved separately.  We are currently building a catalogue of of synthetic emission line profiles with \\shape that should be a useful reference to interpret long-slit observations of PNe with diverse morphologies and orientations.  In this paper we showed applications of \\shape to the planetary nebulae Abell~30 and NGC~6369. As a result of our application of \\shape to new observations of the planetary nebula NGC~6369, we propose that its basic structure is that of a ellipsoidal or barrel-shaped main nebula with bipolar protrusions at a large angle to the symmetry-axis of the main nebula. Abell~30 has been modeled combining different velocity patterns to fit the expanding shell and the inner high-velocity knots."
    },
    "0601/astro-ph0601299_arXiv.txt": {
        "abstract": "We investigate BBN in scalar-tensor theories of gravity with arbitrary matter couplings and self-interaction potentials. We first consider the case of a massless dilaton with a quadratic coupling to matter. We perform a full numerical integration of the evolution of the scalar field and compute the resulting light element abundances. We demonstrate in detail the importance of particle mass thresholds on the evolution of the scalar field in a radiation dominated universe. We also consider the simplest extension of this model including a cosmological constant in either the Jordan or Einstein frame. ",
        "introduction": "\\label{sec1}  The concordance model of cosmology calls for the introduction of a cosmological constant or a dark energy sector. Various candidates have been proposed~\\cite{darkrevue}, among which the possibility that gravity is not described by general relativity on large cosmological scales. It is of interest therefore, to test our theory of gravity in a cosmological context. This can be achieved in two complementary ways, either by designing model independent tests (see e.g. Refs.~\\cite{uam,ctes2,bct} for a discussion of the various possible tests) or by considering a class of well motivated theories and use all available data to determine how close to general relativity we must be.  Among all extensions of general relativity, scalar-tensor theories are probably the simplest in the sense that they consider only the introduction of one~\\cite{stgen} (or many~\\cite{def}) scalar field(s) universally coupled to matter. These theories involve two free functions describing the coupling of the scalar field to matter and its self-interaction potential. They respect % local Lorentz invariance and the universality of free fall of laboratory-size bodies. They are motivated by high-energy theories trying to unify gravity with other interactions which generically involve a scalar field in the gravitational sector. In particular, in superstring theories~\\cite{polchy} the supermultiplet of the 10-dimensional graviton contains a scalar field, the dilaton, and other scalar fields, moduli, appear during Kaluza-Klein dimensional reduction of higher dimensional theories to our usual four dimensional spacetime.  In cosmology, two main properties make these theories appealing. First, an attraction mechanism toward general relativity~\\cite{dn1,dpol} was exhibited. This implies that even if the tests of general relativity in the Solar system set strong constraints on these theories, they may differ significantly from general relativity at high redshift. Second, it was shown that the general mechanism of quintessence was conserved~\\cite{jpu99,autres} if the quintessence field was non-minimally coupled and that the attraction mechanism toward general relativity still held with runaway potentials~\\cite{bp,runaway1,runaway2}. These extended quintessence models are the simplest theories in which there is a long range modification of gravity, since the quintessence field is light, and they allow for a very interesting phenomenology~\\cite{msu5}.  Cosmological data give access to various aspects of these models. The cosmic microwave background (CMB) tests the theory in the linear regime~\\cite{ru02,cmb1,cmb2,cmb3,cmb4} while weak lensing opens a complementary window on the non-linear regime~\\cite{carlo,acqua}. Solar system experiments give information on the theory today and big-bang nucleosynthesis (BBN) allows us to constrain the attraction mechanism toward general relativity~\\cite{dn1} at very high redshift.\\\\   BBN is one of the most sensitive available probes of the very early Universe and of physics beyond the standard model. Its success rests on the concordance between the observational determinations of the light element abundances of D, \\he3, \\he4, and \\li7, and their theoretically predicted abundances \\cite{bbn, bbnb}. Furthermore, measurements of the CMB anisotropies by WMAP \\cite{wmap} have led to precision determinations of the baryon density or equivalently the baryon-to-photon ratio, $\\eta$. As $\\eta$ is the sole parameter of the standard model of BBN, it is possible to make very accurate predictions \\cite{cfo3,coc,pdes,cuoco,cyburt} and hence further constrain physics beyond the standard model \\cite{cfos}.  In particular, the \\he4 abundance is often used as a sensitive probe of new physics.  This is due to the fact that nearly all available neutrons at the time of BBN end up in \\he4 and the neutron-to-proton ratio is very sensitive to the competition between the weak interaction rate and the expansion rate.  Of interest to us here is the effect of modifications to gravity which will directly affect the expansion rate of the Universe through a modified Friedmann equation.  The WMAP best fit assuming a varying spectral index is $\\Omega_\\baryon h^2 = 0.0224 \\pm 0.0009$ and is equivalent to $\\eta_{\\rm 10,{\\rm CMB}} = 6.14 \\pm 0.25$, where $\\eta_{10} = 10^{10} \\eta$. Using the WMAP data to fix the baryon density, the light element abundances \\cite{cfo3,coc,pdes,cuoco,cyburt} can be quite accurately predicted. Some BBN results are displayed in Table 1.   \\begin{table*} \\caption{BBN results for the light element abundances assuming the WMAP-inferred baryon density.} \\begin{center} \\begin{tabular}{l|c|c|c|c} \\hline\\hline Source& $Y_p$ & ${\\rm D}/{\\rm H}$&$^3{\\rm He}/{\\rm H}$ & $^{7}{\\rm Li}/{\\rm H}$\\\\ &&$\\times10^{-5}$&$\\times10^{-5}$&$\\times10^{-10}$\\\\ \\hline\\hline Coc {\\it et al.} (2004)& 0.2479$\\pm0.0004$&$2.60\\pm0.17$&$1.04\\pm0.04$&$4.15\\pm0.46$\\\\ \\hline Cyburt {\\it et al.} (2003) & $0.2484^{+0.0004}_{-0.0005}$ & $2.74^{+0.26}_{-0.16}$ & $0.93^{+0.1}_{-0.67}$ & $3.76^{+1.03}_{-0.38}$\\\\ \\hline\\hline \\end{tabular} \\end{center} \\label{t:yields} \\end{table*}  The effect of scalar-tensor theories of gravity on the production of light elements has been investigated extensively (see e.g. Ref.~\\cite{ctes} for a review). As a first step, it is useful to consider only the speed up factor, $\\xi=H/H_{\\rm GR}$, that arises from the modification of the value of the gravitational constant during BBN~\\cite{speedup1,cfos}. Other approaches considered the full dynamics of the problem but restricted themselves to the particular class of Jordan-Fierz-Brans-Dicke theory~\\cite{bbnJFBD}, of a massless dilaton with a quadratic coupling~\\cite{dp,bbn_quad} or to a general massless dilaton~\\cite{bbnst1}. It should be noted that a combined analysis of BBN and CMB data was investigated in Ref.~\\cite{copi} and Ref.~\\cite{bbnst2}. The former considered $G$ constant during BBN while the latter focused on a non-minimally quadratic coupling and a runaway potential. We stress that the dynamics of the field can modify CMB results so that one needs to be careful while inferring $\\Omega_\\baryon$ from WMAP.  The goal of this article is to implement scalar-tensor theories in an up-to-date BBN code. This will complement our existing set of tools which allows us to confront scalar-tensor theories with observations of  type Ia supernovae  and CMB anisotropies~\\cite{ru02} as well as weak lensing~\\cite{carlo}. In particular, the predictions to be compared with observations can be computed in the same framework for any self-interaction potential and matter-coupling function.  We first recall, in \\S~\\ref{sec_ST_theory}, the equations describing the theory to be implemented in our BBN code and we also discuss local constraints. As a check of our code, we consider, in \\S~ \\ref{sec_modelDP}, the case of a massless dilaton with quadratic coupling~\\cite{dp}. In particular, we perform a full numerical integration up to the present that can be compared with the analytical results of Ref.~\\cite{dp}. We update the constraints on this model by taking into account the latest BBN data discussed above. We reaffirm that only helium-4 is sensitive to the modification of gravity considered here. In \\S~\\ref{sec_model2}, we will consider the simplest extension of this model by introducing a cosmological constant. Such a constant can be introduced as a constant potential either in the Einstein frame, hence keeping the dilaton massless, or in the Jordan frame, hence generalizing the constant energy density component. Both cases are considered and we conclude in Section~\\ref{sec_concl}. Applications to various cases of cosmological interest will be presented in a follow-up article. ",
        "conclusions": "This article described the implementation of general scalar-tensor theories applicable to a BBN code. The formalism allows one to choose any self-interaction potential as well as any coupling to matter. As such it can be applied to models which account for the present day acceleration such as (extended) quintessence models.  The ability to use BBN as a constraint completes our set of tools, which include CMB anisotropies and SNIa~\\cite{ru02} and weak lensing~\\cite{carlo}, to study the cosmological imprints of this set of well-motivated theories of gravity. All observables are computed using the same formalism for compatibility. They can be used conjointly to set constraints on these theories and on deviations from general relativity during the entire evolution of our universe. We emphasize their complementarity since BBN depends only on the background evolution and mainly tests the attraction mechanism toward GR, CMB is mainly sensitive to the evolution of the perturbations in the linear regime while weak lensing probes the non-linear regime.  In this article, we have focused on the case of a quadratic coupling in order to check our code. In the case where $V=0$ our results are compatible with previous analysis~\\cite{dp}. Note however that they do not rely on any specific form of the analytic solution. Also, our evaluation of the Fermi integrals that are necessary to estimate the kick during electron-positron annihilation do not rely on an approximation but rather on a full numerical integration. We have used a complete BBN code with up to date nuclear reaction rates. Current data on the light element abundances have been used to set constraints and we have also investigated the effect of a cosmological constant on these constraints. For this particular model, CMB anisotropies are not affected and we are allowed to infer $\\Omega_\\baryon$ from standard CMB analyses. We emphasize that in general this has to be checked case by case.  Since our approach is fully numerical, it can be applied to any scenario and in particular to extended quintessence scenarios such as models with runaway fields. In these models, during the radiation era, the field evolves to reach a scaling solution. Before this, there may be a kinetic phase. According to when this kinetic phase ends, various effects on BBN can be expected. In particular $\\varphi_{*\\xin}$=const. may not be a good approximation. It was also proposed that the coupling to dark matter may be different to the coupling to standard matter. This hypothesis relaxes the Solar system bound and allows higher values of $\\alpha_{\\rm cdm}$. All these questions, and others, will be adressed in following works.  \\noindent{\\bf Acknowledgements}: We would like to thank T. Damour, G. Esposito-Far\\`ese and C. Schimd for discussions and comments. The work of K.A.O. was supported in part by DOE grant DE--FG02--94ER--40823. The work is also supported by the project ``INSU--CNRS/USA''."
    },
    "0601/astro-ph0601250_arXiv.txt": {
        "abstract": "The galaxy pair AM\\,1353-272 nicknamed ``The Dentist's Chair'' shows two $\\sim$30\\,kpc long tidal tails. Previous observations using multi-slit masks showed that they host up to seven tidal dwarf galaxies. The kinematics of these tidal dwarfs appeared to be decoupled from the surrounding tidal material. New observations of the tip of the southern tidal tail with the VIMOS integral field unit confirm the results for two of these genuine tidal dwarfs but raise doubts whether the velocity gradient attributed to the outermost tidal dwarf candidate is real.  We also discuss possible effects to explain the observational difference of the strongest velocity gradient seen in the slit data which is undetected in the new integral field data, but arrive at no firm conclusion. Additionally, low-resolution data covering most of the two interacting partners show that the strongest line emitting regions of this system are the central parts. ",
        "introduction": "\\label{PmW:Sect:Intro} Following old ideas about the creation of dwarf galaxies during interaction of giant galaxies and detailed investigations of several nearby examples of these Tidal Dwarf Galaxies \\citep[TDGs,][]{DBS+00,DBW+97,DM98,HGvG+94}, we carried out a first small survey of interacting galaxies \\citep{WDF+00} with the aim of better understanding the star-formation history of TDGs and constraining the number of TDGs that are built per interaction \\citep{WDF03}.  During this survey, we studied a system cataloged as AM\\,1353-272, which we called ``The Dentist's Chair'' for its peculiar shape. Fig.~\\ref{PmW:Fig:System} shows the two components of the system, `A', a galaxy with $\\sim$30\\,kpc long tidal tails, and `B', a disturbed disk galaxy, have a distance of $D \\approx 160$\\,Mpc ($H_0 = 75$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$). Within the tails several obvious clumps with blue optical colors are visible. Using optical and near-infrared imaging, evolutionary models, and optical spectroscopy, \\emph{seven} of these clumps were classified as TDG candidates in formation \\citep[][, marked `a' to `d' and `k' to `m']{WFDF02}. The largest velocity gradient with an amplitude of $>$300\\,km\\,s$^{-1}$ appeared in TDG candidate `a', at the very end of the southern tidal tail. This raised the question how an object with relatively low luminosity could exhibit such fast ``rotation''. However, as these observations were done using the multi-slit technique and hence are spatially restricted due to the narrow slit, subsequent observations were planned using an integral field unit (IFU) to cover more of the tidal tails and view the velocity structure of the TDGs in two dimensions.  \\begin{figure} \\centering \\includegraphics[width=9cm]{Weilbacher_3Dspec05_chart_grey.eps} \\caption{The interacting system AM\\,1353-272. The two interacting galaxies (capital letters) and the relevant knots in the tidal tails are marked (lower case letters). Overlayed are the original FORS2 slits and the two VIMOS pointings in low and high resolution mode.}\\label{PmW:Fig:System} \\end{figure} ",
        "conclusions": "\\label{PmW:Sect:Sum} We presented a few tentative results for the interacting system called ``The Dentist's Chair'' from new observations with the VIMOS integral field mode: line emission seems to be                          % concentrated within the centers of the interacting partners while the tidal tails themselves are not detected H$\\beta$ narrowband slices. Three knots, previously identified as TDG candidates near the end of the southern tidal tail, are detected in [O{\\sc iii}]5007 emission. For two of them the velocity profiles were confirmed. However, the strongest velocity gradient in the outermost TDG candidate (knot `a') as measured on FORS2 data is not confirmed by the VIMOS datacube. The reason for this discrepancy is unknown.  To solve this mystery and find more clues to the origin of the velocity fields seen in this interacting system, further, deeper IFU observations, taken with appropriate dither offsets to facilitate more accurate sky subtraction, are required. With other instruments like e.\\,g.\\ the GMOS IFU would be possible to cover both the Ca-triplet and H$\\alpha$ in the same exposures and directly compare the stellar velocity field with ionized gas dynamics. As good S/N is required to detect the absorption lines, this can only be done in the brighter northern tidal tail.  \\begin{acknowledgement} PMW received financial support through the D3Dnet project from the German Verbundforschung of BMBF (grant 05AV5BAA). We are grateful to Ana Monreal-Ibero and Lise Christensen for practical hints on IFU data handling. The data we discuss was taken in service mode at Paranal (ESO Program 074.B-0629). \\end{acknowledgement}"
    },
    "0601/nucl-th0601002_arXiv.txt": {
        "abstract": "The sensitivity of the differential cross section of the interaction between neutrino-electron with dense matter to the possibly nonzero neutrino electromagnetic properties has been investigated. Here, the relativistic mean field model inspired by effective field theory has been used to describe non strange dense matter, both with and without the neutrino trapping. We have found that the cross section becomes more sensitive to the constituent distribution of the matter, once electromagnetic properties of the neutrino are taken into account. The effects of electromagnetic properties of neutrino on the cross section become more significant for the neutrino magnetic moment $\\mu_{\\nu}$ $>$ $10^{-10}\\mu_B$ and for the neutrino charge radius $R$   $>$ $10^{-5} {\\rm MeV}^{-1}$. ",
        "introduction": "\\label{sec_intro} The demand for precise  information on the neutrino transport in the investigations of astrophysical phenomena, such as supernovae explosion and the structure of protoneutron stars, has stimulated several studies on neutrino interactions in matter at high densities~\\cite{reddy1,niembro01,parada04,Horo2,mornas1,mornas2,reddy2,horowitz91,yama,caiwan,margue,chandra,cowel,caroline05}. Even recently, a realistic neutrino opacity in neutron rich matter for the supernovae simulation which takes into account the correlations and weak magnetism of nucleons in the finite temperature calculations has been also considered \\cite{Horo2}.  In the standard model, massless neutrinos have zero magnetic moment and electronic charge. However, there are evidences that the laboratory bounds on the neutrino-electron magnetic moment ($\\mu_{\\nu}$) is smaller than $1.0 \\times 10^{-10} \\mu_B$ at the 90\\% confidence level~\\cite{munu}, where $\\mu_B$ is the Bohr magneton. Stronger bound  of $\\mu_{\\nu} \\leq 3.0 \\times 10^{-12} \\mu_B$ also exists from astrophysical consideration, particularly from the study of the red giant population in the globular cluster~\\cite{Raffelt99}. On the other hand the charge radius of $\\nu_e$ has been bounded by the LAMF experiment~\\cite{vilain} to be $R^2=(0.9 \\pm 2.7)\\times 10^{-32}~{\\rm cm}^{2}=(22.5 \\pm 67.5)\\times 10^{-12}~{\\rm MeV}^{-2}$, while the plasmon decay in the globular cluster star predicts the limit of $e_{\\nu} \\leq 2 \\times 10^{-14}\\,e$~\\cite{Raffelt99}, where $e$ is electron charge. Recent discussions on the effects of the neutrino electromagnetic properties in astrophysics can be seen in, e.g., Ref.~\\cite{Raffelt99}, while for the case of the solar neutrino problem one can consult, e.g., Ref.~\\cite{friedland}.  So far, there has been no calculation of the interaction of neutrinos with dense matter  which considers neutrino electromagnetic form factors. To this end we extend our previous report of  the interaction of neutrinos with electrons gas~\\cite{caroline05} to the interaction of neutrinos with the non strange dense stellar matter which takes into account the effects of the  trapped neutrinos in matter, but still use the zero temperature approximation. The validity of this approximation is fulfilled by the fact that temperature effects on the equation of state (EOS) of supernovae matter and maximum masses of protoneutron stars are smaller than those without neutrino trapping~\\cite{hua03}. We note that considerable efforts have been devoted to study the effective electromagnetic coupling of neutrinos in the thermal background of particles~\\cite{nives1,nives2,sawyer,samina}. However, the particles under investigation so far are electron and nucleon gasses. Clearly, the effect of nucleon correlations has not yet been studied. References~\\cite{nives1,samina} have shown that the densities and temperature yield enhancements to the form factor with standard charge radius, but they do not give a significant effect on the electric and magnetic dipole form factors. Moreover, they can only give a significant contribution if $T/m$ and $\\mu/m$ are large, where $T$, $\\mu$ and $m$ are the temperature, chemical potential and mass of the corresponding particles in the thermal bath, respectively~\\cite{samina}. Therefore, the expectation that the temperature and hadron correlations can strongly reduce the electromagnetic properties of neutrinos seems to be unlikely. For simplicity, the RPA correlations are still neglected. The importance of the neutrino trapping in the supernova dynamical evolution has been, for example, pointed out in Ref.~\\cite{chiapparini}. We choose this kind of matter, because we suspect that the effects of the neutrino electromagnetic properties would be more pronounced in the protons- and electrons rich matter.  Different from Ref.~\\cite{chiapparini}, which used the standard relativistic mean field (RMF) model to calculate the EOS, in this work we use the RMF model inspired by the effective field theory (E-RMF model)~\\cite{Furnstahl96}.  The construction of this paper is as follows. Section~\\ref{sec_nucmod} consists of a brief review of the matter models used in this work. The analytical results of the neutrino electromagnetic form factors effects on the cross section are given in section~\\ref{sec_EFFN}. In section~\\ref{sec_rslt}, numerical results and their discussions are presented. Finally, the conclusion is given in section~\\ref{sec_sum}. ",
        "conclusions": "\\label{sec_sum} In conclusion, we have studied the sensitivity of the neutrino cross section to the neutrino electromagnetic properties. In our calculations, we use the G2* parameter set of the E-RMF model to describe matter. This parameter set predicts a soft EOS at high density and has $\\rho_B= 2.5\\rho_0$ that coincide with the direct URCA process threshold. The calculation has been performed for the cases in which neutrinos are trapped and absent in matter. It is found that in the non strange stellar matter, the electromagnetic form factor has an important role in the neutrino-electron matter cross section if  $\\mu_{\\nu}>10^{-10} \\mu_B$ and $R>10^{-5}~{\\rm MeV}^{-1}$. It is  also found that the effects of the neutrino electromagnetic form factors on the cross sections are more pronounced at higher densities.  Matters with trapped neutrino, like supernovae ones, are more sensitive  to the presence of neutrino electromagnetic properties. This is due to the fact that matters with trapped neutrino have larger fractions of protons and electrons than those without neutrinos.Although we have found that the role of neutrino electromagnetic properties in the neutrino-electron matter interaction is not too crucial, this would not be the case if we considered  the bound of the  neutrino-muon dipole moment given by Refs.~\\cite{allen,kraka}.  With increasing the density of the protoneutron star, strangeness can be liberated and face up in the the filling of the hyperon Fermi seas and/or in the creation of kaon condensates. Occurrence of these exotics in protoneutron star interiors will enhance the neutrino scattering rate and it may have also interesting observational consequences, like softening the EOS, possibility of changing nucleon isospin composition in the star matter evolution, as well as enhancing neutrino emission processes in the neutron star matter evolution, as extensively discussed in Refs.~\\cite{kolo,knorren,baldo1,vidana1,bunta,zuo,banik,hua03,hua2}.  Further, depending on the temperature and the model used, baryon correlations can also reduce the scattering rate~\\cite{Horo2,mornas2,horowitz91,caiwan,cowel}. Therefore, an extension of this calculation by including the strange matter, RPA correlation and a more general condition, i.e., finite temperature might also be interesting and quite relevant to consider in the future.  \\appendix"
    },
    "0601/astro-ph0601700_arXiv.txt": {
        "abstract": "We present results from {\\it Chandra} observations of SDSS~J1004+4112, a strongly lensed quasar system with a maximum image separation of $15\\arcsec$. All four bright images of the quasar, as well as resolved X-ray emission originating from the lensing cluster, are clearly detected. The emission from the lensing cluster extends out to approximately $1\\farcm5$.  We measure the bolometric X-ray luminosity and temperature of the lensing cluster to be $4.7 \\times 10^{44}~{\\rm erg\\,s^{-1}}$ and 6.4~keV, consistent with the luminosity-temperature relation for distant clusters.  The mass estimated from the X-ray observation shows excellent agreement with the mass derived from gravitational lensing.  The X-ray flux ratios of the quasar images differ markedly from the optical flux ratios, and the combined X-ray spectrum of the images possesses an unusually strong Fe~K$\\alpha$ emission line, both of which are indicative of microlensing. ",
        "introduction": "The quadruply lensed quasar SDSS~J1004+4112, first identified as a gravitational lens candidate from optical imaging and spectroscopy from the Sloan Digital Sky Survey \\citep[SDSS;][]{york00}, has an exceptionally large image separation of $\\sim15\\arcsec$ \\citep{inada03,oguri04}. The $z=1.734$ quasar is multiply imaged by a cluster of galaxies (rather than a galaxy) at $z=0.68$. This is the only known quasar strongly lensed by a central part of a massive cluster. The quasar appears to be radio-quiet since it is not detected in radio surveys such as the FIRST survey \\citep{becker95}. Subsequent observations revealed additional intriguing aspects of the system. \\citet{richards04} discovered time variability in the blue wings of several optical broad emission lines, which can be plausiblly interpreted in terms of microlesning.  {\\it Hubble Space Telescope} images detected the fifth image \\citep{inada05} and multiply imaged background galaxies \\citep{sharon05}. SDSS~J1004+4112 appears to be an ideal laboratory for exploration of the structure of a large cluster of galaxies and a quasar.  The uniqueness of SDSS~J1004+4112 affords advantages that strongly argue for multiwavelength observations, of which X-ray observations constitute an essential part. Measurement of the X-ray properties of the lensing cluster offers an opportunity to improve lens models of SDSS~J1004+4112. Moreover, the lensing cluster is the first example of a {\\it strong-lens selected} cluster; it is interesting to determine if the cluster follows the empirical scaling relations between luminosities, temperatures, and masses. Second, we can investigate the magnification ratios of the quasar images at X-ray wavelengths, which may aid the interpretation of the microlensing event that has been observed at optical wavelengths \\citep{richards04}. Indeed, X-ray flux ratios of lensed quasars are frequently different from optical flux ratios; this phenomenon has been attributed to microlensing \\citep{chartas04,blackburne06}.  In this paper, we report our analysis of an observation of SDSS~ J1004+4112 with the {\\it Chandra X-Ray Observatory}.  A key feature of {\\it Chandra} images is the high angular resolution (on the order of $1\\arcsec$); this resolution is required to allow reliable separation of the X-ray flux of the lensing cluster from that of the quadruply imaged quasar.  In this paper we will adopt a cosmological model with the matter density $\\Omega_M=0.27$, the cosmological constant $\\Omega_\\Lambda=0.73$, and the Hubble constant $H_0=70~{\\rm km\\,s^{-1}Mpc^{-1}}$ \\citep{dns03}. At the redshift of the cluster ($z=0.68$), $1\\arcsec$ corresponds to 7.16~kpc.  Unless otherwise specified, quoted errors indicate the 90\\% confidence range. ",
        "conclusions": "We have presented results from {\\it Chandra} observations of SDSS J1004+4112, a large-separation gravitational lens system created by a cluster of galaxies.  We have detected X-ray emission from the lensing cluster as well as four lensed quasar images.  From the cluster X-ray emission, we have constrained the bolometric luminosity and the temperature to be $L_{\\rm X}=4.7\\times 10^{44}~{\\rm erg\\,s^{-1}}$ and $kT=6.4^{+2.3}_{-1.4}$~keV, consistent with the luminosity-temperature relation of distant clusters. We have reconstructed the mass profile of the lensing cluster assuming isothermality and hydrostatic equilibrium, and found that the mass within 100~kpc excellently agrees with that expected from strong lensing.  X-ray emission from the lensed quasar images displays two anomalies: the presence of a strong Fe K$\\alpha$ line and significant differences in the flux ratios from those found in the optical band. Both of these features are suggestive of microlensing. The idea is supported by the fact that image A, which appears most anomalous, is a saddle-point image."
    },
    "0601/astro-ph0601470_arXiv.txt": {
        "abstract": "We calculate theoretical isochrones in a consistent way for five filter pairs near the $J$ and $K$ band atmospheric windows ($J$--$K$, $J$--$K'$, $J$--$K_s$, F110W--F205W, and F110W--F222M) using the Padova stellar evolutionary models of Girardi et al.  We present magnitude transformations between various $K$-band filters as a function of color. Isochrones with extinction of up to 6 mag in the $K$ band are also presented.  As found for the filter pairs composed of $H$ \\& $K$ band filters, we find that the reddened isochrones of different filter pairs behave as if they follow different extinction laws, and that the extinction curves of {\\it Hubble Space Telescope} NICMOS filter pairs in the color-magnitude diagram are considerably nonlinear.  Because of these problems, extinction values estimated with NICMOS filters can be in error by up to 1.3 mag. Our calculation suggests that the extinction law implied by the observations of Rieke et al for wavelengths between the $J$ and $K$ bands is better described by a power-law function with an exponent of 1.66 instead of 1.59, which is commonly used with an assumption that the transmission functions of $J$ and $K$ filters are Dirac delta functions. ",
        "introduction": "\\label{sec:introduction}  Kim et al. (2005, hereafter Paper I) have calculated theoretical isochrones with extinction for some $H$ and $K$ band filters using the Padova stellar evolutionary models by Girardi et al. (2002).  In Paper I, we found that the reddened isochrones of different filter pairs in $H$ and $K$ bands behave as if they follow different extinction laws, and that care is needed when applying an extinction law obtained with one filter pair to other, similar filter pairs.  For example, if the extinction law for the Johnson-Glass $H$ and $K$ filters obtained by Rieke, Rieke, \\& Paul (1989) is directly applied to the photometry from the {\\it Hubble Space Telescope} ({\\it HST}) NICMOS filters (F160W, F205W, and F222M), estimated extinction values can be in error by up to 0.3 mag for true extinction at $K$ of 6 mag or less.  To reduce this error, Paper I introduced an ``effective extinction slope'' for each filter pair and isochrone model.  It was also found that the extinction behavior of isochrones in the color-magnitude diagram (CMD) for filter pair F160W--F222M is highly nonlinear (i.e., the amount of extinction is not proportional to color excess) because of a significant width difference in the two filters.  These problems are certainly not limited to the isochrones for filter pairs in the $H$ and $K$ bands.  This problem will apply to any situation in which one applies an extinction law deduced from one filter pair to other similar filter pairs.  Furthermore, the nonlinear behavior of the extinction vector in the CMD will be problematic for filter pairs with significant difference in width.  In the present paper, we extend the calculations performed in Paper I to the isochrones for filter pairs in the $J$ and $K$ bands.  The filters considered here are the four ground-based filters $J$, $K$ (Johnson et al. 1966), $K'$ (Wainscoat \\& Cowie 1992), and $K_s$ ($K$-short; developed by M. Skrutskie; see the appendix of Persson et al. 1998), and the three NICMOS filters F110W, F205W, and F222M (transmission functions of these filters are shown in Figure~\\ref{fig:filter}).  Out of these seven filters, we consider five filter pairs: $J$--$K$, $J$--$K'$, $J$--$K_s$, F110W--F205W, and F110W--F222M.  We adopt a Vega-based photometric system (VEGAMAG system), which uses Vega ($\\alpha$ Lyr) as the calibrating star.  For photometric zero points of NICMOS filters, we adopt $\\langle f_\\nu^{\\rm Vega} \\rangle$ values from the NICMOS Data Handbook (ver. 5.0): 1775~Jy for F110W, 703.6~Jy for F205W, and 610.4~Jy for F222M.  For the spectra of synthetic stellar atmospheres, we adopt Kurucz ATLAS9 no-overshoot models\\footnote{See NOVER files at http://kurucz.harvard.edu/grids.html.} (Kurucz 1993) calculated by Castelli et al. (1997).  The metallicities of these models cover the values of [M/H] = $-2.5$ to $+0.5$.  A microturbulent velocity $\\xi=2\\,{\\rm km \\, s^{-1}}$ and a mixing length parameter $\\alpha =1.25$ are adopted in the present study.  For the temporal evolution of effective temperature and luminosity as functions of stellar mass (i.e., stellar evolutionary tracks), we adopt the ``basic set\" of the Padova models\\footnote{See http://pleiadi.pd.astro.it.} (Girardi et al. 2002).  We consider isochrones with a metallicity $Z$ = 0.0001, 0.001, 0.019, and 0.03.  The stellar spectral library and the evolutionary tracks we adopted assume a solar chemical ratios.  For more details on the magnitude system, stellar spectral library, and evolutionary tracks that we adopt here, readers are referred to Paper I.  Throughout this paper, we generically refer to the atmospheric wavebands centered near 1.25, 1.65, and 2.2~\\micron, as the $J$, $H$, and $K$ bands, whereas we refer to the Johnson-Glass filters (Johnson et al. 1966; Glass 1974) as the $J$, $H$, and $K$ filters. ",
        "conclusions": "\\label{sec:summary}  We have calculated in a consistent way five near-infrared theoretical isochrones for filter pairs composed of $J$ and $K$ filters: $J$--$K$, $J$--$K'$, $J$--$K_s$, F110W--F205W, and F110W--F222M.  We presented isochrones for a $Z$ of 0.0001--0.03 and an age of $10^7$--$10^{10}$~yr. Even in the same Vega magnitude system, near-infrared colors of the same isochrone can be different by up to $\\sim 0.4$ mag at the bright end of the isochrone for different filter pairs. The difference in intrinsic colors for a red giant for atmospheric filters and the {\\it HST} NICMOS filters is generally 0.2--0.4 mag. We have provided magnitude transformations between $K$-band filters as a function of color from $J$ and $K$ band filters. We also presented isochrones with $A_K$ of up to 6 mag. We found that care is needed when comparing extinction values that are estimated using different filter pairs, in particular when comparing those of atmospheric and NICMOS filter pairs: extinction values inferred using NICMOS filters can be in error by up to 1.3 mag.  To alleviate this problem, we introduced an ``effective extinction slope'' for each filter pair and isochrone model, which describes the extinction-dependent behavior of isochrones in the observed CMD.  We also provided a procedure to accurately estimate the extinction value for NICMOS filter pairs, whose extinction curves in the CMD are highly nonlinear."
    },
    "0601/astro-ph0601193_arXiv.txt": {
        "abstract": "We present 2D MHD simulations of the radiatively driven outflow from a rotating hot star with a dipole magnetic field aligned with the star's rotation axis. We focus primarily on a model with moderately rapid rotation (half the critical value), and also a large magnetic confinement parameter, $\\estar \\equiv \\Bstar^2 \\Rstar^{2}/ \\Mdot \\Vinf = 600$. The magnetic field channels and torques the wind outflow into an equatorial, rigidly rotating disk extending from near the Kepler corotation radius outwards.  Even with fine-tuning at lower magnetic confinement, none of the MHD models produce a stable Keplerian disk. Instead, material below the Kepler radius falls back on to the stellar surface, while the strong centrifugal force on material beyond the corotation escape radius stretches the magnetic loops outwards, leading to episodic breakout of mass when the field reconnects.  The associated dissipation of magnetic energy heats material to temperatures of nearly $10^{8}\\,{\\rm K}$, high enough to emit hard (several keV) X-rays.  Such \\emph{centrifugal mass ejection} represents a novel mechanism for driving magnetic reconnection, and seems a very promising basis for modeling X-ray flares recently observed in rotating magnetic Bp stars like $\\sigma$~Ori~E. ",
        "introduction": "\\label{sec:intro}  Magnetic fields can greatly influence stellar winds by guiding the outflowing plasma along field lines \\citep[e.g.,][]{MacG1988,ShoBro1990}.  The extent of this influence depends largely on the magnetic field strength and the wind mass loss rate.  Two-dimensional (2D) magnetohydrodynamical (MHD) simulations of axisymmetric hot-star winds with no rotation show that the competition between the wind and the field can be quite well characterized by a single wind magnetic confinement parameter, \\estar, as defined by \\citet{udDOwo2002}, and described further in \\S\\ref{sec:alfven} below.  For significant confinement (i.e., $\\estar \\ga 10$), material driven from opposite footpoints of closed magnetic loops is forced to collide near the loop tops, leading to a \\emph{Magnetically Confined Wind Shock} (MCWS) scenario for production of intermediate-hardness X-rays \\citep{BabMon1997a,BabMon1997b}.  For the magnetic hot star \\tOriC, which is inferred to have both a moderately strong wind (with mass loss rate $\\approx 10^{-7}$\\,\\Msun/\\years) and a large-scale, dipole field of strength $\\Bstar \\approx 1100\\,\\gauss$ \\citep{Don2002}, recent MHD simulations based on this MCWS paradigm provide quite good agreement with the X-ray spectrum observed by \\emph{Chandra} \\citep{Gag2005}.  Indeed, because \\tOriC\\ is a relatively slow rotator (with period $\\sim$~15~days), even the observed rotational modulation can be well accounted for by simply varying the observational perspective within a {\\em non-rotating} dynamical simulation model.  The purpose of this letter is to report on initial MHD simulation results for more rapidly rotating hot stars, in which the rotation can have a direct dynamical effect on the overall evolution of the wind outflow and its confinement.  The results have particular relevance in the context of two proposed models for rotating, magnetic hot stars, known generally as the \\emph{Magnetically Torqued Disk} (MTD) model \\citep{Cas2002} and the \\emph{Rigidly Rotating Magnetosphere} (RRM) model \\citep{TowOwo2005}.  The MTD has been proposed as a scenario for producing \\emph{Keplerian} decretion disks, such as inferred from the Balmer line emission of classical Be stars.  The RRM leads instead to \\emph{rigid-body} disks or clouds that provide a promising way to explain the rotationally modulated Balmer emission from strongly magnetic Ap/Bp stars like the B2pV star \\sOriE\\ \\citep{Tow2005}.  After first casting the relative importance of magnetic confinement vs.\\ rotation in terms of characteristic Alfv{\\`e}n vs.\\ Kepler corotation radii (\\S\\ref{sec:alfven}), we show in \\S\\ref{sec:keplerian} that even for parameters fine-tuned to optimize the chances of a producing a Keplerian disk, our 2D MHD simulations for rotationally aligned, dipole fields instead produce a combination of inflow and outflow for material below and above the Kepler corotation radius, and never lead to the kind of extended Keplerian disk envisioned in the MTD scenario.  On the other hand, we show in \\S\\ref{sec:breakout} that a model with strong magnetic confinement ($\\estar = 600$) does lead to a \\emph{rigid-body} disk, much as predicted in the RRM model.  However, the outer regions of this disk are characterized by episodic breakout of material, driven by the strong net outward centrifugal force associated with the spin-up of material into corotation with the underlying star.  Moreover, the energy release associated with magnetic reconnection in these breakout events heats material to extremely high temperatures, providing a promising mechanism for explaining the very hard X-ray flare emission recently observed from magnetic Bp stars such as \\sOriE.  We conclude in \\S\\ref{sec:summary} with a brief summary of our findings, noting also the need for further work on the detailed physics of this somewhat-novel, centrifugally driven reconnection mechanism. ",
        "conclusions": "\\label{sec:summary}  The rigid-body corotation makes the equatorial disks in our simulations quite distinct from the Keplerian disk envisioned in the MTD scenario promoted by \\citet{Cas2002}. We find no evidence of an extended region of Keplerian rotation.  Instead the results correspond more closely to those derived in the RRM analysis of \\citet{TowOwo2005}, although the confinement parameter we consider is near the lower end of the range of validity for the basic rigid-field approach.  Moreover, our numerical simulations further suggest a novel mechanism for magnetic reconnection, driven by episodic centrifugal breakout events.  The strong reconnection heating associated with such breakout represents a promising mechanism for explaining X-ray flares observed in magnetic Bp stars such as \\sOriE.  In future work, we plan to use more-extensive numerical tools to investigate in greater detail the reconnection processes within this centrifugally driven breakout scenario.  We also plan to explore hybrid methods that could allow extensions of MHD simulations to larger confinement parameters, including 3D models of tilted-field cases more appropriate to \\sOriE\\ and other magnetic hot-stars."
    },
    "0601/astro-ph0601646_arXiv.txt": {
        "abstract": "{We study the process whereby quantum cosmological perturbations become classical within inflationary cosmology. By setting up a master-equation formulation we show how quantum coherence for super-Hubble modes can be destroyed by its coupling to the environment provided by sub-Hubble modes. We identify what features the sub-Hubble environment must have in order to decohere the longer wavelengths, and identify how the onset of decoherence (and how long it takes) depends on the properties of the sub-Hubble physics which forms the environment. Our results show that the decoherence process is largely insensitive to the details of the coupling between the sub- and super-Hubble scales. They also show how locality implies , quite generally, that the decohered density matrix at late times is diagonal in the field representation (as is implicitly assumed by extant calculations of inflationary density perturbations). Our calculations also imply that decoherence can arise even for couplings which are as weak as gravitational in strength. }  \\begin{document}    \\makeatletter \\@addtoreset{equation}{section} \\makeatother \\renewcommand{\\theequation}{\\thesection.\\arabic{equation}}  \\setcounter{page}{1} \\pagestyle{plain} \\renewcommand{\\thefootnote}{\\arabic{footnote}} \\setcounter{footnote}{0} ",
        "introduction": "Recent precision observations \\cite{WMAP,boomerang,dasi} of temperature fluctuations within the Cosmic Microwave Background (CMB) agree well \\cite{WMAPInflation} with the predictions made for them by inflationary models, according to which the Universe once underwent an accelerated expansion during its remote past \\cite{Inflation}. According to the inflationary picture, the observed CMB temperature variations are seeded by tiny primordial density perturbations in the much earlier universe, perturbations which began as quantum fluctuations during the earlier inflationary epoch.  More precisely, within the inflationary picture the primordial density fluctuations derive from the quantum fluctuations, $\\langle \\phi(x) \\phi(y) \\rangle$, of the quantum inflaton field, $\\phi$, whose dynamics drives the signature accelerated expansion. The success of the inflationary picture requires these {\\it quantum} fluctuations to be converted into classical {\\it spatial} variations in the energy density whose amplitude varies stochastically as one moves from one Hubble patch to another within the Universe during the much later recombination epoch.  At present, the classicalization of primordial perturbations is understood in terms of the properties of the quantum state into which the fluctuation fields evolve (see, for example \\cite{Guth:1985ya,Polarski:1995jg}). It is known that both scalar and tensor density fluctuations evolve during inflation into a particular kind of quantum state -- one very similar to the squeezed state of quantum optics \\cite{Grishchuk:1990bj,Albrecht:1992kf}. In particular, quantum expectation values of products of fields in a highly squeezed state turn out to be identical to stochastic averages calculated from a stochastic distribution of classical field configurations, up to corrections which vanish in the limit of infinite squeezing \\cite{Polarski:1995jg,Kiefer:1998qe,Lesgourgues:1996jc}. Quantum autocorrelation functions like $\\langle \\phi(x) \\phi(y) \\rangle$ become indistinguishable from autocorrelations for $c$-number fields $\\varphi(x)$ computed within a classical stochastic process, when computed in the limit of large squeezing.  Our purpose in this note is to describe what these analyses do {\\it not} address: the decoherence process through which the initial quantum state evolves into the corresponding stochastic ensemble of field configurations. That is, we wish to understand how the density matrix, $\\rho[\\varphi(x),\\varphi'(y)] = \\langle \\varphi(x)| \\rho | \\varphi'(y) \\rangle$, for each of the various quantum fields, $\\phi(x)$, evolves from an initially pure state, \\be \\label{PureState} \\rho[\\varphi(x),\\varphi'(y)] = \\Psi[\\varphi(x)] \\Psi^*[\\varphi'(y)] \\,, \\ee describing the initial conditions (where $\\Psi[\\varphi(x)] = \\langle \\varphi(x)|\\psi \\rangle$), to a mixed state, \\be \\label{DecoheredState} \\rho[\\varphi(x), \\varphi'(y)] = P[\\varphi(x)] \\, \\delta[\\varphi(x) - \\varphi'(y)] \\,, \\ee describing the later classical stochastic probability distribution for the field amplitudes. Such a transition is implicit in the standard predictions of inflationary implications for the CMB, and our goal is to understand to what extent these predictions might depend on the details of how this decoherence is achieved.  Decoherence is a much-studied process outside of cosmology, although not all of the conceptual issues have been resolved. In particular, the transition from pure to mixed state we seek typically arises whenever the degrees of freedom of interest interact with an `environment' consisting of other degrees of freedom whose properties are not measured. The measured degrees of freedom can make a transition from a pure to a mixed state once a partial trace is taken over this environmental sector.  We are most interested in those modes which are responsible for the correlations which are observed in the CMB and in structure formation -- modes which have only come within the Hubble scale relatively recently. We take the relevant decohering environment to consist of those modes having much shorter wavelength, which entered the Hubble scale much earlier and whose correlations are not directly visible in the sky. By borrowing techniques from atomic and condensed-matter physics we set up a master equation for the density matrix of the observed modes which allows us to robustly identify how these modes decohere provided only that ($i$) the environment is large enough to not be overly perturbed by the modes whose decoherence we follow; ($ii$) the interactions with the environment are weak; and ($ii$) the characteristic correlation time of fluctuations in the environment is much smaller than the timescales of interest to the decoherence process.  We use this formalism to identify how the decoherence depends on the properties of the short-wavelength modes which are assumed to make up the environment. Other authors have also examined the decoherence problem, often using the formalism of influence functionals \\cite{Feynman} (see, for example \\cite{Kiefer:1998qe,Sakagami:1987mp,Brandenberger:1990bx, Calzetta:1995ys,Lombardo:2005iz,Grishchuk:1989ss}). Unfortunately, the complexity of this formulation often forces calculations to be performed only for free fields, which couple to one another only by mixing in their mass or kinetic terms, leaving the uncertainty as to whether the results are restricted to the special choices which are made for calculational purposes. The main virtue of our treatment in terms of the master equation is that it allows us to identify the nature of the decoherence process in a controllable approximation which can encompass very general, realistic environment-system interactions.  Our analysis leads to the following results: \\begin{itemize} \\item We find that the quantity in the environment which controls the decoherence process is the autocorrelation function of the interaction Hamiltonian, $V(t)$, which couples the observed modes to the environment: $\\Tr_{\\rm env} [\\rho \\,\\delta V(t) \\, \\delta V(t')]$, where the trace is over the environmental sector, $\\rho$ denotes the system's density matrix and $\\delta V(t) = V(t) - \\Tr_{\\rm env}[\\rho \\, V(t)]$ represents the fluctuations of $V(t)$ (in the interaction representation). More complicated dependence on the properties of the environment only arise at higher order in $V(t)$. \\item We find that the generic situation is that the observed modes decohere into a density matrix of the form of eq.~\\pref{DecoheredState}, which is diagonal when expressed in terms of field eigenstates. This basis typically plays a special role because locality usually dictates that the important interactions between the modes of interest and their environment are diagonal when written in this basis. This point has previously been emphasized in ref.~\\cite{Kiefer:1998jk} \\item We show that the final stochastic probability distribution in eq.~\\pref{DecoheredState} is given in terms of the initial pure state by $P[\\varphi(x)] = |\\langle \\varphi(x)|\\psi\\rangle|^2$, and thereby justify more precisely the standard arguments which implicitly use the classical stochastic field distribution which reproduces the computed initial quantum fluctuations of the inflaton. \\item We find that short wavelength modes of the environment have the potential to decohere longer-wavelength modes, but need not do so. In particular they should not do so if the environment is prepared with its modes in their adiabatic vacua. It is this observation which explains why we are able to measure quantum states at all in the lab, given that there are always an enormous number of short-wavelength modes to which any given experiment does not have access. \\item We find that for simple choices for the short-distance environment there is ample time for decoherence to occur during or after inflation, even if the modes of interest are coupled to the environment with gravitational strength. Since all modes couple gravitationally, this shows that decoherence is ubiquitous for modes produced during inflation which are now observable on the largest scales. In the special case that decoherence arises due to Planck-suppressed couplings between super-Hubble inflaton modes and thermal sub-Hubble modes after reheating, decoherence is only now beginning to occur for those modes which cross the horizon during the recombination epoch, raising the intriguing possibility that this might have observable implications. \\end{itemize}  We next present these results in more detail. We do so by first briefly reviewing our approach to decoherence with an emphasis on the limits of validity of our calculations. We then examine two candidate environments, both of which are known to exist in most inflationary scenarios: ($i$) the short-wavelength modes of the metric and inflaton themselves, during inflation and afterwards (for which we find our calculations break down, although we give arguments as to why decoherence from this source may be unlikely); and ($ii$) those short-wavelength modes which are the first to thermalize after inflation ends (for which we argue our calculations apply, and show that decoherence could reasonably occur in the time available).  Our purpose is not to argue that either of these sources must provide the dominant source of decoherence, but rather to argue that our calculational tools suffice to identify that environments exist which have sufficient time to decohere initial quantum fluctuations. Our hope is that these tools can be used to compare potential sources of decoherence, with a view to identifying those which are the most important and how the detailed features of decoherence depend on the nature of the environment assumed. In particular, we hope to find observational consequences of the decoherence process which might be possible to use to test further the inflationary paradigm for generating primordial density fluctuations. We group some of the details of the calculations into appendices. ",
        "conclusions": "In this paper we present a formalism within which it is possible to follow explicitly the development of decoherence within a system's long-wavelength modes due to their interactions with various kinds of short-wavelength physics. The formalism relies for its validity on there being a large hierarchy between the correlation-time for fluctuations in the short-distance sector and the time-frame for evolution of the long-wavelength modes.  By applying this formalism to several possible choices for the short-distance environment we show that there is ample time during inflation for decoherence to occur between reheating and horizon re-entry for the modes of interest for CMB observations. We emphasize that our calculations do not yet establish what the most important source of decoherence might be for inflationary perturbations. In particular, we cannot compute with these tools how decoherence proceeds due to environments whose correlation times are not much shorter than the Hubble scale, or for short-wavelength modes prepared within their Bunch-Davies vacuum during inflation itself, although, we argue that decoherence during inflation by short-wavelength modes in their vacuum is unlikely, given the absence of such decoherence in non-cosmological situations.  Rather, our intent is merely to show that plausible sources could easily have sufficed to decohere the observed modes. Interestingly, if the decoherence is due to gravitational-strength interactions with the thermal (or other mixed) state produced during reheating, it could well be that the process is only just occurring for modes whose amplitudes differ by $O(H)$ as they re-enter the horizon near radiation-matter equality. If this should prove to be the most important source of decoherence we believe it would be worth exploring whether this might have observable consequences.  We argue that the decohered system tends to a final density matrix which is diagonal in a basis which also diagonalizes the relevant interaction Hamiltonian. On grounds of locality this is typically a basis for which the long-distance field operators themselves are diagonal. Since such a density matrix corresponds physically to the establishment of a stochastic distribution of classical field configurations, it is precisely the kind of late-time state which is implicitly assumed in standard calculations of the impact of inflationary fluctuations on the CMB.  \\medskip\\begin{center} {\\it Note Added:} \\end{center}  While preparing this manuscript we discovered a paper, ref.~\\cite{pm}, posted to the arXiv by one of our collaborators on this work. It describes an analysis which overlaps the one presented here on the calculation of the effects of decoherence by squeezed states during inflation. We believe that it differs in its conclusions due to errors in the calculations presented there, including (but not restricted to) the application of the master-equation formalism outside of its domain of validity."
    },
    "0601/astro-ph0601652_arXiv.txt": {
        "abstract": "Recent discovery of the afterglow emission from short gamma-ray bursts suggests that binary neutron star or black hole-neutron star binary mergers are the likely progenitors of these short bursts. The accretion of neutron star material and its subsequent ejection by the central engine implies a neutron-rich outflow. We consider here a neutron-rich relativistic jet model of short bursts, and investigate the high energy neutrino and photon emission as neutrons and protons decouple from each other. We find that upcoming neutrino telescopes are unlikley to detect the 50 GeV neutrinos expected in this model. For bursts at $z\\sim 0.1$, we find that {\\it GLAST} and ground-based Cherenkov telescopes should be able to detect prompt 100 MeV and 100 GeV photon signatures, respectively, which may help test the neutron star merger progenitor identification. ",
        "introduction": "The nature and origin of short gamma-ray bursts (SGRBs), with $t_\\gamma \\lesssim 2$ s duration, were largely conjectural until recently. This changed dramatically with the detection of the afterglow and the identification of the host galaxy of GRB 050509b at a redshift $z=0.226$ \\citep{getal05}, as well as subsequent observations of other short burst afterglows, e.g., GRB 050709 \\citep{fetal05, vetal05}, GRB 050724 \\citep{betal05}, etc. These observations revealed that the short bursts are located mainly in elliptical hosts, while some are in spirals or irregulars, as expected from an old population such as NS-NS or NS-BH mergers. They also provided, for the first time, details about the X-ray, and in some cases optical and radio afterglow lightcurves, including in two cases evidence for a jet break, qualitatively similar to the relativistic jet afterglows of long GRBs [see, e.g., \\cite{zm04} for reviews of long GRBs]. The average redshifts are lower than for long bursts, indicating an isotropic-equivalent energy of SGRBs approximately two orders of magnitude lower than for long ones.  We investigate here a model of SGRB jet which is initially loaded with neutrons and fewer protons, as expected from a neutron star composition.  High energy emission in this model is then expected from the photosphere of the relativistic jet outflow, including both thermal and non-thermal components, in addition to the radiation expected outside of it.  We describe our jet model in Section 2, high energy emission processes in Section 3, and expected neutrino and photon fluxes in Section 4 and 5 respectively. We summarize and discuss our results in Section 6. ",
        "conclusions": "If SGRBs are produced as a result of double neutron star or neutron star-black hole binary mergers, the neutron-rich outflow may result naturally in a high proton Lorentz factor $\\Gamma_{p,f}\\gtrsim 600$. Besides the usual non-thermal (e.g. internal shock) synchrotron spectrum, this implies the presence of a thermal photospheric component peaking around $\\sim 10$ MeV. The $n$-$p$ decoupling will lead to higher energy photons ($\\gtrsim 0.1$ GeV) from neutral and charged pion decay cascades, and neutrinos ($\\gtrsim 30$ GeV) directly from pion and muon decays. There may be additional X-ray emission when the neutron-decay proton shell with initial lower Lorentz factor impact with the decelerating proton shell initially moving with a high Lorentz factor \\citep{da06}. The afterglow light curves can also be affected by the neutron decay \\citep{dkk99b,rbr05}. Here, we have concentrated on the prompt ($t\\lesssim 2$ s) high energy (MeV to tens of GeV) signatures of neutron-rich outflows. For short bursts at $z\\sim 0.1$ the $\\sim 50$ GeV neutrinos are unlikely to be detectable with currently planned Gigaton Cherenkov detectors. However, the pion decay photons at $\\sim 60$ GeV escaping from whithin a $\\tau'_{\\rm Th}\\sim 1$ skin-depth of the photosphere should be detectable with high area, high duty cycle Cherenkov detectors such as {\\it Milagro}, at an estimated rate of 5 bursts per year. The pion decay photons created below the photosphere will result in a Comptonized cascade spectrum which, peaking in the range $\\gtrsim 0.3$ GeV, dominates the non-thermal shock synchrotron-IC spectrum, and should be detectable with {\\it GLAST}.  Both the prompt GeV photon components discussed here are essentially contemporaneous with the usual MeV range short GRB prompt emission, $t_\\gamma\\lesssim 2$ s.  Their detection or lack thereof could constrain the neutron content of the outflow and the nature of the short burst progenitors."
    },
    "0601/astro-ph0601378_arXiv.txt": {
        "abstract": "We present results from a short {\\em Chandra}/ACIS-S observation of NGC 7618, the dominant central galaxy of a nearby ($z$=0.017309, d=74.1 Mpc) group. We detect a sharp surface brightness discontinuity 14.4 kpc N of the nucleus subtending an angle of 130$^\\circ$ with an X-ray tail extending $\\sim$ 70 kpc in the opposite direction. The temperature of the gas inside and outside the discontinuity is 0.79$\\pm$0.03 and 0.81$\\pm$0.07 keV, respectively. There is marginal evidence for a discontinuous change in the elemental abundance ($Z_{inner}$=0.65$\\pm$0.25,$Z_{outer}$=0.17$\\pm$0.21 at 90\\% confidence), suggesting that this may be an `abundance' front. Fitting a two-temperature model to the ASCA/GIS spectrum of the NGC 7618/UGC 12491 pair shows the presence of a second, much hotter ($T$=$\\sim$2.3 keV) component. We consider several scenarios for the origin of the edge and the tail including a radio lobe/IGM interaction, non-hydrostatic `sloshing', equal-mass merger and collision, and ram-pressure stripping.  In the last case, we consider the possibility that NGC 7618 is falling into UGC 12491, or that both groups are falling into a gas poor cluster potential. There are significant problems with the first two models, however, and we conclude that the discontinuity and tail are most likely the result of ram pressure stripping of the NGC 7618 group as it falls into a larger dark matter potential. ",
        "introduction": "The complex cluster morphology seen in X-ray images and in galaxy distributions gave support to the hypothesis that large structures form hierarchically; that is, that small groups of galaxies merge to form low-mass subclusters, which then merge to form a massive rich cluster with the infalling groups aligned along large filaments. Galaxies and groups are the building blocks of the observable Universe and contain the bulk of the observable baryons.  Groups are estimated to contain a significant fraction, 20-30\\%, of the total matter in the Universe. Thus groups are important cosmological indicators of the distribution and properties of the dark matter.  However, because they are not as luminous as clusters, they have received less study than their more massive cousins.  Observations of rich clusters show many examples of pending or ongoing mergers of subclusters.  In particular in X-ray cluster catalogs, about 40\\% of rich clusters show substructure \\citep{jon84, jon99, moh95}. Virtually all stages of cluster mergers have been thoroughly investigated with both the {\\em Chandra} and XMM-Newton observatories \\citep{mar00,vik01,bri04,hen04}. However, less attention has been paid to the merging of groups and the formation of low-mass clusters due to their lower X-ray luminosity and the paucity of examples in the local Universe. To our knowledge, the only nearby example of the early stages of the merger of two roughly equal mass groups is the NGC 499/NGC 507 pair \\citep{kim95,kra03}. In the hierarchical scenario, group mergers represent a critical transitional phase in the formation of larger scale structure. An understanding of the group merger process is thus fundamental to understanding the growth of structure.  In this paper, we report results from analysis of a short {\\em Chandra}/ACIS-S observation of the nearby elliptical galaxy NGC 7618 ($z$=0.017309 or d$_L$=74.1 Mpc for WMAP cosmology \\citep{spe03} - 1$''$=350 pc).  The X-ray luminosity of NGC 7618 is $\\sim$7$\\times$10$^{42}$ ergs s$^{-1}$ in the 0.1-10 keV band, typical of groups, not isolated elliptical galaxies, although it appears to be optically isolated \\citep{col01}. We find a sharp surface brightness discontinuity in the X-ray emission north of of the nucleus of NGC 7618, and an extended tail to the south. We conclude that these features are either the result of a major group-group merger with UGC 12491, a group that lies 14.1$'$ on the sky from NGC 7618 at virtually identical redshift ($z$=0.017365) \\citep{ebl02}, or ram-pressure stripping due to infall of NGC 7618 into a larger gravitational potential (which may include UGC 12491). This pair has been poorly studied in large part because of its relatively low galactic latitude ($\\ell=105.575$, $b=-16.909$, $A_V$=0.97, $N_H$=1.19$\\times$10$^{21}$ cm$^{-2}$)). ",
        "conclusions": "We have observed a sharp surface brightness discontinuity in the X-ray emission from the hot gas in the NGC 7618 group, and an X-ray `tail' extending 70 kpc in the opposite direction in an 8 ks {\\em Chandra}/ACIS-S observation. Archival ASCA/GIS observations indicate the presence of a hotter (2.3 keV) component, although the morphology of this gas is poorly constrained. We conclude that there are three possible explanations for these features. First, the NGC 7618/UGC 12491 pair underwent a recent `near-miss' flyby. If this is the case, this pair is the nearest early-stage merger of two roughly equal mass groups. Second, UGC 12491 may be at the center of a cluster, and NGC 7618 is falling into it.  Third, NGC 7618 and UGC 12491 are both falling into a gas poor cluster with no dominant central elliptical galaxy. Whether the observed features are the result of a group-group merger, or the infall of two groups into a larger dark matter potential, the NGC 7618/UGC 12491 pair is one of the best examples of an ongoing merger in the local Universe. Deeper X-ray observations are required to better constrain the thermodynamic parameters of the gas in the central regions and the larger scale halo.  Radio observations will be critical in assessing the role of radio plasma/IGM interactions. If the observed X-ray features are the result of ram-pressure stripping or a merger interaction between the groups, the effects on the relic radio halo are likely to have been dramatic. A detailed optical census including velocities of the other galaxies in the both groups will be useful to constrain their dynamics and their relationship to the larger dark matter potential."
    },
    "0601/astro-ph0601187_arXiv.txt": {
        "abstract": "We have used a proper combination of multiband high-resolution {\\it HST-WFPC2} and wide-field ground based observations to image the galactic globular cluster NGC~6266 (M62).  The extensive photometric data set allows us to determine the center of gravity and to construct the most extended radial profile ever published for this cluster including, for the first time, detailed star counts in the very inner region.  The star density profile is well reproduced by a standard King model with an extended core ($\\sim 19''$) and a modest value of the concentration parameter ($c=1.5$), indicating that the cluster has not-yet experienced core collapse.  The millisecond pulsar population (whose members are all in binary systems) and the X-ray emitting population (more than 50 sources within the cluster half mass radius) suggest that NGC~6266 is in a dynamical phase particularly active in generating binaries through dynamical encounters. UV observations of the central region have been used to probe the population of blue straggler stars, whose origin might be also affected by dynamical interactions. The comparison with other globular clusters observed with a similar strategy shows that the blue straggler content in NGC~6266 is relatively low, suggesting that the formation channel that produces binary systems hosting neutron stars or white dwarfs is not effective in significantly increasing the blue straggler population. Moreover, an anticorrelation between millisecond pulsar content and blue straggler specific frequency in globular cluster seems emerging with increasing evidence. ",
        "introduction": "\\label{intro} In the last decades it has become obvious that some stars observed in Globular Clusters (GCs) are not the products of the passive evolution of isolated stars. Stellar encounters can produce binary systems and affect the evolution of the new produced binaries as well as any binaries that might have existed since the formation of the cluster. Thus the stellar content of a GC is intimately linked with the dynamical history of the cluster. In fact, the Color-Magnitude Diagrams (CMDs) of the core regions of GCs recently observed using HST-WFPC2 have shown the presence of a variety of exotic stellar objects, whose formation and evolution may be strongly affected by dynamical interactions. These include Low-Mass X-ray Binaries (LMXBs), (e.g. Heinke et al. 2001, Edmonds et al. 2003), Cataclysmic Variables (CVs) (e.g. Cool et al. 1998), and Blue Straggler Stars (BSSs) (e.g. the cases of 47 Tuc, Ferraro et al 2004; M3, Ferraro et al. 1997; M80, Ferraro et al., 1999a; M30, Guhathakurta et al., 1998, Ferraro et al. 2003a, hereafter F03).  BSS were first discovered by Sandage (1953) in the globular cluster M3, and manifest themselves in the CMD as an extension of the Main Sequence (MS) above the Turn-Off (TO), mimicking MS stars with larger initial masses. There are two viable proposed mechanisms for BSS generation: the first is mass exchange in a primordial binary system and the second is the merger of two stars induced by stellar interactions (either single or binary) in a dense stellar environment. The fact that BSSs have been found in all properly observed stellar systems make them a powerful tracer of dynamical history of the parent cluster.  Millisecond Pulsars (MSPs) may also be the by-product of collisions in the dense environment of a cluster. They form in binaries containing a neutron star (NS) which is eventually spun up through mass accretion from the evolving companion. Since they are point-like objects and (in many cases) extremely stable clocks, they are invaluable tools to investigate the binary evolution in a very dense stellar environment (e.g.  Rappaport et al. 2001). It is important to study the overall properties of GCs hosting MSPs and to compare the distributions of the objects which may be significantly affected by dynamical encounters. Hence we have undertaken a long-term project to observe in the optical band the GCs which are rich in known MSPs, with particular emphasis on the BSS population. After having examined 47~Tucanae (Ferraro et al. 2001, 2004; Mapelli et al. 2004) and NGC~6752 (Ferraro et al. 2003b, Sabbi et al.  2004), we now turn our attention to NGC~6266.  NGC~6266 (M62) is a high density ($\\log \\rho_c \\sim$ 5.34), highly reddened ($E(B-V)\\sim$ 0.47) GC (Jacoby 2002; Possenti et al. 2003, hereafter P03) and is one of the most massive ($M_v=-9.19$, Harris 1996{\\footnote{For all references to Harris (1996) we have used the updated data set at the web site http://www.physics.mcmaster.ca/Globular.html}}).  It has been classified as a possible Post Core Collapse (PCC) GC in the compilation by Djorgovski \\& Meylan (1993, hereafter DM93). Six binary MSPs have been recently discovered in this cluster (D'Amico et al. 2001a, 2001b; Jacoby et al. 2002; P03). NGC~6266 ranks fourth of the GCs in wealth of MSPs, after Terzan~5, 47~Tucanae and M15. Surprisingly, all MSPs in NGC~6266 are in binary systems (P03). As discussed in P03, the absence of known isolated MSPs in NGC~6266 cannot simply be ascribed to a selection effect since, for a given spin period and flux density, an isolated MSP is easier to detect than a binary MSP. If such a lack of isolated pulsars is not a statistical fluctuation, it must be somehow related to their formation mechanism and to the dynamical state of the cluster.  Moreover, recent deep {\\it Chandra} X-ray observations of NGC~6266 have revealed a very large number of X-ray sources---51 sources were detected within the cluster 1\\farcm 23 half mass radius (Pooley et al. 2003), indicating that an high number of cataclysmic (and/or interacting) binaries should be present. These observational facts may indicate that NGC~6266 is now in a dynamical state where the rate ${\\cal R}_{\\rm form}$ of formation (and of hardening) of binary systems containing a neutron star and/or a white dwarf is much larger than the rate of disruption ${\\cal R}_{\\rm disr}$ of such systems.  In order to investigate in more detail the dynamical properties of this cluster, we have used HST-WFPC2 high resolution images of the central core and ESO-WFI wide field observations to sample the more external regions.  The combination of these two data-sets allows us to derive an accurate star density profile over the entire cluster extension. \\S~\\ref{obs} describes the datasets and the reduction procedures, while \\S~\\ref{center} presents the derived CMD. \\S~\\ref{profile} describes the determination of the cluster center of gravity and the star density profile whereas \\S~\\ref{param} is focused on the determination of the new cluster parameters. \\S~\\ref{bss} is devoted to the study of the Blue Straggler population in the central regions of NGC~6266 and, finally we discuss the results in \\S~\\ref{disc}. ",
        "conclusions": "\\label{disc}  By adopting the cluster parameters computed in Section~\\ref{param}, we can now investigate some dynamical properties of NGC~6266. We computed the mean time interval between single-single ($t_{ss}$) and binary-binary ($t_{bb}$) collisions using equations 13 and 14 respectively from Leonard (1989):  $$ t_{ss}=5.5\\times10^9\\left(\\frac{1\\,{\\rm pc}}{r_c}\\right)^{3} \\left(\\frac{10^3\\,{\\rm pc}^{-3}}{n_0}\\right)^2\\times \\left(\\frac{V_{rms}}{5\\,{\\rm km\\,s}^{-1}}\\right)\\left(\\frac{0.5M_{\\odot}}{m_{*}}\\right) \\left(\\frac{0.5R_{\\odot}}{R_{*}}\\right)[yr], $$   $$ t_{bb}=1.7\\times10^7\\left(\\frac{1\\,{\\rm pc}}{r_c}\\right)^{3} \\left(\\frac{10^3\\,{\\rm pc}^{-3}}{n_0}\\right)^2\\times \\left(\\frac{V_{rms}}{5\\,{\\rm km\\,s}^{-1}}\\right)\\left(\\frac{0.5M_{\\odot}}{m_{*}}\\right) \\left(\\frac{1AU}{a}\\right)[yr], $$  \\noindent where $a$ is the initial semimajor axis of each binary which can be estimated for each cluster according to the relation (F03) $$ a=\\frac{GM}{9\\sigma^2_0}, $$ $n_0$ is the central number density (number of stars per parsec$^3$) and $V_{rms}$ is the relative root-mean-square velocity, which can be approximated with the central velocity dispersion ($\\sigma_0$, see Table 5). Leonard (1989) derived his equation (14) under the assumption that the binary frequency is 100\\% in the core and the star mass is $m_{\\ast}=0.5M_{\\odot}$. Assuming an average mass of $m_{\\ast}=0.2M_{\\odot}$ (Kroupa 2001) as suggested in F03, and $R_{\\ast}=0.5R_{\\odot}$, we computed the number of single-single ($ss$) and binary-binary ($bb$) encounters occurring per Gyr ($N_{ss,bb}$) in the core of NGC~6266 to be: $$ N_{ss,bb}=\\frac{10^9}{t_{ss,bb}} \\;\\;\\;\\;\\;\\;   [{\\rm encounters\\,Gyr^{-1}}]. $$ In Table~\\ref{bb} we give the results obtained for NGC~6266, compared to the 8 GCs which have been previously searched for BSS in the central regions by using the same observational strategy.  The values of the encounter rate are normalized to those obtained for 47 Tuc. NGC~6266 has a predicted {\\it{bb}} encounter rate $N_{bb}$ which is much greater than for any other analyzed GC. As shown in Table \\ref{bb}, the same result holds for the close encounters between two single stars leading to tidal capture. Both of these classes of dynamical interactions (together with the binary-single [$bs$] encounters) may lead to the merging of the two stars, thus potentially producing a BSS.  Since all collision modes should be effective in increasing the population of BSS, one would expect to see a relatively larger specific frequency of BSS in NGC~6266. This expectation is not born out by the results of Table \\ref{bss_freq}. The discrepancy might be reduced if the binary fraction in NGC~6266 were significantly smaller than in the other GCs, but this hypothesis can hardly be reconciled with the estimated high rate of binary formation ${\\cal R}_{\\rm form}$ computed for this cluster (see later).  Perhaps the results given in Tables \\ref{bss_freq} and \\ref{bb} can be understood if the BSS in NGC~6266 result from primordial binaries which have survived in the outer part of the cluster eventually wandering into the center where interactions drive them to merge (as we suggested for 47~Tuc, see Ferraro et al. 2004).  On another side, the very high rate of binary-binary and binary-single star interactions in NGC~6266 may well explain why all of its known MSPs are in binary systems.  In fact, P03 suggested that NGC~6266 is in a particular dynamical phase in which the rate of binary system formation ${\\cal R}_{\\rm form}$ is larger than their destruction rate (${\\cal R}_{\\rm disr}$). Following the prescription of Verbunt (2003), we recomputed those ratios for all clusters listed in Table~\\ref{coll}.  According to Verbunt (2003) ${\\cal R}_{\\rm form}\\propto\\rho_0^{1.5}r_c^2$ and ${\\cal R}_{\\rm disr}\\propto\\rho_0^{0.5}r_c^{-1}$, where $\\rho_0$ is the central density (expressed in $M_{\\odot}\\,{\\rm pc}^{-3}$) and $r_c$ is the core radius of the cluster. In Table~\\ref{coll} we give the values of ${\\cal R}_{\\rm form}$ and ${\\cal R}_{\\rm disr}$ normalized to 47 Tuc (as done by P03). A quite useful quantity is the ratio ${\\cal R}_{\\rm form}/{\\cal R}_{\\rm disr}$ which quantifies the binary survival rate in each environment. NGC~6266 shows, by far, the largest value of this ratio among the clusters listed in Table~\\ref{coll}. The 2nd ranked cluster is M3 in which all the MSPs discovered so far are also in binaries.  In summary, we have been able to recompute the ratio ${\\cal R}_{\\rm form}/{\\cal R}_{\\rm disr}$ and to show that {\\it (i)} its high value nicely agrees with the hypothesis (supported by observations of MSPs and X-ray sources) that NGC~6266 has experienced (or is experiencing) a phase of very high production of binary systems. Additionally, we show that {\\it (ii)} the highly effective binary production phase is occurring while the cluster has not yet undergone the collapse of the core, and that {\\it (iii)} the formation channel that produces a wealth of binary systems hosting neutron stars or white dwarfs is not efficient in producing BSS.  We can also compare NGC~6266 with the sample of GCs searched for BSS with an approach similar to that used in this work. In particular, inspection of Tables \\ref{bss_freq} and \\ref{coll} shows that in GCs hosting MSPs the specific frequency of BSS never exceeds 0.28. This sample includes the post-core collapsed globular cluster, NGC~6752, in which only a few BSS have been found (Sabbi et al. 2004). Conversely, up to now no MSPs have been found in M80 and NGC~288, the two clusters with the largest specific frequency of BSS measured so far (probably generated by different channels: primordial binaries in NGC~288 and collisional binaries in M80).  The difficulty in modeling the observational biases involved in discovering radio pulsars in GCs prevents any firm conclusion to date, but an anticorrelation between MSP content and BSS specific frequency seems emerging. GCs with large populations of BSS seem to have a small population of MSP."
    },
    "0601/astro-ph0601464_arXiv.txt": {
        "abstract": "The 2-10 keV continuum of AGN is generally well represented by a single power law. However, at smaller energies the continuum displays an excess with respect to the extrapolation of this power law, called the ``soft X-ray excess''. Until now this soft X-ray excess was attributed, either to reflection of the hard X-ray source by the accretion disk, or to the presence of an additional comptonizing medium, giving a steep spectrum. An alternative solution proposed by Gierli\\' nski and Done (2004) is that a single power law well represents both the soft and the hard X-ray emission and the impression of the soft X-ray excess is due to absorption of a primary power law by a relativistic wind. We examine the advantages and drawbacks of reflection versus absorption models, and we conclude that the observed spectra can be well modeled, either by absorption (for a strong excess), or by reflection (for a weak excess). However the physical conditions required by the absorption models do not seem very realistic: we would prefer an ``hybrid model''. ",
        "introduction": "Now more than 50\\% of well studied Seyfert 1 galaxies and many quasars are known to possess absorbers (e.g., Blustin et al. 2005). One main issue is to explain the apparent change of slope in the overall X-ray spectrum at $\\sim$ 1 keV in Sy1s and QSOs. When fitting an observed X-ray spectrum with a power law plus absorption plus (eventually) the Compton reflection component plus (eventually) the iron line and (eventually) narrow spectral features, the model usually underpredicts the observed spectrum in the soft X-ray range. An additional component -- a soft X-ray excess -- is needed (Wilkes \\& Elvis 1987). Apart from an usual additional continuum or a strongly ionized reflection, this component is well modeled by absorption of an originally rather soft power law intrinsic spectrum due to an absorber having a random or bulk velocity of several thousands of km~s$^{-1}$ (Gierli\\' nski \\& Done 2004). We consider the advantages and drawbacks of the reflection vs. absorption models (Chevallier et al.~2005), using our code TITAN (Dumont et al.~2003). ",
        "conclusions": "Absorption models could account for some strong soft X-ray excesses, but require a kind of ``fine tuning\" in order to constrain the 1 keV trough, which otherwise could have any strength (e.g., constant density models). We have suggested a medium in total -- gas and radiation -- pressure equilibrium, which leads to a maximum intensity of the trough, as well as a ``universal\" shape of this maximum trough, due to the thermal instability mechanism. A complete grid of constant total pressure models, very demanding in computation time, is necessary to pursue this study. In the absorption model, either a thick accretion flow, or a relativistic wind is required. None of them seem realistic from a physical point of view, and moreover both require an additional source of UV emission, like a geometrically thin accretion disk. On the other hand, the traditional reflection models involving a hot corona emitting X-rays which are reflected by a cold disk cannot account for the observations, unless the X-ray source is hidden from our view. Therefore we favor an ``hybrid\" model, where the primary UV-X source could be produced by a disk-corona system, and then absorbed by a modest relativistic wind."
    },
    "0601/astro-ph0601522_arXiv.txt": {
        "abstract": "Where do quasars reside? Are quasars located in environments similar to those of typical $L^*$ galaxies, and, if not, how do they differ? An answer to this question will help shed light on the triggering process of quasar activity. We use the Sloan Digital Sky Survey to study the environment of quasars and compare it directly with the environment of galaxies. We find that quasars ($M_i \\le -22$, $z \\le 0.4$) are located in higher local overdensity regions than are typical $L^*$ galaxies. The enhanced environment around quasars is a local phenomenon; the overdensity relative to that around $L^*$ galaxies is strongest within $100\\,$kpc of the quasars. In this region, the overdensity is a factor of 1.4 larger than around $L^*$ galaxies. The overdensity declines monotonically with scale to nearly unity at $\\sim1\\,h_{70}^{-1}\\,$Mpc, where quasars inhabit environments comparable to those of $L^*$ galaxies. The small-scale density enhancement depends on quasar luminosity, but only at the brightest end: the most luminous quasars reside in higher local overdensity regions than do fainter quasars. The mean overdensity around the brightest quasars ($M_i \\le -23.3$) is nearly three times larger than around $L^*$ galaxies while the density around dimmer quasars ($M_i = -22.0$ to $-23.3$) is $\\sim1.4$ times that of $L^*$ galaxies. By $\\sim0.5\\,$Mpc, the dependence on quasar luminosity is no longer significant. The overdensity on all scales is independent of redshift to $z \\le 0.4$. The results suggest a picture in which quasars typically reside in $L^*$ galaxies, but have a local excess of neighbors within $\\sim0.1\\,-\\,0.5\\,$Mpc; this local density excess likely contributes to the triggering of quasar activity through mergers and other interactions. ",
        "introduction": "\\label{sec.intro}  For more than two decades, significant effort has been spent attempting to understand the triggering mechanism of quasar activity, as well as the relation between quasars and their host galaxies. Since \\citet{bell_1969}, it has become widely accepted that quasars are fueled by accretion of gas onto super-massive black holes. Observations have shown that a number of nearby galaxies have a central black hole whose mass correlates with the luminosity of the spheroid of the host galaxy. This connection suggests that the formation of the black hole is linked to the formation of the galaxy which, in turn, is known to strongly depend on its environment. To develop a better understanding of the quasar phenomenon, it is therefore important to investigate and quantify the relation between quasars and their environments. Despite the importance of this issue, our knowledge of the quasar environment is still limited.  Quasar environments have been studied on different scales ranging from those of the host galaxy to those of large scales. Such studies have provided important but controversial results regarding the environment of quasars. It has been known for more than three decades that quasars are associated with enhancements in the spatial distribution of galaxies \\citep{bahcall_1969}. Studies have shown that, in the nearby universe, quasars reside in environments ranging from small to moderate groups of galaxies rather than in rich clusters (\\citealt{bahcall_1991b,fisher_1996,mclure_2001}).  Early observations of quasar environments \\citep{stockton_1978, yee_1984, yee_1987, boyle_1988, smith_1990, ellingson_1991}, revealed a positive association of bright quasars with neighboring galaxies at a level somewhat higher than that of normal galaxies and comparable to the environment of small- to intermediate-richness groups of galaxies \\citep[e.g.,][]{bahcall_1991}. Observations of quasar environment from the Hubble Space Telescope snapshot survey \\citep{bahcall_1997} further support this density enhancement around bright quasars. All these observations focused on small scales, typically within $\\sim 0.5\\,$Mpc of the quasars, and used relatively small samples of objects.  Early observations of the clustering properties of quasars themselves, as measured by the quasar auto-correlation function, suggest that quasars are significantly more strongly clustered than galaxies on scales up to $10\\,$Mpc and greater \\citep[e.g.,][]{shaver_1988, shanks_1988, chu_1988, chu_1989, crampton_1989}, but less clustered than rich clusters of galaxies \\citep[e.g.,][]{bahcall_1991}. This finding suggests that quasars are located in high overdensity regions, more so than $L^*$ galaxies, since higher overdensity regions are clustered more strongly than lower overdensity regions \\citep[e.g.,][]{bahcall_1983, kaiser_1984, bardeen_1986}. An overdense environment would indeed be expected if the quasar activity was triggered by galaxy interactions.  On the other hand, new generation surveys, such as the Two Degree Field (2dF) and the Sloan Digital Sky Survey (SDSS), have given rise to different results. Using significantly larger complete samples of quasars and galaxies, these surveys have shown that on large scales, i.e. from 1 to $10\\,$Mpc, the quasar-galaxy cross-correlation and the quasar auto-correlation are comparable to the correlation function of $L^*$ galaxies. This suggests, in conflict with previous results, that quasars and Active Galactic Nuclei (AGNs) inhabit environments similar to those of $L^*$ galaxies \\citep[e.g.,][]{smith_1995, croom_1999, croom_2003, miller_2003, kauffmann_2004, wake_2004}.  The results also suggest that the quasar correlation function does not depend significantly on either quasar luminosity or redshift within the ranges studied.  Recent work on sub-Mpc scales using the SDSS to find quasar pairs suggests that the quasar-quasar auto-correlation function may be enhanced relative to the galaxy-galaxy distribution \\citep{hennawi_2005}, consistent with the earlier results on small scales, as discussed above.  In this paper we use the SDSS survey to determine the galaxy environment around quasars as a function of scale. The SDSS is uniquely suited for this investigation: it is the largest complete survey available of both galaxies and quasars, carried out in a well-calibrated, self-consistent manner. The data used in this study covers 4000 deg$^2$, with $\\sim2\\times10^3$ quasars of redshift $z\\le0.4$ and ten million photometric galaxies to a magnitude limit of $i = 21$. We use these data to determine the mean galactic environment around quasars as a function of quasar luminosity and redshift. For comparison, the same analysis is then repeated to find the local environment around $10^5$ spectroscopic galaxies in the SDSS area, as well as around random positions in the survey. All the analyses are carried out using the same ten million photometric galaxies. This technique allows a direct comparison between the environment around quasars with that around random points as well as with the environment around $L^*$ galaxies, thereby minimizing potential selection effects and systematics.  The outline of the paper is as follows: we discuss the data in Section~\\ref{sec.data}, the analysis in Section~\\ref{sec.analysis}, and the results in Section~\\ref{sec.results}. The conclusions are summarized in Section~\\ref{sec.conclusions}. Throughout this paper, we use a cosmological model with $H_0\\,=\\,70\\,{\\rm km^{-1}\\,Mpc^{-1}}$, $\\Omega_M\\,=\\,0.3$, and $\\Omega_{\\Lambda}\\,=\\,0.7$ for both absolute magnitudes and distance measures. All distances are measured using comoving coordinates. ",
        "conclusions": "\\label{sec.conclusions}  We use the SDSS data to investigate the local environment of quasars and compare it with the environment around galaxies and around random positions. This study provides a direct comparison between the environment of quasars and galaxies.  We use bright quasars with $-24.2 \\le M_i \\le -22$ ($h=0.7$) and redshift $z \\le 0.4$ (2028 quasars over $4000\\,$deg$^2$) to study the density of photometric galaxies with $i = 14$ to 21 (sample of $\\sim10^7$ galaxies) located within $25\\,$kpc to $1\\,$Mpc ($h=0.7$) of the parent quasars. We compare the results with the same analysis carried out at random positions in the SDSS survey, $\\ngr$; this yields the galaxy overdensity, over random, around quasars, i.e., $\\ngq/\\ngr$. We investigate this overdensity as a function of quasar redshift and luminosity. The same analysis is then repeated for determining the overdensity around $10^5$ spectroscopic galaxies in the SDSS data ($i \\le 18.5$) in the same redshift range ($z = 0.08$ to $0.4$). This allows a direct comparison of the quasar overdensity to the overdensity around galaxies. The overdensities are studied as a function of scale, luminosity, and redshift. This comparison of quasar environment with that of galaxies provides a self-consistent comparison of the host environments.  Our results are summarized below.  1. At all radii, from $25\\,$kpc to $\\sim1\\,$Mpc, quasars are located in higher density environments than are $L^*$ galaxies. The overdensity around quasars relative to that of $L^*$ galaxies increases with decreasing scale: the overdensity is greatest closest to the quasars. At a distance of $40\\,$kpc of the quasars, the mean overdensity for the brightest quasars ($-24.2 \\le Mi \\le -23.3$, $z \\le 0.4$) is nearly a factor of three times larger than the overdensity around $L^*$ galaxies. The mean overdensity around the brightest quasars relative to that of $L^*$ galaxies decreases with increasing scale to a value of 2.2 within $0.1\\,$Mpc, 1.2 at $0.3\\,$Mpc, and approaches unity at $\\sim0.5 - 1\\,$Mpc. On these larger scales, quasars reside in environments similar to those of $L^*$ galaxies.  2. The brightest quasars are found to be located in higher overdensity regions than are fainter quasars, especially at small separations ($< 0.1\\,$Mpc). On larger scales, bright and faint quasars live in similarly dense environments.  3. The mean overdensity around quasars 0 ($<~\\ngq/\\ngr~> = 2.12\\pm0.08$ within $0.1\\,$Mpc, decreasing to $1.13\\pm0.01$ within $1\\,$Mpc) is independent of redshift for $z \\le 0.4$. The mean overdensity is also independent of luminosity except for the brightest quasars, which are located in higher density environments. This dependence of quasar environment on luminosity, showing enhancement only for the most luminous quasars, is consistent with recent galaxy merging models of quasars \\citep{hopkins_2005}.  4. The enhanced mean overdensity around quasars is observed to be a local phenomenon, affecting mostly the $\\sim$ $0.1\\,$Mpc region closest to the quasars. On these scales, very close to the quasars, the high overdensity of galaxies likely affects the formation and triggering of the quasar activity through mergers and other interactions. On scales of $\\sim 1\\,$Mpc, the quasars inhabit similar environments to those of normal $L^*$ galaxies."
    },
    "0601/astro-ph0601008_arXiv.txt": {
        "abstract": "Numerical simulations of the intergalactic medium have shown that at the present epoch a significant fraction ($40-50\\%$) of the baryonic component should be found in the ($T\\sim 10^6$K) Warm-Hot Intergalactic Medium (WHIM) - with several recent observational lines of evidence indicating the validity of the prediction. We here recompute the evolution of the WHIM with the following major improvements: (1) galactic superwind feedback processes from galaxy/star formation are explicitly included; (2) major metal species (O V to O IX) are computed explicitly in a non-Equilibrium way; (3) mass and spatial dynamic ranges are larger by a factor of 8 and 2, respectively, than in our previous simulations. Here are the major findings: (1) galactic superwinds have dramatic effects, increasing the WHIM mass fraction by about $20\\%$, primarily through heating up warm gas near galaxies with density $10^{1.5}-10^4$ times the mean density. (2) the fraction of baryons in WHIM is increased modestly from the earlier work but is $\\sim 40-50\\%$. (3) the gas density of the WHIM is broadly peaked at a density $10-20$ times the mean density, ranging from underdense regions to regions that are overdense by $10^3-10^4$. (4) the median metallicity of the WHIM is $0.18\\zsun$ for oxygen with 50\\% and $90\\%$ intervals being (0.040,0.38) and (0.0017,0.83). ",
        "introduction": "Cosmological hydrodynamic simulations have strongly suggested that most of the previously ``missing\" baryons may be in a gaseous phase in the temperature range $10^5-10^7$K and at moderate overdensity (Cen \\& Ostriker 1999, hereafter \"CO\"; Dav\\'e 2001), called the warm-hot intergalactic medium (WHIM), with the primary heating process being hydrodynamic shocks from the formation of large-scale structure at scales currently becoming nonlinear. The reality of the WHIM has now been quite convincingly confirmed by a number of observations from HST, FUSE, Chandra and XMM-Newton (Tripp, Savage, \\& Jenkins 2000; Tripp \\& Savage 2000; Oegerle \\etal 2000; Scharf \\etal 2000; Tittley \\& Henriksen 2001; Savage \\etal 2002; Fang \\etal 2002; Nicastro \\etal 2002; Mathur, Weinberg, \\& Chen 2002; Kaastra \\etal 2003; Finoguenov, Briel, \\& Henry 2003; Sembach \\etal 2004; Nicastro \\etal 2005).  In addition to shock heating, feedback processes following star formation in galaxies can heat gas to the same WHIM temperature range. What is lacking theoretically is a satisfactory understanding of the known feedback processes on the WHIM and how the WHIM may be used to understand and calibrate the feedback processes. Another unsettled issue is how the predicted results will change if one has a more accurate, non-equilibrium calculation of the major metal species, such as O VI, O VII, O VIII, since time scales for ionization and recombination are not widely separated from the Hubble time scale. We have investigated these important processes. Additional improvements include significantly larger dynamic ranges of the simulation, a WMAP normalized cosmological model and an improved radiative transfer treatment. Our current work significantly extends previous theoretical works by our group and others (CO; Dav\\'e 2001; Cen \\etal 2001; Fang \\etal 2002; Chen \\etal 2003; Furlanetto \\etal 2004,2005a,b; Yoshikawa \\etal 2003; Ohashi \\etal 2004; Suto \\etal 2004; Fang \\et al. 2005). In this paper we focus on the feedback effects due to star formation and in a companion paper (Cen \\& Fang 2005) we will present additional effects due to non-equilibrium treatments of metal species and observables of the WHIM. The outline of this paper is as follows: the simulation details are given in \\S 2; in \\S 3 we give detailed results and discussion and conclusions are presented in \\S 4. ",
        "conclusions": "Numerical simulations of the intergalactic medium have shown that at the present epoch a significant fraction ($40-50\\%$) of the baryonic component should be found in the ($T\\sim 10^{5-7}$K) Warm-Hot Intergalactic Medium (WHIM) - with several recent observational lines of evidence indicating the validity of the prediction. We here recompute the evolution of the WHIM with the following major improvements: (1) galactic superwind feedback processes from galaxy/star formation are explicitly included; (2) major metal species (O IV to O IX) are computed explicitly in a non-equilibrium way; (3) mass and spatial dynamic ranges are larger by a factor of 8 and 2, respectively, than our previous simulations.  Our significantly improved simulations confirm previous conclusions based on earlier simulations: {\\it nearly one half of all baryons at the present epoch should be found in the WHIM} - a filamentary network in the temperature range of $10^5-10^7$~K.  Here are the major findings: (1) Overall, the fraction of baryons in WHIM is slightly increased from the earlier work but consistent with the fraction at $\\sim 40-50\\%$. (2) galactic superwinds have three significant effects, first by increasing the WHIM mass fraction by about $20\\%$, through shock heating photoionized gas adjacent to filaments, second by contaminating nearby gas with extra metals and additional heating, and third by dispersing much of the metals within galaxies. (3) the gas density of the WHIM is broadly peaked at a density $10-20$ times the mean density ranging from underdense regions to regions overdense by $10^3-10^4$. (4) the median metallicity of the WHIM is $0.18\\zsun$ for oxygen with 50\\% and $90\\%$ intervals being $(0.040,0.38)\\zsun$ and $(0.0017,0.83)\\zsun$.   The physics behind this robust conclusion is largely dictated by the dominance of gravitational heating when large-scale structures begin to collapse in the recent past in the cosmic history. Simply put, the average temperature of the intergalactic medium at any epoch is determined by the mass scale that goes nonlinear at that epoch, which, in present time, is close to the $8h^{-1}$Mpc scale which fixes the abundance of rich clusters. This explains the relative insensitivity of the computed mass fraction in the WHIM on cosmological parameters, so long as each model is properly normalized to reproduce the well determined abundance of observed rich clusters of galaxies today; only a very small extrapolation is needed to go from this well observed scale to the nonlinear scale, so the estimated temperature of collapsed regions that we find will be common to all models based on the gravitational growth of structure - as normalized to local cluster observations.    It was not entirely clear how important feedback processes from galaxy formation on IGM are, based on our early simulations compared to the simulations by others (Dav\\'e \\etal 2001). In this work we demonstrate quantitatively, for the first time, this effect. We find that galactic superwinds generated collectively by supernova explosions in galaxies, while still subdominant to gravitational heating of large-scale structure formation, have important effects on the WHIM as well as other components of the IGM. The mass fraction of the WHIM is increased by about 20\\%, when GSW are included. We show that this increase in the WHIM mass fraction comes largely at the expense of the warm IGM phase over a broad range $30-10^4$ times the mean density, which is often independently invoked to suppress star formation to alleviate the overcooling problem. A perhaps more important effect of GSW is to transport metals to the IGM from inside galaxies to a distance of several hundred kiloparsecs up to about one megaparsec, a necessary process to enrich the IGM. It is therefore likely that detailed, joint analyses and comparisons with observations of the detailed structure of the WHIM density, temperature and metal abundances should be able to provide useful information on the GSW, which currently can only be modeled crudely. Nevertheless, with the adopted approximate treatment of GSW, we are able to show the significant effects on many observables, including the properties of major absorption lines and emissions lines in UV and soft X-ray."
    },
    "0601/astro-ph0601003.txt": {
        "abstract": "{ We constrain the evolution of the galaxy mass and luminosity functions from the analysis of (public) multi-wavelength data in the Chandra Deep Field South (CDFS) area, obtained from the GOODS and other projects, and including very deep high-resolution imaging by HST/ACS. Our reference catalogue of faint high-redshift galaxies, which we have thoroughly tested for completeness and reliability, comes from a deep ($S_{3.6}\\geq 1\\ \\mu$Jy) image by IRAC on the Spitzer Observatory. These imaging data in the field are complemented with extensive optical spectroscopy by the ESO VLT/FORS2 and VIMOS spectrographs, while deep K-band VLT/ISAAC imaging is also used to derive further complementary statistical constraints and to assist the source identification and SED analysis.    We have selected a highly reliable IRAC 3.6$\\mu$m sub-sample of 1478 galaxies with $S_{3.6}\\geq 10\\ \\mu$Jy, 47\\% of which have spectroscopic redshift, while for the remaining objects both COMBO-17 and $Hyperz$ are used to estimate the photometric redshift. This very extensive dataset is exploited to assess evolutionary effects in the galaxy luminosity and stellar mass functions, while luminosity/density evolution is further constrained with the number counts and redshift distributions. The deep ACS imaging allows us to differentiate these evolutionary paths by morphological type, which our simulations show to be reliable at least up to $z\\sim 1.5$ for the two main early- (E/S0) and late-type (Sp/Irr) classes.  These data, as well as our direct estimate of the stellar mass function above $M_\\ast h^2=10^{10} M_\\odot$ for the spheroidal subclass, consistently evidence a progressive dearth of such objects to occur starting at $z\\sim 0.7$, paralleled by an increase in luminosity. A similar trend, with a more modest decrease of the mass function, is also shared by spiral galaxies, while the irregulars/mergers show an increased incidence at higher z. Remarkably, this decrease of the comoving density with redshift of the total population appears to depend on galaxy mass, being stronger for moderate-mass, but almost absent until $z=1.4$ for high-mass galaxies, thus confirming previous evidence for a \"downsizing\" effect in galaxy formation. % %Simple evolutionary models, fitting the fast convergence %of the number counts and redshift distributions, and the evolutionary %mass function suggest the main episodes for spheroidal build-up %(of either old or newly-formed stellar populations) %to happen between $z\\sim 2$ and $z\\leq 1$ for such field population. % Our favoured interpretation of the evolutionary trends for the two galaxy categories is that of a progressive morphological transformation (due to gas exhaustion and, likely, merging) from the star-forming to the passively evolving phase, starting at $z\\geq 2$ and keeping on down to $z\\sim 0.7$. The rate of this process appears to depend on galaxy mass, being already largely settled by $z\\sim 1$ for the most massive systems. ",
        "introduction": "\\label{intro}  Subject of active, as much as inconclusive, investigation during the last 40 years or so, the cosmological origin of the Hubble galaxy morphological sequence can now be very effectively constrained by combining the unique imaging capabilities of HST/ACS with the wide IR multi-wavelength coverage offered by the Spitzer Space Telescope and the remarkable photon-collecting power and multiplexing of spectrographs on large ground-based telescopes (VLT, Keck). As much a complex process as it might have been -- involving both gravity and hydrodynamics (see e.g. Baugh et al. 2005), and possibly other physical processes such as black-hole formation and accretion, tidal interactions and merging, and feedback from stellar and nuclear activity (Springel et al. 2005) -- we have now a definite chance to observe it in operation.  At the current stage, however, the subject remains still rather controversial. While slow infall of primordial gas may explain disk formation in a relatively simple way, (e.g. Mo, Mao \\& White 1998), we still lack an adequate understanding of spheroid formation. On one side, the homogeneity of the early-type population and tightness of the fundamental plane might suggest that these galaxies have formed from a single monolithic collapse, an early aggregation of lumps of gas turning into stars in the remote past ($z_{form} \\ge$3) via a huge burst-like episode followed by quiescence (Eggen et al. 1962; Larson et al. 1975; Chiosi and Carraro 2002).  This however is in apparent contradiction with recently favoured models of hierarchical galaxy formation which postulate that early-type galaxies are assembled at later times by stochastic merging of lower-mass galaxies, either accompanied by strong star formation activity (e.g. White et al. 1978; White et al. 1991; Somerville \\& Primack 1999; Cole et al. 2000), or through more \"silent\" dry-merging and dynamical aggregation (Bell et al. 2005a, Tran et al. 2005). In such case, ellipticals would be formed over timescales comparable to the Hubble time, with a major fraction of the mass assembly taking place around $z \\sim  1$ (e.g. Somerville et al. 2001), and virtually all massive galaxies disappearing by $z \\geq 1.5$. Benson et al. (2002) find popular hierarchical models to produce as many spheroids with highly inhomogeneous colors as observed locally, but to underpredict the proportion of homogeneous, passive objects at redshifts $z \\sim 1$. This suggests that while the star formation rate in spheroidals at low redshifts ($z \\le 1$) is predicted correctly, the formation rate at higher redshifts is underestimated. On the other hand, recent results from the K20 project (see Daddi et al. 2004a and references therein) claim that semianalytic models underestimate the number of massive galaxies at $z \\sim 2$ by about a factor of 30 and suggest that the assembly of massive galaxies took place at substantially earlier epochs than predicted by these models.  Observational constraints on the star formation history have been inferred from the broad-band colors, line strength indices and stellar chemical abundances. When referred to massive ellipticals, these observations often suggest that the bulk of stars might have been formed in a remote past. However, some secondary activity of star formation in the recent past is also evident: nearby ellipticals show a large variety of morphological and kinematical peculiarities (e.g. Longhetti et al. 2000) and a considerable spread of stellar ages, particularly for the field population (Thomas et al. 2005). Strong evolution in the population of early-type galaxies has been reported by Kauffmann, Charlot \\& White (1996) and Kauffmann \\& Charlot (1998), which has been considered to support the hierarchical galaxy formation models.  Published results from high redshift galaxy surveys appear not unfrequently in disagreement with each other, and conflicting conclusions are reported (see Faber et al. 2005 for a recent review about galaxy evolution at $z<1$). This is partly due to the small sampled areas and the corresponding substantial field-to-field variance. However, a more general problem stems from the apparent conflict between reports of the detection of massive galaxies at very high redshifts (e.g. Cimatti et al. 2004; Glazebrook et al. 2004; Labb\\'e et al. 2005; Daddi et al. 2005a,b) and indications for a fast decline in the in the comoving number density at $z>1$ (Franceschini et al. 1998; Fontana et al. 2004).  In summary, whereas early monolithic collapse and hierarchical models imply radically different histories for spheroids, neither the theoretical predictions nor the observational constraints for field galaxies have yet been sufficiently definitive for precise conclusions to be drawn in favour of one or the other scenario.  The most direct way of constraining the evolutionary history of galaxies and trying to resolve the discrepancies would be to derive the redshift-dependent luminosity and mass functions from deep unbiassed surveys.  This has been pursued by a number of teams, relying on either U-band optical selection as a probe of SFR density (Lilly et al. 1996; Madau et al. 1996) and UV selection by GALEX (Schiminovich et al. 2005), or observations in the K-band (Cowie et al. 1996; Cimatti et al. 2002; Dickinson et al. 2003; Franx et al. 2003; Drory et al. 2004; Fontana et al. 2004; Bundy, Ellis \\& Conselice 2005).  We contribute to this effort by exploiting in this paper very deep public imaging by the IRAC photometric camera on the Spitzer Observatory to select a most unbiased sample of  high-z ($z\\leq 2$) near-IR galaxies. We use for this the IRAC Channel-1 3.6$\\mu$m data over 160 arcmin$^2$ in the Chandra Deep Field South (CDFS) taken within the GOODS project (Dickinson et al. 2004). The other IRAC imaging data in the field are either redundant (Channel-2 at 4.5$\\mu$m, too close to channel-1 and somewhat less sensitive), or include non-stellar contributions by the galaxy ISM (the longer wavelength Channels-3 and 4) which would far complicate the interpretation.  Near-IR surveys are best suited for the study of faint high-redshift galaxy populations, for various reasons.  Compared to UV-optical selection, the observed fluxes are minimally affected by dust extinction. At the same time they are good indicators of the stellar mass content of galaxies (Dickinson et al. 2003; Berta et al. 2004), and  closer to provide a mass-selection tool. For typical spectra of evolved galaxies, the IRAC Channel-1 3.6$\\mu$m also benefits by a K-correction particularly favourable for the detection of high-redshift galaxies, particularly if we consider that the $H^-$ opacity minimum (corresponding to a typical peak in galaxy's SEDs) at $\\lambda\\simeq 1.7\\ \\mu$m in stellar athmospheres (Simpson \\& Eisenhardt 1999 ) falls within the waveband of Channel-1 for z=1 objects.   In addition to the Spitzer observations, the GOODS project has provided the community with an unprecedented amount of high quality optical and near-IR data in CDFS, particularly the very deep 4-band ACS imaging, allowing the most accurate morphological analysis currently possible over an appreciable area. Building on our previous experience of faint galaxy imaging and statistical modelling (Rodighiero et al. 2001; Cassata et al. 2005), on our tools for spectro-photometric analysis (Poggianti, Bressan \\& Franceschini 2001; Berta et al. 2004), and expertises for reduction of deep IR imaging data from space (e.g. Rodighiero et al. 2004 for ISO data; Lonsdale et al. 2004, Hatziminaoglou et al. 2005, Rodighiero et al. 2005 for Spitzer data), we illustrate in this paper the power of combining such multi-wavelength information in the analysis of the evolutionary mass and luminosity functions of faint galaxies.  Since the spectroscopic follow-up is currently only partial in the field, and to avoid as far as possible confusion problems in the IRAC data, we limit our analysis to moderate depths. In spite of this, the constraints on the history of massive galaxy evolution are already relevant. Pushing the analysis to the IRAC sensitivity limits will allow to extend our conclusions further down in mass/luminosity and up in $z$.  Though substantially larger than HDF's, our survey field is still of moderate size. Eventually, a complete understanding of the influence of large-scale structure on the evolutionary history and to reduce the effects of cosmic variance will need new-generation datasets on large areas (e.g. COSMOS, Scoville et al. 2004).  The paper is structured as follows. Section \\ref{data} describes the optical and IR (Spitzer) data, the near-IR data and ACS/HST imaging, the spectroscopic data used in our analysis, and our criteria for catalogue combination and merging. Section \\ref{morph} details on our quantitative morphological analysis of faint galaxies, and Section \\ref{photoz} summarizes our effort for the photometric redshift estimate. Our statistical analyses are reported in Section \\ref{stati}, while Section \\ref{mf} is dedicated to derive the mass function in stars of our galaxies. A comparison of observed and model number counts is discussed in Section \\ref{model}. Sections \\ref{discussion} and \\ref{conclusion} summarize our results and conclusions.  We adopt in the following a standard set of values for the cosmological parameters $\\Omega_M$=0.3 and $\\Omega_{\\Lambda}=0.7$, while for ease of comparison with other published results we express the dependence on the Hubble constant in terms of the parameter $h\\equiv H_0/100\\ Km/s/Mpc$ and provide the relevant scaling factors. ",
        "conclusions": "This paper is devoted to a systematic exploitation of public multi-wavelength data from the GOODS survey in the Chandra Deep Field South to derive observational constraints on the emergence of the Hubble galaxy morphological sequence through cosmic time. Critical data for this purpose are made available, in particular, by the very deep multicolor high-resolution imaging by HST/ACS and by the Spitzer Space Telescope deep photometric infrared imaging. We also make use of extensive optical spectroscopic observations by the ESO VLT/FORS2 and VIMOS spectrographs.  Our main selection for faint high-redshift galaxies is based on deep images by IRAC on Spitzer. We have selected from them a highly reliable IRAC 3.6$\\mu$m sample of 1478 galaxies with $S_{3.6}\\geq 10\\ \\mu$Jy for detailed statistical analyses and for the derivation of mass and luminosity functions in bins of redshift. We have also extended the morphologically-differentiated number counts down to a flux limit of $S_{3.6}= 1\\ \\mu$Jy. We have carefully analyzed and thoroughly tested these data for completeness and reliability, based on simulations.   Forty-seven percent of the sample objects have spectroscopic redshift from the VVDS, K20 and GOODS projects. For the remaining, we have used photometric redshifts from COMBO-17 for galaxies below $z\\sim 1$, while, for galaxies for which the COMBO-17 guess was above 1, we have re-estimated the photometric redshifts with $Hyperz$. Deep K-band VLT/ISAAC imaging in the field is also used to derive further complementary statistical constraints and to assist the source identification and SED analysis.  This very extensive dataset is then used to assess evolutionary effects in the galaxy stellar mass and luminosity functions, while luminosity/density evolution is further constrained with the number counts and redshift distributions. The estimate of galaxy stellar masses benefits in particular by the constraint set by the IRAC 3.6 $\\mu$m flux on the number of low-mass stars. The deep ACS imaging has allowed us to differentiate these evolutionary paths by morphological type, that our simulations show to be reliable at least up to $z\\sim 1.5$.  The main results of the paper are hereby summarized.  \\begin{itemize}   \\item We have derived luminosity functions at 3.6$\\mu$m for various galaxy populations as a function of redshift up to $z=1.4$. After careful calibrations of the $M/L$ ratio, based on a detailed spectral fitting analysis to the observed SED's for each sample galaxy, we have also obtained estimates of the evolutionary stellar mass functions. On one side, the 3.6$\\mu$m luminosity functions that we have derived show evidence for a positive, moderate luminosity evolution as a function of redshift (by $\\sim0.7$mag in the L-band from $z=0$ to 1.2 for the most massive galaxies, $M h^2 >10^{11}M_\\odot$, in agreement with Treu et al. 2005a), likely due to stellar ages in galaxies becoming younger at increasing $z$. On the other hand, the corresponding global mass function shows evidence for an exponential decrease in the comoving density of galaxies ($\\rho_\\ast\\propto exp(-[2+z]^4/141.6)$) at the corresponding redshifts.   \\item The galaxy number counts, z-distributions, the $\\left<V/V_{max}\\right>$ test, as well as our direct estimate of the stellar mass function above $M_\\ast h^2 =10^{10} M_\\odot$, provide consistent evidence for a progressive dearth (by a factor $\\sim2.5$ by $z=1.2$ for the stellar mass density, see Fig. \\ref{zint} and eq. \\ref{rho2}) of the spheroidal galaxy population to occur starting at low-$z$ and becoming quite significative at $z\\geq 0.7$, paralleled by an increase in luminosity (half a mag in L-band). Simple evolutionary models, fitting the fast convergence of the number counts and redshift distributions, and the evolutionary mass function, require the main episodes for spheroidal build-up (of either old or newly-formed stellar populations) to happen between $z\\sim 2$ and $z\\leq 1$ for such field population, on average.  \\item This decrease in comoving density of galaxies with redshift shows, however, a remarkable dependence on galaxy mass, being strong for moderate-mass, but almost absent until $z=1.4$ for high-mass galaxies, thus confirming previous evidence for a \"downsizing\" effect in galaxy formation (e.g. Cowie et al. 1996; Franceschini et al. 1998). By comparison with dynamical studies of the high-redshift spheroidal population (Treu et al. 2005a,b), it is concluded that both stellar mass and total \"dynamical\" mass are  driving parameters of this differential evolution. This evolutionary pattern may also help explaining some inconsistencies in the evolution of galaxies at high redshifts previously reported by different teams. Our results appear consistent with recent reports by independent teams and selection functions (Fontana et al. 2004; Bundy et al. 2005).   \\item As for the complementary class of actively star-forming (irregular/merger) galaxies, deep Spitzer/IRAC 3.6$\\mu$m and K-band observations show them to evolve towards moderately higher luminosities and number densities up to $z\\sim 1$ to 2, while normal spirals (those with asymmetry indices $A<0.4$)  show similar, though slower, convergence at $z>1$ to that of spheroids.  \\item Our favored interpretation of the estimated mass functions and evolutionary trends for the two broad galaxy categories is that of a progressive morphological transformation (due to gas exhaustion and, likely, merging) from the star-forming to the passively evolving phase starting at $z\\geq 2$ and keeping on down to $z\\sim 0.7$. The rate of this process appears to depend on galaxy mass, being already largely concluded by $z\\sim 1.4$ for the most massive systems.  \\item We finally discuss how well this evidence for a differential rate of galaxy build up with galactic mass compares with estimates of the SFR history based on deep far-IR surveys (e.g. Perez-Gonzales et al. 2005). A match between the two complementary views, of the history of SFR by the best star-formation tracer (the bolometric flux) on one side, and the rate of stellar mass build up traced by the near-IR emission on the other, would be achieved just by assuming that the progenitors of the most massive galaxies are the most (bolometrically) luminous sources at high-$z$.  Ample evidence is accumulating in favor of the latter.  \\end{itemize}  If the evolution pattern for galaxies to $z\\sim 1.4$ is now close to be understood, the knowledge of what exactly happened in the critical higher redshift era is still limited by very poor statistics in the number of detected sources and lack of spectroscopic follow-up. How is the mass function behaving in detail at such high redshifts? Does the \"downsizing\" trend continue there, as it might seem natural to expect? In principle, the sensitivity of Spitzer/IRAC would allow accurate stellar mass determinations at these high-$z$, but quite further substantial effort with powerful spectrographs is needed before we get credible answers. It is encouraging that much along this line has already been undertaken (among others, by GDDS, Abraham et al. 2004, Juneau et al. 2005; FIRES, Franx et al. 2003; GMASS, Cimatti et al., in progress; COSMOS \\& z-COSMOS, Scoville et al., Lilly et al., in progress).    \\vspace{0.75cm} \\par\\noindent {\\bf ACKNOWLEDGMENTS} \\par \\noindent This work is based on observations made with the {\\it Spitzer Space Telescope}, which is operated by the Jet Propulsion Laboratory, California Institute of Technology under NASA contract 1407. Support for this work, part of the Spitzer Space Telescope Legacy Science Program, was provided by NASA through an award issued by the Jet Propulsion Laboratory, California Institute of Technology under NASA contract 1407.  ACS was developed under NASA contract NAS 5-32865, and this research has been supported by NASA grant NAG5-7697. We are grateful for an equipment grant from  Sun Microsystems, Inc. The Space Telescope Science Institute is operated by AURA Inc., under NASA contract NAS5-26555.  Many of the observations bringing to these results have been carried out using the Very Large Telescope at the ESO Paranal Observatory. %under Program ID(s): LP168.A-0485  This work makes use the GalICS/MoMaF Database of Galaxies (http://galics.iap.fr). We thank L. Silva for making available to us her code's results in tabular form and she, Alvio Renzini and Andrea Cimatti for useful comments. We warmly thank Laurence Tresse, the referee, for a careful reading of the paper and numerous useful comments."
    },
    "0601/astro-ph0601244_arXiv.txt": {
        "abstract": "Metallic gas detected in the $\\beta$~Pictoris circumstellar debris disk raises many questions. The origin of this gas is unclear and its very presence is difficult to explain: many constituents of the gas are expected to be radiatively accelerated outward, yet their motion appears to be consistent with Keplerian rotation out to at least $300$~AU. Hydrogen has previously been hypothesized to exist in the disk, acting as a braking agent, but the amount required to brake individual elements conflicted with observed upper limits.  To resolve this discrepancy, we search for alternative braking mechanisms for the metallic gas. We find that all species affected by radiation force are heavily ionized. Frequent Coulomb collisions couple the ions into a single fluid, reducing the radiation force on species feeling the strongest acceleration. For a gas of solar composition, the resulting total radiation force still exceeds gravity, while a gas of enhanced carbon abundance could be self-braking.  We also explore two other braking agents: collisions with dust grains and neutral gas. Grains surrounding $\\beta$\\,Pic are photoelectrically charged to a positive electrostatic potential. If a significant fraction of the dust grains are carbonaceous ($10\\%$ in the midplane and larger at higher altitudes), ions can be slowed down to satisfy the observed velocity constraints. In this case, both the gas kinematics and spatial distributions are expected to coincide with those of small grains, the latter being indeed observed. For neutral gas to brake the coupled ion fluid, we find the minimum required mass to be $\\approx 0.03$\\,M$_{\\earth}$, consistent with observed upper limits of the hydrogen column density, and substantially reduced relative to previous estimates.  Our results favor a scenario in which metallic gas is generated by grain evaporation in the disk, perhaps during grain-grain collisions. We exclude a primordial origin for the gas, but cannot rule out the possibility of its production by falling evaporating bodies near the star. We discuss the implications of this work for observations of gas in other debris disks. ",
        "introduction": "\\label{s:intro}  The evolution of circumstellar disks is intimately connected to the process of star and planet formation. In the standard isolated low-mass star formation scenario (e.g., \\citealt*{shuadamslizano}), disks begin their existence as \\emph{protostellar} accretion disks, later turning into \\emph{protoplanetary} disks, once the newborn star reaches the T~Tauri phase. The late evolutionary stages of circumstellar disks are not well understood. Once the star has reached the main sequence and the planet formation process is complete, the disk is thought to have been depleted of its primordial gas, keeping small amounts generated by either stellar winds or evaporation of solid bodies (e.g., \\citealt{lagrange2000}). In the few cases where late disks are observable, lifetimes of small dust grains ($\\lesssim$ 1\\,$\\mu$m) are shorter than the estimated ages of those systems, implying a dust replenishment mechanism \\citep{backman93}. These kind of late circumstellar disks are usually referred to as \\emph{debris disks}. The gas lifetime in those systems, currently unknown, determines the maximum period over which protoplanetary cores can accrete gas and form gaseous planets (e.g., \\citealt{bodenheimer_lin}), and is related the formation of terrestrial planets and cores of giant planets by the increased efficiency of planetesimal accretion due to the presence of gas \\citep{rafikov04}.  So far, the best studied debris disk is that of $\\beta$~Pictoris, a nearby (19.3 $\\pm$ 0.2 pc, \\citealt{crifo97}), A5V star. The dust component of the disk has been extensively studied since it revealed itself as infrared excess in the stellar spectrum \\citep{aumann85}. The favorable edge-on orientation has allowed direct imaging in thermal emission \\citep{lagage94}, as well as in scattered light (\\citealt{smith84}, \\citealt{heap00}), indicating a disk size $\\sim$1\\,000\\,AU. The gas component has been more elusive. \\citet{freudling95} set a $3\\sigma$ upper limit on the column density of H\\,I, $N_\\mathrm{H\\,I}\\leq 10^{19}$\\,g\\,cm$^{-2}$, by non-detection of the 21\\,cm line, while \\citet{thi01} and \\citet{lecavelier01} searched for signatures of H$_\\mathrm{2}$ in the IR and UV, respectively, obtaining discrepant results: 50\\,M$_\\earth$ in the first case (corresponding to N$_{\\mathrm{H}_2}\\sim 3\\times 10^{21}$\\,cm$^{-2}$), and an upper limit N$_{\\mathrm{H}_2}\\leq 10^{18}$\\,cm$^{-2}$ ($\\sim$0.1\\,M$_\\earth$) in the second. Metallic elements have been detected through stable and variable circumstellar lines \\citep{hobbs85, ferlet87}. The variable, redshifted lines have been attributed to falling evaporating bodies (FEBs) in highly eccentric orbits \\citep{ferlet87}, whereas the stable component has no well established explanation for its origin. \\citet{lagrange98} proposed three possible gas formation scenarios: stellar wind, star-grazing comet evaporation, and dust grain evaporation near the star. Recent improvements in the gas characterization include spatially resolved emission lines from stable metallic species, which revealed them to be consistent with Keplerian rotation \\citep{olofsson01}, and spatially extended out to at least $300$\\,AU from the star \\citep{brandeker04}.  It was realized early on (e.g., \\citealt{beust89}) that certain elements in the gas disk of $\\beta$\\,Pic should be subject to a very strong radiation force from the star, overcoming gravitational attraction. However, besides Keplerian rotation, stable lines do not show any substantial radial velocities relative to the star (e.g., \\citealt{lagrange98}, \\citealt{brandeker04}). This implies the presence of a braking mechanism acting in the disk.  The first model of how gas might be braked and produce the stable absorption lines was presented by \\citet{lagrange98}. They proposed a scenario in which gas is injected into the disk due to FEBs. Since neutral hydrogen is not affected by the stellar radiation force, it would accumulate in the region of gas generation, in the form of an annulus located a few AU from the star. The elements affected by radiation force would then be slowed down temporarily in this annular region, producing the observed stable absorption features. After passing this hydrogen-rich region, elements would then be accelerated again, leaving the system at high velocities. This model reproduces the stable absorption lines, and is consistent with the FEBs scenario for the variable lines. However, the fact that elements affected by radiation force reach high velocities after passage by the annulus, which extends only out to a few AU from the star, is incompatible with the emission from $\\sim$300\\,AU. \\citet{brandeker04} calculated the amount of material necessary to slow down the most strongly accelerated particles to within observational constraints by means of neutral-neutral collisions, finding that $\\sim$50\\,M$_\\earth$ of hydrogen are needed. According to \\citet{thebault05}, there could be no more than 0.4\\,M$_\\earth$ of gas in the disk, as it would affect the dynamics of the dust to the degree that dust production models become incompatible with observations.  As an alternative braking mechanism, \\citet{alexis_thesis} explored the slowing down of ionized species by a large-scale magnetic field, concluding that only a toroidal field could remove the radial velocity components of ions. \\citet{alexis_thesis} showed that unless the field is stronger than $\\sim$1\\,mG at 100\\,AU, this configuration would be unstable to currents generated by radiation force. He also raised the question of the origin of such an ad-hoc field geometry.  In this paper we address the braking of the stable metallic gas in the entire $\\beta$\\,Pic disk. Given the spectral type of the star and the fact that the disk is optically thin, we expect several gaseous species to be ionized. This prompts us to explore the role of Coulomb collisions among ionized gas particles, which may be an important braking agent due to the long range of electromagnetic interactions. We proceed then to explore the interaction of this ionized gas with dust grains, which are photoelectrically charged. Finally, we extend the study of \\citet{lagrange98} to explore whether ion-neutral collisions are able to brake particles affected by radiation force far out in the disk, and the amount of neutral gas required to satisfy the observed velocity constraints.  The paper is organized as follows: in \\S\\ref{s:stage} we explore the physical conditions of the gas, and the constraints on the possible braking mechanisms that this imposes. In \\S\\ref{s:brake} we study ion-ion collisions, ion-grain collisions, and revisit ion-neutral interactions. In \\S\\ref{s:discuss} we discuss the implications of our results for gas generation and extend it to similar systems. Our conclusions follow in \\S\\ref{s:conclusions}. ",
        "conclusions": "\\label{s:conclusions}  Motivated by the apparent contradiction between the strong radiation force acting on the gas in the $\\beta$~Pictoris disk and spatially resolved observations showing it to be consistent with Keplerian rotation, we have explored different braking mechanisms that are likely to operate in the disk. We have found that: \\begin{itemize} \\item[a)] All species affected strongly by radiation force are heavily ionized. The velocities that the short-lived neutral particles achieve as a result of radiative acceleration, before becoming ionized, are consistent with observational constraints. \\item[b)] Ions are dynamically coupled due to their high Coulomb collision frequency. They feel a single effective radiation force coefficient $\\beta_\\mathrm{eff}$. For solar composition, this coefficient is $\\sim$5, in which case ions cannot brake by themselves. If carbon is over-abundant by a factor $\\gtrsim$10 on the other hand, the ion fluid may indeed be self-braking. \\item[c)] If a significant fraction of the dust in the disk is carbonaceous, the ion ensemble can be slowed down to drift velocities below the local sound speed by photoelectrically charged dust grains. This fraction is as low as $10\\%$ near the midplane and rises with altitude above the midplane. \\item[d)] Ions can also be slowed down by collisions with neutral gas (e.g., H$_2$, H$_2$O). The amount of neutral material required to satisfy observational constraints on drift velocities is $<$1\\,M$_\\earth$, substantially less than previous estimates. The UV absorption upper limit on $H_2$ is $\\sim 0.1 M_\\earth$ \\citep{lecavelier01}. \\item[e)] Reasonable assumptions on gas drift velocities imply gas lifetimes of order $10^4-10^5$\\,yr, much lower than the age of the system and which require a replenishment mechanism, in analogy with the dust. \\end{itemize}  The required gas production rate $10^{-13}$\\,M$_\\sun$\\,yr$^{-1}$ is comparable to the FEB scenario prediction (e.g., \\citealt{lagrange88}), although we argue that gas is more likely to be produced by local dust grain evaporation. A primordial origin of the gas is ruled out, since for having a lifetime of $\\sim$10$^7$\\,yr it would require a substantial quantity of neutral material, which would have already been detected.  Further gas observations in this and other systems would be able to test several predictions arising from this work, such as the absence of P and Be in $\\beta$\\,Pic, and the relative importance of ion-grain and ion-neutral collisions for slowing down ions. A more detailed study of the dust composition would shed light on the feasibility of the dust braking scenario."
    },
    "0601/astro-ph0601558_arXiv.txt": {
        "abstract": "We present an analysis of the mid-infrared (MIR) and optical properties of type 1 (broad-line) quasars detected by the {\\em Spitzer Space Telescope}.  The MIR color-redshift relation is characterized to $z\\sim3$, with predictions to $z=7$.  We demonstrate how combining MIR and optical colors can yield even more efficient selection of active galactic nuclei (AGN) than MIR or optical colors alone.  Composite spectral energy distributions (SEDs) are constructed for 259 quasars with both Sloan Digital Sky Survey and {\\em Spitzer} photometry, supplemented by near-IR, {\\em GALEX}, VLA and {\\em ROSAT} data where available.  We discuss how the spectral diversity of quasars influences the determination of bolometric luminosities and accretion rates; assuming the mean SED can lead to errors as large as a factor of 2 for individual quasars.  Finally, we show that careful consideration of the shape of the mean quasar SED and its redshift dependence leads to a lower estimate of the fraction of reddened/obscured AGNs missed by optical surveys as compared to estimates derived from a single mean MIR to optical flux ratio. ",
        "introduction": "Access to the mid-infrared (MIR) region opens up new realms for quasar science as we are able to study large numbers of objects with high signal-to-noise ratio data in this bolometrically important band for the first time.  At least four distinct energy generation mechanisms are at work in active galactic nuclei (AGN) from jets in the radio, dust in the IR, accretion disks in the optical-UV-soft-Xray, and Compton upscattering in the hard-X-ray.  All of these spectral regions need to be sampled with high precision if we are to understand the physical processes governing AGN emission.  The {\\em Spitzer Space Telescope} \\markcite{wrl+04}({Werner} {et~al.} 2004) allows the first robust glimpse of the physics of the putative dusty torus in AGNs out to $z\\sim2$--3.  MIR photometry from {\\em Spitzer} has provided a better census of active nuclei in galaxies than has been previously possible \\markcite{lss+04}(e.g., {Lacy} {et~al.} 2004a).  Optical surveys are biased against heavily reddened and obscured objects, and even X-ray surveys may fail to uncover Compton-thick sources.  Thus the MIR presents an attractive window for determining the black hole accretion history of the Universe.  To that end, {\\em Spitzer} will be of considerable utility in helping to decipher the nature of the $M_{\\rm BH}-\\sigma$ relation \\markcite{tgb+02}(e.g., {Tremaine} {et~al.} 2002), both in terms of the census of AGNs and in determining accurate bolometric luminosities (and, in turn, accretion rates).  This paper builds upon and extends the results from recent papers describing the {\\em Spitzer} MIR color distribution of AGNs. \\markcite{lss+04}{Lacy} {et~al.} (2004a) showed that MIR colors alone can be used to select AGNs with both high efficiency and completeness, including both dust reddened and optically obscured (type 2) AGNs that may otherwise be overlooked by optical selection techniques.  We will show that the addition of optical colors and morphology can be used to improve the MIR-only selection efficiency of type 1 quasars (including those that are moderately reddened).  \\markcite{seg+04}{Stern} {et~al.} (2005) also describe a MIR selection technique for AGN, making statistical arguments that the obscured AGN fraction may be as high as 76\\%.  We reconsider their arguement in light of the influence that the shape of the mean quasar spectral energy distribution (SED) has on determining the obscured quasar fraction.  Such considerations allow us to demonstrate that the true obscured AGN fraction must be lower than that determined by \\markcite{seg+04}{Stern} {et~al.} (2005).  Finally, \\markcite{hpp+04}{Hatziminaoglou} {et~al.} (2005) investigated the combined optical+MIR color distribution of quasars by combining data from the ELAIS-N1 field in the {\\em Spitzer} Wide-Area Infrared Extragalactic Survey (SWIRE; \\markcite{lsr+03}{Lonsdale} {et~al.} 2003) with data from the Sloan Digital Sky Survey (SDSS; \\markcite{yaa+00}{York} {et~al.} 2000).  Using the data from 35 SDSS quasars they determine the mean optical-MIR SED of type 1 quasars and investigate their mass and bolometric luminosity distribution.  We expand on these results by determining a number of different ``mean'' SEDs as a function of color and luminosity for 259 SDSS quasars in the {\\em Spitzer} Extragalactic First Look Survey\\footnote{http://ssc.spitzer.caltech.edu/fls/} (XFLS), SWIRE\\footnote{http://swire.ipac.caltech.edu/swire/} ELAIS-N1/N2, and SWIRE Lockman Hole areas.  We use these SEDs to demonstrate that the diversity of quasar SEDs must be considered when determining bolometric luminosities and accretion rates for individual quasars --- as was emphasized in the seminal SED work of \\markcite{ewm+94}{Elvis} {et~al.} (1994).  Section~2 reviews the data sets used in our analysis.  In \\S~3 we explore the MIR color-redshift relation and MIR-optical color-color space occupied by type 1 quasars.  In addition to showing these relations for the data, we also show the predicted relations derived from two quasar SEDs convolved with the SDSS and {\\em Spitzer} filters curves: one SED derived largely from broad-band photometry \\markcite{ewm+94}({Elvis} {et~al.} 1994), the other from a mean optical+IR spectral template \\markcite{ghw05}({Glikman et al.} 2005).  Section~4 presents a brief discussion of the determination of the type 1 to type 2 ratio of quasars. In \\S~5 we discuss the radio through X-ray SED of quasars and construct new MIR-optical templates from our sample. We present an overall mean SED along with mean SEDs for subsets of optically luminous/dim, MIR luminous/dim, and optically blue/red quasars in order to explore how different optical/MIR properties are related to the overall SED. Section~6 discusses the implications of our new SED templates on the determination of bolometric luminosities and accretion rates.  Our conclusions are presented in \\S~7.  Throughout this paper we will distinguish between normal type 1 quasars, dust reddened/extincted type 1 quasars, and type 2 quasars. By type 1 quasars we mean those quasars having broad lines and colors/spectral indices that are roughly consistent with a Gaussian spectral index distribution of $\\alpha_{\\nu}=-0.5\\pm0.3$ ($f_{\\nu}\\propto\\nu^{\\alpha}$).  Reddened type 1 quasars are those quasars that have broad lines but have spectral indices that are redder than about $\\alpha_{\\nu}=-1$.  Optical surveys can find such quasars up to $E(B-V)\\sim0.5$, but are increasingly incomplete above $E(B-V)\\sim0.1$ \\markcite{rhv+03}({Richards} {et~al.} 2003).  By type 2 quasars, we mean those that lack rest-frame optical/UV broad emission lines and with nuclei that are completely obscured in the optical such that the optical colors are consistent with the host galaxy.  Throughout this paper we use a $\\Lambda$-CDM cosmology with $H_{\\rm o}=70$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{\\Lambda}=0.7$, and $\\Omega_{\\rm m}=0.3$, consistent with the WMAP cosmology \\markcite{svp+03}({Spergel} {et~al.} 2003). ",
        "conclusions": "We have compiled a sample of 259 SDSS type 1 quasars with 4-band {\\em Spitzer}-IRAC detections.  These data are supplemented with multiwavelength data spanning the radio to X-ray where available.  We have shown that while MIR-only selection of AGNs is indeed very efficient, adding optical morphology and color information allows for even more efficient selection of type 1 quasars.  This multiwavelength data set was used to construct new mean SEDs of quasars as a function of color and luminosity, such as are needed to span the diversity of quasars in terms of computing accurate bolometric luminosities and accretion rates.  It was shown that computing a bolometric luminosity by assuming a single mean quasar SED can lead to errors as large as a factor of two, which translates directly to an error in the presumed accretion rate.  Finally, judicious use of the knowledge of the mean quasar SED makes it possible to estimate the fraction of AGNs that are obscured in the optical by considering the redshift dependence of the relative flux limits between the optical and MIR."
    },
    "0601/astro-ph0601591_arXiv.txt": {
        "abstract": "Recent calculations of atomic data for \\ion{Fe}{15} have been used to generate theoretical line ratios involving {\\em n} = 3--4 transitions in the soft X-ray spectral region ($\\sim$52--83 \\AA), for a wide range of electron temperatures and densities applicable to solar and stellar coronal plasmas. A comparison of these with solar flare observations from a rocket-borne spectrograph (XSST) reveals generally good agreement between theory and experiment. In particular, the 82.76 \\AA\\ emission line in the XSST spectrum is identified, for the first time to our knowledge in an astrophysical source, as the 3{\\em s}3{\\em d} $^{3}$D$_{3}$--3{\\em s}4{\\em p} $^{3}$P$_{2}$ transition of \\ion{Fe}{15}. Previous suggested identifications of the 53.11, 63.97 and 69.65 \\AA\\ features as the 3{\\em s}$^{2}$ $^{1}$S--3{\\em s}4{\\em p} $^{3}$P$_{1}$, 3{\\em p}$^{2}$ $^{1}$D--3{\\em s}4{\\em f} $^{1}$F and 3{\\em s}3{\\em p} $^{1}$P--3{\\em s}4{\\em s} $^{1}$S lines of \\ion{Fe}{15}, respectively, are confirmed. However the former is blended (at the $\\sim$ 50\\%\\ level) with the \\ion{S}{9} 2{\\em s}$^{2}$2{\\em p}$^{4}$ $^{3}$P$_{1}$--2{\\em p}$^{3}$3{\\em s} $^{3}$P$_{2}$ transition. Most of the \\ion{Fe}{15} transitions which are blended have had the species responsible clearly identified, although there remain a few instances where this has not been possible. The line ratio calculations are also compared with a co-added spectrum of Capella obtained with the {\\em Chandra} satellite, which is probably the highest signal-to-noise observation achieved for a stellar source in the $\\sim$25--175 \\AA\\ soft X-ray region. Good agreement is found between theory and experiment, indicating that the \\ion{Fe}{15} lines are reliably detected in {\\em Chandra} spectra, and hence may be employed as diagnostics to determine the temperature and/or density of the emitting plasma. However the line blending in the {\\em Chandra} data is such that individual emission lines are difficult to measure accurately, and fluxes may only be reliably determined via detailed profile fitting of the observations. The co-added Capella spectrum is made available to hopefully encourage further exploration of the soft X-ray region in astronomical sources. ",
        "introduction": "Emission lines arising from {\\em n} = 3--3 transitions in \\ion{Fe}{15} are widely detected in solar extreme ultraviolet (EUV) spectra covering the $\\sim$200--400 \\AA\\ wavelength interval (Dere 1978; Thomas \\& Neupert 1994).  The 3{\\em s}$^{2}$ $^{1}$S--3{\\em s}3{\\em p} $^{1}$P resonance line at 284 \\AA\\ is one of the most intense emission features in the solar EUV spectrum, and has also been widely observed in stellar coronal sources by the {\\em Extreme Ultraviolet Explorer} (EUVE) satellite (Laming \\& Drake 1999). Bely \\& Blaha (1968) first noted the potential of EUV lines of \\ion{Fe}{15} as diagnostics for the emitting plasma, and since then many authors have calculated emission line intensities for this ion and compared these with solar observations (see, for example, Kastner \\& Bhatia 2001 and references therein).  However \\ion{Fe}{15} also shows a rich emission line spectrum in the soft X-ray region, $\\sim$50--80 \\AA, arising from {\\em n} = 3--4 transitions.  By contrast to the EUV lines, there has been little work on the soft X-ray transitions in the astrophysical literature, probably as a result of the limited availability of high quality solar spectra for this wavelength region.  There were several rocket-borne flights in the 1960's which covered the soft X-ray region, but these detected few or no \\ion{Fe}{15} transitions (see, for example, Austin et al. 1966; Widing \\& Sandlin 1968; Behring, Cohen, \\& Feldman 1972; Malinovsky \\& Heroux 1973).  The most detailed measurements of the {\\em n} = 3--4 lines in \\ion{Fe}{15} to date were made by the XSST spectrograph during a rocket flight in 1982 (Acton et al. 1985).  However, even these observations have only been subject to a brief analysis by Bhatia, Mason, \\& Blancard (1997), as part of their work to produce a large atomic dataset for this ion. Finally, Campbell \\& Brickhouse (2001) presented preliminary results from an investigation of \\ion{Fe}{15} and \\ion{Fe}{16} lines in {\\it Chandra} X-ray spectra of the late-type coronally-active stars Capella (G1~III + G8~III) and Procyon (F5~IV).  In this paper we use the most recent atomic physics calculations for \\ion{Fe}{15} to generate soft X-ray line ratios for a wide range of electron temperatures and densities applicable to coronal plasmas.  We compare these in detail with the XSST observations of Action et al. (1985), to assess the usefulness of the \\ion{Fe}{15} transitions as plasma diagnostics. However, in addition we analyse a {\\em Chandra} spectrum of Capella constructed from multiple observations.  This co-added X-ray spectrum has one of the highest signal-to-noise ratios ever achieved for an astrophysical coronal plasma in the $\\sim$25--175 \\AA\\ soft X-ray region, hence allowing us to perform an investigation of the \\ion{Fe}{15} {\\em n} = 3--4 lines in an astronomical source other than the Sun.  The Capella spectrum is also made freely available, to encourage further research on the soft X-ray spectral region. For example, in conjunction with the published XSST solar line list, the Capella data should be useful for reliably identifying new emission features. In addition, there are potentially many useful diagnostic lines in the soft X-ray wavelength range, especially for solar flare-type plasmas (see, for example, Brown et al. 1986; Feldman et al. 1992), and hence further exploration and assessment of these is warranted. ",
        "conclusions": "In Tables 2 and 4 we list the observed \\ion{Fe}{15} emission line ratios from the XSST solar flare spectrum and the {\\em Chandra} Capella observations, respectively.  Also shown in the tables are the theoretical results from Figures 1--7 at the temperature of maximum fractional abundance in ionization equilibrium for \\ion{Fe}{15}, T$_{e}$ = 10$^{6.3}$ K (Mazzotta et al. 1998), and at electron densities derived for the solar and Capella observations from emission line ratios in \\ion{O}{7} (Brown et al. 1986; Phillips et al. 2001). The temperature-sensitive $G$-ratio for \\ion{O}{7} in the XSST and Capella data indicates that T$_{e}$ = 10$^{6.3}$ K in both features, similar to that of maximum fractional abundance for \\ion{Fe}{15}.  In the case of Capella, whose coronal emission measure distribution rises quite steeply to a peak at T$_{e}$ = 6$\\times$10$^6$~K (see, for example, Brickhouse et al. 2000), it is possible that this $G$-ratio slightly underestimates the mean temperature of line formation, as was found empirically by Testa, Drake \\& Peres (2004) for a sample of active stars.  However, the ion populations of both \\ion{O}{7} and \\ion{Fe}{15} peak at very similar temperatures, and hence the \\ion{O}{7} density should still accurately reflect that of the \\ion{Fe}{15} emitting region of the plasma.  We have also verified that the slightly higher temperatures of line formation that would result from the relatively steep Capella emission measure distribution do not significantly change the predicted \\ion{Fe}{15} line ratios from those listed in Table~4.  The error bars on the theoretical results are based on the estimated $\\pm$20\\%\\ accuracy of the line ratio calculations (see \\S\\ 2).  An inspection of Table 2 reveals very good agreement between the XSST observations and the theoretical predictions for the ratios  R$_{5}$ and R$_{7}$. This indicates that the 63.97 and 69.65 \\AA\\ lines are well detected in the XSST spectrum, and must be relatively free from blends. In particular, we confirm the identification by Bhatia et al. of the 63.97 \\AA\\ feature as the 3{\\em p}$^{2}$ $^{1}$D--3{\\em s}4{\\em f} $^{1}$F transition of \\ion{Fe}{15}, and not the \\ion{Al}{8} 2{\\em s}$^{2}$2{\\em p}$^{2}$ $^{3}$P$_{1,2}$--2{\\em s}2{\\em p}$^{2}$3{\\em p} $^{3}$D$_{2,3}$ lines as originally classified by Acton et al. (1985). These authors also identify the 69.65 \\AA\\ line as being partially due to \\ion{Si}{8} and \\ion{Fe}{14} transitions, with an unknown feature responsible for most of the line flux. However our result for R$_{7}$ clearly shows that the line is primarily the \\ion{Fe}{15} 3{\\em s}3{\\em p} $^{1}$P--3{\\em s}4{\\em s} $^{1}$S transition, as suggested by Bhatia et al. We have confirmed this by generating a synthetic spectrum using {\\sc chianti}, which indicates that the \\ion{Si}{8} 2{\\em s}$^{2}$2{\\em p}$^{3}$ $^{4}$S--2{\\em s}$^{2}$2{\\em p}$^{2}$3{\\em s} $^{4}$P$_{5/2}$ line only contributes about 15\\%\\ to the total 69.65 \\AA\\ flux, and that no \\ion{Fe}{14} transitions are present.  Similarly, the observed R$_{1}$, R$_{9}$, R$_{10}$, R$_{11}$ and R$_{12}$ ratios in Table 2 are only slightly larger than the theoretical values, indicating that the 52.91, 69.98, 70.05, 73.47 and 82.76 \\AA\\ lines must be mostly due to \\ion{Fe}{15}. Indeed, within the observational and theoretical errors in the line ratios, all the observed fluxes may arise from this ion. This is supported by the {\\sc chianti} synthetic spectrum, which predicts no significant blending species for the 52.91, 69.98 and 70.05 \\AA\\ features, and contributions of only about 10\\%\\ from \\ion{Ne}{8} 1{\\em s}$^{2}$2{\\em p} $^{2}$P$_{1/2}$--1{\\em s}$^{2}$4{\\em d} $^{2}$D$_{3/2}$ and 20\\%\\ from \\ion{Ne}{9} 1{\\em s}2{\\em p} $^{1}$P--1{\\em s}3{\\em s} $^{1}$S to the 73.47 and 82.76 \\AA\\ lines, respectively. Our classification of the 82.76 \\AA\\ feature as the \\ion{Fe}{15} 3{\\em s}3{\\em d} $^{3}$D$_{3}$--3{\\em s}4{\\em p} $^{3}$P$_{2}$ line is, to our knowledge, the first time this transition has been identified in an astrophysical source.  The good agreement between theory and observation for the above line ratios indicates that they should provide useful diagnostics for the \\ion{Fe}{15} emitting region of a plasma. In particular, R$_{7}$, R$_{9}$, R$_{10}$ and R$_{11}$ are all temperature and/or density sensitive, and hence with careful use should allow these parameters to be derived. For example, R$_{7}$ could be employed to evaluate T$_{e}$, as it does not vary with N$_{e}$ (Figure 1), and the latter could then be determined from R$_{10}$ or R$_{11}$ (Figures 6 and 7).  The observed values of R$_{2}$, R$_{3}$, R$_{4}$, R$_{6}$ and R$_{8}$ in the XSST spectrum are all significantly larger than theory, implying that the 53.11, 55.78, 56.17, 66.25 and 69.93 \\AA\\ lines of \\ion{Fe}{15} are blended to varying degrees. For the 53.11 \\AA\\ feature, the blending line contributes about 50\\%\\ of the measured flux, and the {\\sc chianti} synthetic spectrum indicates that this is most likely the \\ion{S}{9}  2{\\em s}$^{2}$2{\\em p}$^{4}$ $^{3}$P$_{1}$--2{\\em p}$^{3}$3{\\em s} $^{3}$P$_{2}$ transition. This is predicted to have an intensity of approximately 10\\%\\ that of the \\ion{Fe}{16} 3{\\em p} $^{2}$P$_{1/2}$--4{\\em d} $^{2}$D$_{3/2}$ line at 54.13 \\AA, i.e. about 16 photons cm$^{-2}$ s$^{-1}$ arcsec$^{-2}$. The measured flux of the 53.11 \\AA\\ feature is 31 photons cm$^{-2}$ s$^{-1}$ arcsec$^{-2}$, so \\ion{S}{9} should contribute 50\\%, in agreement with observation. We note that Acton et al. (1985) did not classify the 53.11 \\AA\\ feature, the \\ion{Fe}{15} identification being suggested by Bhatia et al. (1997). However we can now confirm that the line is the 3{\\em s}$^{2}$ $^{1}$S--3{\\em s}4{\\em p} $^{3}$P$_{1}$ transition of \\ion{Fe}{15}, although it is blended.  For the 55.78 \\AA\\ line, {\\sc chianti} indicates that the \\ion{Si}{10} 2{\\em s}2{\\em p}$^{2}$ $^{2}$S--2{\\em s}2{\\em p}3{\\em s} $^{2}$P$_{1/2,3/2}$ and \\ion{Si}{9} 2{\\em s}$^{2}$2{\\em p}$^{2}$ $^{3}$P$_{1}$--2{\\em s}$^{2}$2{\\em p}3{\\em d} $^{1}$D transitions should together have an intensity about 70\\%\\ that of the \\ion{Fe}{15} line. Our observation of R$_{3}$ implies that the blending species may be stronger, and contribute about 150\\%\\ of the \\ion{Fe}{15} flux. However when the uncertainties in the experimental and theoretical R$_{3}$ ratios are taken into account, the data are compatible with a 70\\%\\ contribution from \\ion{Si}{10} and \\ion{Si}{9}. Similarly, a comparison of theory and observation for R$_{6}$ indicates that \\ion{Fe}{15} only contributes about 10\\%\\ to the intensity of the 66.25 \\AA\\ feature. This is in agreement with {\\sc chianti}, which predicts that 90\\%\\ of the observed flux is due to the \\ion{Fe}{16} 3{\\em d} $^{2}$D$_{3/2}$--4{\\em f} $^{2}$F$_{5/2}$ line.  By contrast, it is not clear from {\\sc chianti} what the blending species are for the 56.17 and 69.93 \\AA\\ lines, with no transitions predicted to have intensities more than 10\\%\\ those of the relevant \\ion{Fe}{15} features. It is however possible that the blending may be due to strong lines observed by XSST in second- or third-order, which are appearing in first-order at these wavelengths. For example, in the case of the 56.17 \\AA\\ line, {\\sc chianti} indicates no strong transitions in second-order (i.e. around 28.09 \\AA), but the intense \\ion{Ca}{18} 1{\\em s}$^{2}$2{\\em s} $^{2}$S--1{\\em s}$^{2}$3{\\em p} $^{2}$P$_{1/2}$ transition is predicted in third-order (18.73 \\AA). Similarly, for the 69.93 \\AA\\ line there is a predicted third-order feature at 23.31 \\AA, namely the \\ion{Ca}{16} 2{\\em s}2{\\em p}$^{2}$ $^{2}$P$_{3/2}$--2{\\em s}2{\\em p}3{\\em d} $^{2}$D$_{5/2}$ transition. There are unfortunately no \\ion{Ca}{18} lines in the XSST spectrum for comparison purposes, but there is a \\ion{Ca}{16} transition at 21.45 \\AA\\ (2{\\em s}$^{2}$2{\\em p} $^{2}$P$_{1/2}$--2{\\em s}$^{2}$3{\\em d} $^{2}$D$_{3/2}$), with an intensity I = 37 photons cm$^{-2}$ s$^{-1}$ arcsec$^{-2}$. However the predicted {\\sc chianti} intensity ratio is I(23.31 \\AA)/I(21.45 \\AA) $<$ 10$^{-3}$, indicating that the \\ion{Ca}{16} 23.31 \\AA\\ transition cannot be responsible for the blend in the 69.93 \\AA\\ line. Clearly, the identification of the blends for the 56.17 and 69.93 \\AA\\ features will require further work.  In the case of the {\\em Chandra} observations of Capella, agreement between theory and experiment in Table 4 is very good for the majority of the line ratios. Where there are discrepancies, such as for R$_{8}$, these tend to mirror those found for the XSST spectrum, with the observed value being somewhat larger than theory, indicative of some blending.  However, in most instances the observed line ratios actually show smaller discrepancies with theory than do the XSST results. To a certain extent, such agreement may be judged to be unsurprising, as in our analysis of the Capella observations we added line components in an ad hoc fashion until the profiles of blending features could be empirically matched (see \\S\\ 3.2). Nevertheless, our results indicate that profile fitting of the {\\em Chandra} data, rather than simple measurements of emission line intensities, does allow the reliable detection of \\ion{Fe}{15} soft X-ray features in the 50--75 \\AA\\ wavelength range. Indeed, our work represents (to our knowledge) the first time that the \\ion{Fe}{15} {\\em n} = 3--4 lines have been identified in an astrophysical source other than the Sun. Hopefully, our results will encourage the future use of these lines as plasma diagnostics for this spectral region.  In summary, we have performed a detailed analysis of \\ion{Fe}{15} soft X-ray lines in the $\\sim$52--83 \\AA\\ wavelength region, observed in the solar spectrum by the XSST spectrograph, and confirmed identifications for several emission features and also detected a new transition. Additionally, the \\ion{Fe}{15} lines have been measured in {\\em Chandra} observations of Capella, and appear to provide potentially useful plasma diagnostics. We therefore hope that our work will stimulate renewed interest in this relatively unexplored spectral region, in particular through our provision of the co-added {\\em Chandra} spectrum for Capella, which are probably the highest signal-to-noise data available. Many dozens of lines in the solar soft X-ray spectrum are without identifications, and the Capella spectrum (in combination with the XSST solar line list) should provide an aid to the identification of soft X-ray transitions and the investigation of their usefulness as plasma diagnostics. This will be important for future space missions, and in particular the EUV Variability Experiment (EVE), which is part of the {\\em Solar Dynamics Observatory} (SDO), due for launch in 2008 \\footnote{See http://lasp.colorado.edu/eve/}. One of the EVE instruments is the Multiple EUV Grating Spectrograph (MEGS), which will observe the 50--1050 \\AA\\ region at a spectral resolution of 1 \\AA. Clearly, the lines in this wavelength range will need to be identified and understood in advance of MEGS observations, as this instrument will not be able to resolve them individually."
    },
    "0601/astro-ph0601072_arXiv.txt": {
        "abstract": "We use numerical simulations to study the effect of nonlinear MHD waves in a stratified, self-gravitating molecular cloud that is bounded by a hot and tenuous external medium. In a previous paper, we had shown the details of a standard model and studied the effect of varying the dimensionless amplitude $\\atil$ of sinusoidal driving. In this paper, we present the results of varying two other important free parameters: $\\beta_0$, the initial ratio of gas to magnetic pressure at the cloud midplane, and $\\nutil$, the dimensionless frequency of driving. Furthermore, we present the case of a temporally random driving force. Our results demonstrate that a very important consideration for the actual level of turbulent support against gravity is the ratio of driving wavelength $\\lam$ to the the size of the initial non-turbulent cloud; maximum cloud expansion is achieved when this ratio is close to unity. All of our models yield the following basic results: (1) the cloud is lifted up by the pressure of nonlinear MHD waves and reaches a steady-state characterized by oscillations about a new time-averaged equilibrium state; (2) after turbulent driving is discontinued, the turbulent energy dissipates within a few sound crossing times of the expanded cloud; (3) the line-width--size relation is obtained by an ensemble of clouds with different free parameters and thereby differing time-averaged self-gravitational equilibrium states. The best consistency with the observational correlation of magnetic field strength, turbulent line width, and density is achieved by cloud models with $\\beta_0 \\approx 1$. We also calculate the spatial power spectra of the turbulent clouds, and show that significant power is developed on scales larger than the scale length $H_0$ of the initial cloud, even if the input wavelength of turbulence $\\lam \\approx H_0$, The cloud stratification and resulting increase of \\Alf speed toward the cloud edge allows for a transfer of energy to wavelengths significantly larger than $\\lam$. This explains why the relevant time scale for turbulent dissipation is the crossing time over the cloud scale rather than the crossing time over the driving scale. ",
        "introduction": "Interstellar molecular clouds have nonthermal line widths which are typically supersonic and increase with the scale of the cloud size that is measured (e.g., Larson 1981; Myers 1983; Solomon et al. 1987; Brunt \\& Heyer 2002). For a cloud scale $\\approx 1$ pc, the velocity dispersion $\\sigma \\approx 1 \\kms$, which is several times greater than the isothermal sound speed $\\cs = 0.19 \\kms$ for a typical cloud temperature $T = 10 \\K$ (Goldsmith \\& Langer 1978). When detectable through the Zeeman effect, magnetic field strengths are such that their influence is comparable to that of gravity (Crutcher 1999) and velocity dispersions are correlated with the mean \\Alf speeds $\\bar{V}_A$ such that $\\sigma/\\bar{V}_A \\approx 0.5$ (Crutcher 1999; Basu 2000). Furthermore, maps of polarized emission from dust imply that the magnetic field morphology is relatively well-ordered and therefore not dominated by turbulence (e.g., Schleuning 1998; Matthews \\& Wilson 2000; Houde et al. 2004). It seems possible that turbulence within molecular clouds comprises a spectrum of nonlinear magnetohydrodynamic (MHD) waves generated from various localized sources and from self-gravitational motions. In contrast to the large-scale structuring by turbulence and magnetic fields, we note that the embedded dense cores in regions of low-mass star formation have characteristically subsonic or transonic turbulence; evidence for gravitational collapse in cores also comes in the form of subsonic infall motions where detectable (see discussion and references in Myers 2005).  Although not possessing internal motions as highly supersonic as some larger, more distant molecular clouds, the well-resolved Taurus molecular cloud complex (distance $\\approx 140$ pc) reveals a hierarchy of density and size scales that is characteristic of star-forming regions (see Ungerechts \\& Thaddeus 1987; Mizuno et al. 1995; Onishi et al. 1998; Onishi et al. 2002). The cloud has a common (turbulent) envelope of molecular column density $N \\approx 1 \\times 10^{21} \\cms - 2 \\times 10^{21} \\cms$ (corresponding to visual extinction $A_V \\approx 1 - 2$), embedded parsec-scale dark clouds with $A_V \\gtrsim 3$ and velocity dispersion $\\sigma \\approx 0.6 \\kms \\approx 3 \\, \\cs$, which in turn contain dense regions with an average $A_V \\approx 5$ that are forming small clusters of stars. The typical separation of young stellar objects within the dense regions may be understood by a simple Jeans-like fragmentation process for an isothermal dense sheet or filament (Hartmann 2002). Such an approach ignores the effect of turbulence, which is not measured to be large in those regions anyway. However, the subsonic infall motions imply that a fragmentation process mediated by magnetic fields and ion-neutral friction may be more appropriate (Basu \\& Ciolek 2004).  In this paper, we extend a previous set of models (see Kudoh \\& Basu 2003, hereafter Paper I) which approximate conditions within the molecular cloud envelopes and outer regions of dark clouds. These regions are expected to be sufficiently ionized by the far-ultraviolet background starlight that the ideal MHD limit of flux-freezing is applicable (McKee 1989). Some key questions about cloud envelopes are as follows. What is the rate of dissipation of the waves and what is the primary dissipation mechanism? Is continual driving of the turbulence necessary to explain their ubiquitous presence in all but the densest regions? Furthermore, observed properties of the turbulence such as the line-width--size relation (e.g., Solomon et al. 1987) and the correlation between magnetic field strength, line width, and density (Myers \\& Goodman 1988; Basu 2000) deserve an explanation.  Since the late 1990's, a set of three-dimensional models of MHD turbulence have been developed which use a {\\em local} approach, i.e., use periodic boundary conditions (e.g., Stone, Gammie, \\& Ostriker 1998; Mac Low et al. 1998; Mac Low 1999; Padoan \\& Nordlund 1999; Ostriker, Stone, \\& Gammie 2001). This means that the modeled regions represent a very small region deep within a molecular cloud. However, they are driven with highly supersonic motions characteristic of large-scale molecular cloud structure. Although these models are local, the turbulent driving is global, in that large amplitude momentum fluctuations are input at every grid point. The fluctuations are drawn from a Gaussian random field with a specified root mean squared amplitude and power spectrum. Although this method is convenient, a more likely scenario for molecular cloud turbulence is that it is driven from local sources and propagates to fill the clouds. A main result of the periodic-box models is that the decay time of the MHD turbulence is comparable to the crossing time over the driving scale of the turbulence. For instance, Stone et al. (1998) find a turbulent dissipation rate $\\Gamma_{\\rm turb} \\approx \\rho_0 \\sigma^3/\\lam$, where $\\rho_0$ is the fixed mean density in the periodic simulation box, $\\sigma$ is the one-dimensional velocity dispersion, and $\\lam$ is the driving scale of turbulence. A direct application of this result to large molecular clouds has been questioned by Basu \\& Murali (2001), who point out that dissipation proportional to the crossing time across internal driving scales would lead to a CO emission luminosity $L_{\\rm CO}$ far exceeding that observed for clouds whose size is much larger than the driving scale. Alternatively, they point out that the observed scaling of $L_{\\rm CO}$ with cloud mass can be explained if the relevant length scale for dissipation is the largest scale $L$ of any cloud; this also yields the correct order of magnitude for $L_{\\rm CO}$ itself. In the context of a periodic box model, this means that the driving would have to occur only on the largest scales of each cloud. McKee (1999) also questions extremely fast dissipation, on the grounds that: (1) turbulence should not be so ubiquitous (and present for example in clouds with or without embedded OB associations, a possible major source of stellar-driven turbulence) if it is decaying so rapidly; (2) dissipation times of at least a few cloud crossing times (as opposed to a free-fall time) may explain the relatively low Galactic star formation rate; and (3) current simulations cannot resolve the turbulence within the dense clumps that quickly develop. Sugimoto, Hanawa, \\& Fukuda (2004) have supported the last point with second-order accurate three-dimensional numerical studies of \\Alf wave propagation, finding that $\\simeq 30$ grid points per wavelength are necessary to reduce dissipation to 1\\% per wavelength propagated. Such a high resolution remains prohibitive for three-dimensional simulations that seek to model large-scale cloud dynamics and dense core formation. Cho \\& Lazarian (2003) have also questioned a common interpretation of three-dimensional periodic box models, i.e., that the fast dissipation is due to the coupling of \\Alf modes to the more dissipative fast MHD and slow MHD modes (e.g., Stone et al. 1998). Through modal analysis of their three-dimensional simulations, they find that the coupling of modes is rather weak in the inertial range of turbulence and that dissipation of \\Alf waves is dominated by a nonlinear cascade. Furthermore, Cho, Lazarian, \\& Vishniac (2002) have found that in the case of incompressible MHD turbulence, a directionality in the flux of energy (as may be expected with localized driving sources) leads to a significantly lower dissipation rate.   In addition to the detailed analysis of individual wave modes as described above, new approaches to the numerical study of MHD turbulence include extremely high resolution models of the {\\em global} structure of molecular clouds (Kudoh \\& Basu 2003; Folini, Heyvaerts, \\& Walder 2004). These one-dimensional models employ a numerical resolution that remains beyond the reach of three-dimensional simulations. Kudoh \\& Basu (2003, hereafter Paper I) have modeled the global structure of a one-dimensional cold cloud that is embedded in a hotter external medium. The main advantage over a local (periodic box) approach is the inclusion of the effects of cloud stratification and a cloud edge. Furthermore, the turbulent driving is localized and the propagation of disturbances within the cloud as well as their leakage from the cloud can be studied. Our model in Paper I had 50 numerical cells within the scale length $H_0$ of the initial state, and 70 cells within the wavelength $\\lam$ associated with turbulent driving. Results included the generation of significant density structure and shock fronts within the cloud; however, most of the wave energy remained in transverse, rather than longitudinal motions (in qualitative agreement with the results of Cho \\& Lazarian 2003) and the time-averaged density profile of the cloud (which may be similar to an observed spatial average of many layers along the line-of-sight) was relatively smooth. The localized driving led to turbulence which quickly filled the cloud, causing cloud expansion and large-scale oscillations in the outer regions; all this is aided by a clear transfer of energy to the largest scales of the expanded cloud, due to the stratification and consequent increase of \\Alf speed toward the cloud edge. The work of Folini et al. (2004), who model a very similar system with similar turbulent driving, confirm the existence of highly disturbed density structure, oscillatory motions, and a dominance of wave energy in transverse modes. They focus most extensively on the density structuring and demonstrate that progressively finer scale structuring and enhanced cloud support is revealed by higher resolution simulations; for typical dimensional values of physical parameters, their numerical resolution is 0.001 pc, as in our models in Paper I. In Paper I, we also considered many global properties of the clouds. We found that random motions were greatest in the outer, low-density regions of the clouds in a way that the velocity dispersion $\\sigma$ and cloud scale $L$ satisfied the observed line-width--size relation. Our model also showed that even when turbulence decays away due to lack of driving, the large-scale oscillations are the longest lived component. These results can help to explain the observations of apparent oscillations of B68 (Lada et al. 2003), a dark cloud in which nonthermal motions are small. We also showed that the overall dissipation of turbulent energy (after driving is discontinued) is proportional to the crossing time across the expanded cloud, rather than the crossing time for the driving wavelength $\\lam$. This result points to a way out of the problems with internal driving pointed out by Basu \\& Murali (2001). Stratification leads to turbulence at the largest scales, so that the cloud scale $L$ is more relevant to turbulent dissipation than the internal driving scale $\\lam$, as needed on empirical grounds. We quantify the effect of turbulence on large scales further in this paper by considering the power spectrum of fluctuations in the turbulent clouds.   This paper expands upon the study of Paper I and summarizes the results of a parameter study. The numerical model is presented in \\S\\ 2. The results of the simulations are described in \\S\\ 3. A summary and discussion are presented in \\S\\ 4. ",
        "conclusions": "We have performed 1.5-dimensional numerical simulations of nonlinear MHD waves in a gravitationally stratified molecular cloud that is bounded by a hot and tenuous external medium. Using the same basic model as presented in Paper I (Kudoh \\& Basu 2003), we have carried out a parameter survey by varying the frequency of the driving force and the magnetic field strength of the cloud. We found that the key parameter for the evolution of the cloud is the Alfv\\'en wavelength of the driving force. If the wavelength is larger than the size of the cloud, the cloud is affected less by the waves. The wavelength that is the same order of the cloud size is the most effective in expanding the cloud. This means that turbulent expansion is different than the usual notion of expansion due to a ``turbulent pressure'' in which the wavelengths need to be much smaller than the cloud size. We also studied the effect of a driving force that varies randomly with time and found that significant cloud expansion occurs since there is always a component of the driving that is at a favorable frequency for expanding the cloud. The evolution of our model clouds reveal the following general features: (1) Under the influence of a driving source of Alfv\\'enic disturbances, a cloud shows significant upward and downward motions, with an oscillation time scale that is comparable to the cloud crossing time; (2) After driving is discontinued, the turbulent energy dissipates exponentially with time, with an $e$-folding time essentially equal to the crossing time across the expanded cloud rather than the crossing time of the driving scale; (3) The power spectra show that significant power is generated on scales larger than that of the turbulent driving ($k < H_0^{-1}$); (4) The line-width--size relation is obtained by an ensemble of clouds with different physical parameters which are individually in a time-averaged self-gravitational equilibrium state. The largest amplitude random motions occur in the outer low density regions of a stratified cloud. The magnetic field data is best fit by models with $\\beta_0 \\approx 1$, which implies that the mass-to-flux ratio is very close to the critical value for collapse.   In the case of the random driving force, we require a much larger driving amplitude $\\tilde{a}_d$ than for the sinusoidal case in order to get a comparable expansion of the cloud (see Table 1). This is due to the multiplication by a random number in the range $[-1,1]$ at each time step (see eq. [\\ref{eq:drivingr}]) in the random case. The resulting driving efficiency is quite low since the driving force $F(z,t)$ is often out of phase with $v_y$, resulting in a lesser amount of work done in the random case than in the sinusoidal case for a given value of $\\tilde{a}_d$. A physically meaningful comparison of the energetics of the random and sinusoidal cases can be done by calculating the energy flux of \\Alf waves emerging from just outside the region of the artificial driving. We define the \\Alf energy flux in the $z$-direction to be \\begin{equation} {\\cal F}_A=-\\frac{B_z B_y}{4\\pi} (v_y-\\langle v_m\\rangle) , \\end{equation} and perform the calculation at $z=z_a=0.1 H_0$. We find that the time-averaged energy flux of a random driving model with parameters $\\tilde{a}_d=5000$ and $\\beta_0=1$ is comparable to (actually about twice as large as) that of a sinusoidal driving model with parameters $\\tilde{a}_d=30$, $\\beta_0=1$, and $\\nu_0=1$.  For the random driving case ($\\tilde{a}_d=5000$, $\\beta_0=1$), the time-averaged energy flux measured near the surface of the cloud (at the full-mass position) is about 20\\% of that at $z=z_a$, so that about 80\\% of the input energy is dissipated in the cloud. For the sinusoidal driving case ($\\tilde{a}_d=30$, $\\beta_0=1$, $\\nu_0=1$), the energy flux near the surface is about 30\\% of that at $z=z_a$. The small difference may arise because the random model contains many high frequency waves which dissipate quickly. In each case, the dissipated energy flux measured while the driving term is applied is comparable to $E_T/t_d$, where $E_T$ is the steady-state value of the sum of fluctuating magnetic and kinetic energies during the turbulent driving phase, and $t_d$ is the dissipation time measured after turbulent driving is discontinued. Our obtained value $t_d \\simeq 10t_0$, which does not depend sensitively on most parameters, is somewhat longer than that estimated from three-dimensional periodic box simulations. We believe that part of the discrepancy comes from the generation of longer wavelength modes as the waves travel to low-density regions near the cloud's surface. However, one-dimensional simulations are also known to have lower dissipation rates than two- or three-dimensional ones (Ostriker et al. 2001). Therefore, higher dimensional simulations that include density-stratification of the cloud will give the final answer in the future.  In our ideal MHD model there is no wave damping due to ion-neutral friction, although there is wave damping due to a nonlinear cascade and nonlinear steepening, as well as leakage of wave energy into the external medium (see Paper I for a more extensive discussion). We expect that ion-neutral friction would be most effective at damping the shortest wavelength modes in our simulation (Kulsrud \\& Pearce 1969), and that the power spectra of $B_y$ and $v_y$ may become steeper at high wavenumber. However, our current results show that the wavelengths that are comparable to the initial cloud size are most effective at expanding the cloud, so it is not clear exactly how effective ion-neutral friction would be in removing turbulent cloud support. Furthermore, we are modeling the large scale evolution of molecular clouds with mean column densities of up to a few times 10$^{21} \\, \\cms$, or visual extinction $A_V \\lesssim 3$. At these column densities, the clouds are expected to be significantly ionized by background far-ultraviolet starlight so that ion-neutral friction will take a prohibitively long time (McKee 1989). Our results show that the dense midplane of the cloud has transonic or subsonic motions, and we expect these regions to decouple from the more turbulent low density envelope and develop higher column density regions in which the far-ultraviolet background is shielded (McKee 1989). In these regions, the ionization fraction is significantly lesser, and ambipolar diffusion can enhance the process of core formation. Therefore, models of fragmentation of a thin-layer which include the effects of magnetic fields and partial ionization due to cosmic rays (Basu \\& Ciolek 2004) form a complementary approach to ours.  Our model has the advantage of being global by virtue of including the effect of cloud stratification and a cloud boundary, rather than being a local periodic box model. The 1.5-dimensional approximation has also allowed extremely high resolution which can more accurately track the propagation of MHD waves. Our localized turbulent driving source is an artificial construct that mimics the various possible sources of disturbances in a real cloud. However, a general result is that the wave energy gets quickly redistributed throughout the cloud in such a way that most of the energy is on the largest scales. This result provides a way out of the ``luminosity problem'' posed by Basu \\& Murali (2001), since the dissipation rate is then set by the crossing time of the cloud scale rather than that of the internal driving scale. Our modeled clouds have a steady-state size $\\sim 1$ pc and large-scale velocity dispersion $\\sim 1 \\kms$,  which makes them comparable to observed dark clouds which can form groupings known as the giant molecular clouds (GMC's; see e.g., Blitz \\& Williams 1999). The turbulent dissipation typically occurs with an $e$-folding time of a few Myr and is completely dissipated by $\\sim 10^7$ yr. Even in the case that real molecular clouds and GMC's start their life in a turbulent state and then evolve with {\\em no} further external turbulent input, our results show that the turbulent lifetime is less than but within reach of estimated GMC lifetimes, $\\sim 10^7$ yr; see discussion in Larson (2003). This provides another important motivation for future {\\em global} three-dimensional models which can follow the actual sources of internal turbulent driving. As turbulence initially dissipates and a cloud contracts, we expect a transformation of some released gravitational potential energy back into MHD wave energy. In addition, much of the internal driving may be attributable to the outflows commonly launched from the near environs of protostars (see e.g., Bachiller \\& Tafalla 1999)."
    },
    "0601/astro-ph0601628_arXiv.txt": {
        "abstract": "We present a joint gravitational lensing and stellar dynamical analysis of fifteen massive field early-type galaxies -- selected from the {\\sl Sloan Lens ACS} (SLACS) Survey -- using {\\sl Hubble Space Telescope} ACS images and luminosity weighted stellar velocity dispersions obtained from the Sloan Digital Sky Survey database.  The sample of lens galaxies is well-defined (see Paper I), with a redshift range of $z$=0.06--0.33 and an average stellar velocity dispersion of $\\langle \\sigma_{\\rm ap} \\rangle = 263$\\,km\\,s$^{-1}$ (rms of 44\\,km\\,s$^{-1}$) inside a 3-arcsec fiber diameter. The following numerical results are found: (i) A joint-likelihood gives an average logarithmic density slope for the {\\sl total} mass density of $\\langle \\gamma' \\rangle = 2.01^{+0.02}_{-0.03}$ (68\\% C.L.; $\\rho_{\\rm tot}\\propto r^{-\\gamma'}$) inside $\\langle {\\rm R}_{\\rm Einst} \\rangle = 4.2 \\pm 0.4$\\,kpc (rms of 1.6\\,kpc). The inferred {\\it intrinsic} rms spread in logarithmic density slopes is $\\sigma_{\\gamma'}=0.12$, which might still include some minor systematic uncertainties. A range for the stellar anisotropy parameter $\\beta=[-0.25, +0.25]$ results in $\\Delta\\langle \\gamma'\\rangle=[+0.05,-0.09]$. Changing from a Hernquist to a Jaffe luminosity density profile increases $\\langle \\gamma' \\rangle$ by~0.05. (ii) The average position-angle difference between the light distribution and the total mass distribution is found to be $\\langle \\Delta \\theta\\rangle = 0 \\pm 3$~degrees (rms of 10 degrees), setting an upper limit of $\\langle \\gamma_{\\rm ext}\\rangle \\la 0.035$ on the average external shear.  The total mass has an average ellipticity $\\langle q_{\\rm SIE}\\rangle$=0.78$\\pm$0.03 (rms of 0.12), which correlates extremely well with the stellar ellipticity, $q_*$, resulting in $\\langle q_{\\rm SIE}/q_* \\rangle = 0.99 \\pm 0.03$ (rms of 0.11) for $\\sigma \\ga 225$~km\\,s$^{-1}$. At lower velocity dispersions, inclined S0 galaxies dominate, leading to a higher ratio (up to 1.6). This suggests that the dark-matter halo surrounding these galaxies is less flattened than their stellar component. Assuming an oblate mass distribution and random orientations, the distribution of ellipticities implies $\\langle q_3 \\rangle\\equiv \\langle(c/a)_\\rho\\rangle =0.66$ with an error of $\\sim$0.2.  (iii) The average projected dark-matter mass fraction is inferred to be $\\langle f_{\\rm DM} \\rangle =0.25 \\pm 0.06$ (rms of 0.22) inside $\\langle {\\rm R}_{\\rm E}\\rangle$, using the stellar mass-to-light ratios derived from the Fundamental Plane as priors.  (iv) Combined with results from the {\\sl Lenses Structure \\& Dynamics} (LSD) Survey at $z \\ga 0.3$, we find no significant evolution of the total density slope inside one effective radius for galaxies with $\\sigma_{\\rm ap}\\ge 200$~km\\,s$^{-1}$: a linear fit gives $\\alpha_{\\gamma'} \\equiv d\\langle \\gamma' \\rangle/dz=0.23\\pm0.16$ (1\\,$\\sigma$) for the range $z$=0.08--1.01. We conclude that massive early-type galaxies at $z$=0.06--0.33 on average have an isothermal logarithmic density slope inside half an effective radius, with an intrinsic spread of at most 6\\,\\% (1\\,$\\sigma$). The small scatter and absence of significant evolution in the inner density slopes suggest a collisional scenario where gas and dark matter strongly couple during galaxy formation, leading to a total mass distribution that rapidly converge to dynamical {\\sl isothermality}. ",
        "introduction": "Massive early-type galaxies are postulated to be latecomers in the hierarchical formation process (e.g. Blumenthal et al.\\ 1984; Frenk et al.\\ 1985), formed via mergers of lower-mass (disk) galaxies (e.g.\\ Toomre \\& Toomre 1972; Schweizer 1982; Frenk et al.\\ 1988; White \\& Frenk 1991; Barnes 1992; Cole et al. 2000). As such, a detailed study of their structure (e.g.\\ Navarro, Frenk \\& White 1996; Moore et al.\\ 1998), formation and subsequent evolution provides a powerful test of the concordance $\\Lambda$CDM paradigm (e.g.\\ Riess et al.\\ 1998; Perlmutter et al.\\ 1999; Spergel et al.\\ 2003; Tegmark et al.\\ 2004) at galactic scales.  In this context, the merging of low-mass galaxies to form more massive ones naively seems to imply a continuous evolution of their mass structure (e.g.\\ Bullock et al.\\ 2001), both in their outer regions and dense inner regions (e.g.\\ smaller galaxies accrete and sink to the center through dynamical friction). On the one hand, the inner regions of massive ellipticals can contract into an increasingly denser structure, if significant mass in dissipational gas is accreted (e.g.\\ Blumenthal et al.\\ 1986; Ryden \\& Gunn 1987; Navarro \\& Benz 1991; Dubinski 1994; Jesseit, Naab \\& Burkert 2002; Gnedin et al.\\ 2004; Kazantzidis et al.\\ 2004). If this process occurs at $z$\\,$\\la$\\,1 (e.g.\\ Kauffmann, Charlot \\& White 1996; Kauffmann \\& Charlot 1998) and results in star-formation activity, one can test this scenario directly by using high-quality data of early-type galaxies, obtained with space and 8--10\\,m class ground-based telescopes (e.g.\\ Menanteau et al.\\ 2001a\\&b; Gebhardt et al. 2003; McIntosh et al.\\ 2005; Tran et al.\\ 2005). On the other hand, the mass inside the inner $\\sim$10\\,kpc of the most massive ellipticals (i.e.\\ $>$L$_*$) seems to remain nearly constant from $z\\gg1$ to the present day -- as suggested by collisionless dark-matter simulations -- and additionally accreted dark matter replaces already-present collisionless matter (e.g.\\ Wechsler et al.\\ 2002; Zhao et al.\\ 2003; Gao et al.\\ 2004). If most gas is turned into (collisionless) stellar mass before mergers (e.g. van Dokkum et al. 1999), one expects it to behave similarly to dark matter during assembly in to the more massive galaxies seen at $z \\la 1$.  Based on the notion that the velocity function of massive early-type galaxies at $z$=0 (Sheth et al.\\ 2003) is remarkably close to that of the inner regions (inside $\\sim$10~kpc) of the most massive simulated galaxies at $z$\\,$\\approx$\\,6 -- even though these galaxies continue to accrete collisionless matter below that redshift -- Loeb \\& Peebles\\ (2003) suggest that the inner regions might behave as {\\sl dynamical attractors}, whose phase-space density is nearly invariant under the accretion of collisionless matter (see also e.g.\\ Gao et al.\\ 2004; Kazantzidis, Zentner \\& Kravtsov 2005). In this scenario, one might expect less structural evolution of the inner regions of massive early-type galaxies at $z<1$, compared to models where most gas had not yet turned into stars {\\sl before} the mass assembly of their inner regions took place.  Hence, one way to study the formation scenario of massive ellipticals, is to quantify the evolution of the mass distribution in their inner regions from redshifts $z$=1 to~0.  \\begin{figure*}[t] \\begin{center} \\leavevmode \\vbox{% \\vbox{% \\epsfxsize=0.70\\hsize \\epsffile{fig1a.epsi} \\epsfxsize=0.70\\hsize \\epsffile{fig1b.epsi} \\epsfxsize=0.70\\hsize \\epsffile{fig1c.epsi} \\epsfxsize=0.70\\hsize \\epsffile{fig1d.epsi}} } \\end{center} \\caption{Non-parametric lens-image reconstructions of confirmed SLACS lens systems. For each system, the observed (galaxy-subtracted) HST-ACS F814W image is shown (left panel), the best reconstruction of the system (middle panel) and source model (right panel), assuming a SIE mass model. From top to bottom are shown: J0037$-$0942, J0216$-$0813, J0737+3216, and J0912+0029 (see Paper I and Table~\\ref{tab:results}).} \\label{fig:siemodels} \\end{figure*}  \\addtocounter{figure}{-1}  \\begin{figure*}[t] \\begin{center} \\leavevmode \\vbox{% \\epsfxsize=0.70\\hsize \\epsffile{fig1e.epsi} \\epsfxsize=0.70\\hsize \\epsffile{fig1f.epsi} \\epsfxsize=0.70\\hsize \\epsffile{fig1g.epsi} \\epsfxsize=0.70\\hsize \\epsffile{fig1h.epsi} } \\end{center} \\caption{(Continued) From top to bottom are shown: J0956+5100, J0959+0410, J1250+0523 and J1330$-$0148.} \\end{figure*}  \\addtocounter{figure}{-1}  \\begin{figure*}[t] \\begin{center} \\leavevmode \\vbox{% \\epsfxsize=0.70\\hsize \\epsffile{fig1i.epsi} \\epsfxsize=0.70\\hsize \\epsffile{fig1j.epsi} \\epsfxsize=0.70\\hsize \\epsffile{fig1k.epsi} \\epsfxsize=0.70\\hsize \\epsffile{fig1l.epsi} } \\end{center} \\caption{(Continued) From top to bottom are shown: J1402+6321, J1420+6019, J1627$-$0055 and J1630+4520.} \\end{figure*}  \\addtocounter{figure}{-1}  \\begin{figure*}[t] \\begin{center} \\leavevmode \\vbox{% \\epsfxsize=0.70\\hsize \\epsffile{fig1m.epsi} \\epsfxsize=0.70\\hsize \\epsffile{fig1n.epsi} \\epsfxsize=0.70\\hsize \\epsffile{fig1o.epsi} } \\end{center} \\caption{(Continued) From top to bottom are shown: J2300+0022, J2303+1422 and J2321$-$0939.} \\end{figure*}  From the observational point of view, a significant effort has been devoted in the past two decades to the study of the mass structure of early-type galaxies in the local Universe ($z$$\\la$0.1) through stellar dynamical tracers and X-ray studies (e.g.\\ Fabbiano 1989; Mould et al. 1990; Matsushita et al. 1998; Loewenstein \\& White 1999; Saglia et al. 1992; Bertin et al. 1994; Arnaboldi et al. 1996; Franx et al. 1994; Carollo et al. 1995; Rix et al. 1997; Gerhard et al. 2001; Seljak 2002; Borriello et al. 2003; Romanowsky et al. 2003). In a comprehensive study, Gerhard et al.\\ (2001) conclude that massive ellipticals have, on average, flat circular velocity curves with a scatter of $\\sim$10\\% in their inner two effective radii. This in itself should impose stringent constraints on any numerical simulations of elliptical galaxies (e.g.\\ Meza et al.\\ 2003; Kawata \\& Gibson 2003).  Strong gravitational lensing provides a complementary approach (Kochanek 1991) to study early-type galaxies at higher redshifts. Lensing analysis has been used to demonstrate the presence of dark matter around early-type galaxies and, in some systems, to provide evidence for ``isothermal'' (i.e. $\\rho_{\\rm tot}\\propto r^{-2}$) mass density profiles equivalent to the flat rotation curves observed for spiral galaxies (e.g., Kochanek 1995; Rusin \\& Ma 2001; Ma 2003; Rusin et al. 2002, 2003a; Cohn et al. 2001; Munoz et al. 2001; Winn et al. 2003; Wucknitz et al. 2004; Rusin \\& Kochanek 2005). However, the mass-profile (e.g.\\ Wucknitz 2002) and mass-sheet degeneracies (Falco et al.\\ 1985) often prevent a truly accurate determination of the logarithmic density slope {\\sl at} the Einstein radius.  To answer the question ``{\\sl What is the mass structure inside the inner regions of early-type galaxies and how does it evolve with time?\\,}'', we therefore combine constraints from strong gravitational lensing and stellar kinematics. The former provides an accurate mass measurement inside the Einstein radius (Kochanek 1991), whereas the latter provides a measurement of the mass gradient. The average logarithmic density slope inside the Einstein radius can then be determined -- independent of the mass-sheet degeneracy that is associated with the galaxy mass distribution -- with the same fractional accuracy as is obtained on the luminosity-weighted stellar velocity dispersion (e.g.\\ see Treu \\& Koopmans 2002 and Koopmans 2004).  The density slopes of {\\sl individual} early-type galaxy can be correlated with redshift, to determine any structural evolution in the population.  In an ongoing study of massive early-type lens galaxies between $z \\approx 0.5$ and 1, as part of the {\\sl Lenses Structure \\& Dynamics} (LSD) Survey (Koopmans \\& Treu 2002, 2003; Treu \\& Koopmans 2002, 2003, 2004; hereafter TK04) -- plus two additional systems that were studied to measure the Hubble Constant (Treu \\& Koopmans 2002; Koopmans et al 2003) -- this technique has successfully been applied, to place the first constraints on the inner density slopes and dark-matter halos of early-type galaxies to $z\\approx 1$ (TK04), finding a logarithmic density slope close to isothermal, although the results were limited by the small sample size.  In the first paper of this series (Bolton et al.\\ 2006; hereafter Paper I), we reported on the discovery of nineteen new early-type lens galaxies from the {\\sl Sloan Lens ACS} (SLACS) Survey at $z \\la 0.3$, each with {\\sl Hubble Space Telescope} (HST) F435W and F814W images and a stellar velocity dispersion measured from their SDSS spectra (e.g.\\ Bolton et al.\\ 2004). Some systems have integral field spectroscopy (IFS) of their lensed sources, obtained with Magellan and/or Gemini (see also Bolton et al.\\ 2005). As far as their photometric properties are concerned, they are representative of Luminous Red Galaxies (LRG; Eisenstein et al.\\ 2004) with similar redshifts and similar stellar velocity dispersions (see Paper I). They also lie on the Fundamental Plane (FP; Dressler et al. 1987; Djorgovski \\& Davis 1987) of early-type galaxies, have old stellar populations, and have very homogeneous mass density profiles (Treu et al.\\ 2006; hereafter Paper II).  In this paper, we focus on the analysis of a sub-sample of fifteen isolated early-type lens galaxies from SLACS, combining the constraints from {\\sl Hubble Space Telescope} (HST) images with the stellar velocity dispersion obtained from the SDSS database. The goals are to quantify their inner mass structure and to assess whether any evolution of their inner regions has occurred at~$ z \\la 1$. In \\S\\,2, we present non-parametric lens models for each system. The simplicity of the models supports their lensed nature and provides the necessary input for subsequent analysis. In \\S\\,3, we use the enclosed mass from lensing in a joint stellar-dynamical analysis, to determine the inner density slopes of each early-type galaxy.  In combination with results from the LSD survey, we analyze the redshift behavior of the density slope in \\S\\,4 to quantify its evolution. In \\S\\,5, we summarize our results and draw conclusions. Throughout this paper, we assume H$_0$=70\\,km\\,s$^{-1}$\\,Mpc$^{-1}$, $\\Omega_{\\rm m}=0.3$ and $\\Omega_\\Lambda = 0.7$. ",
        "conclusions": "The {\\sl Sloan Lens ACS} (SLACS) Survey has provided the largest uniformly selected sample of massive early-type lens galaxies to date (Papers I \\& II). We used a sub-sample of fifteen early-type lens galaxies with a redshift range of $z=0.06-0.33$ and an unweighted average stellar velocity dispersion of $\\langle \\sigma_{\\rm ap} \\rangle = 263 \\pm 11$\\,km\\,s$^{-1}$ for a joint lensing and dynamical analysis, finding the following results:  \\begin{itemize}  \\item The average logarithmic density slope of the {\\sl total} mass density of $\\langle \\gamma' \\rangle = 2.01^{+0.02}_{-0.03}$ (68\\% C.L.) assuming a total density profile of $\\rho_{\\rm tot}\\propto r^{-\\gamma'}$ and no anisotropy (i.e.\\ $\\beta=0$). Systematic uncertainties (e.g.\\ orbital anisotropy and different luminosity density cusps) are expected to be $\\la 5-10\\%$ [see \\S\\,3.2].  \\item The {\\sl intrinsic} spread in the logarithmic density slope is at most 6\\%, i.e.\\ $\\sigma_{\\gamma'}=0.12$ (the 1\\,$\\sigma$ of the assumed Gaussian spread) [see \\S\\,3.2], after accounting for measurement errors.  \\item The average position-angle difference between the stellar light component and the total mass component is found to be $\\langle \\Delta \\theta\\rangle = 0 \\pm 3$~degrees with 10~degrees rms, setting an upper limit of $\\langle \\gamma_{\\rm ext}\\rangle \\la 0.035$ on the average external shear. [see \\S\\,2.4].  \\item The ellipticity of the total surface-density is $\\langle q_{\\rm SIE}\\rangle$=0.78$\\pm$0.03 (rms of 0.12) and $\\langle q_{\\rm SIE}/q_* \\rangle = 0.99 \\pm 0.03$ (rms of 0.11) for $\\sigma \\ga 225$~km\\,s$^{-1}$. Assuming an oblate mass distribution and random orientations, this implies $\\langle q_3 \\rangle\\equiv \\langle(c/a)_\\rho\\rangle =0.7$ with an error of 0.2 [see \\S\\,2.4].  \\item The unweighted average projected dark-matter mass fraction is $\\langle f_{\\rm DM} \\rangle =0.25 \\pm 0.06$ (rms of 0.22) inside $\\langle {\\rm R}_{\\rm Einst} \\rangle = 4.2 \\pm 0.4$\\,kpc (rms of 1.6\\,kpc) [see \\S\\,3.5].  \\item The evolution of the total density slope for galaxies with $\\sigma_{\\rm ap}\\ge 200$~km\\,s$^{-1}$ ($>$L$_*$), inside half an effective radius, $\\langle \\gamma'\\rangle (z)=(2.10 \\pm 0.07) - (0.23\\pm0.16)\\cdot z$ for the range $z=0.08-1.01$ (combined sample from the SLACS and LSD Surveys) [see \\S\\,4].  \\end{itemize}  \\noindent {\\it Summarising: Massive early-type galaxies below $z$\\,$\\approx$\\,1 have remarkably homogeneous inner mass density profiles, i.e.\\ $\\rho_{\\rm tot} \\propto r^{-2}$ (equivalent to a flat rotation curve for a rotation supported system), and very close alignment between stellar and total mass. There is no evidence for significant evolution in the ensemble average logarithmic density slope of dark plus stellar mass below a redshift of one in their inner half to one effective radius.}"
    },
    "0601/astro-ph0601134_arXiv.txt": {
        "abstract": "We study the environment-induced decoherence of cosmological perturbations in an inflationary background. Splitting our spectrum of perturbations into two distinct sets characterized by their wavelengths (super and sub-Hubble), we identify the long wavelength modes with our system and the remainder with an environment. We examine the effects of the interactions between our system and the environment. This interaction causes the long-wavelength modes to decohere for realistic values of the coupling and we conclude that interactions due to backreaction are more than sufficient to decohere the system within 60 e-foldings of inflation. This is shown explicitly by obtaining an analytic solution to a master equation detailing the evolution of the density matrix of the system. ",
        "introduction": "As is well known, temperature fluctuations in the CMB and the inhomogeneities that seed structure formation in the universe share a common origin. Both are a result of the scalar metric perturbations produced during inflation. However, these perturbations are of a purely quantum mechanical nature while no cosmological systems of interest (CMB anisotropies, clusters etc.) display any quantal signatures. Presumably, for this to be the case, the primordial density perturbations underwent a quantum-to-classical transition some time between generation during inflation and recombination, when structure first became apparent.  Decoherence is a much studied process (see \\cite{Giulini:1996nw} for a comprehensive review). Although not all conceptual issues have been resolved, it is understood that it can occur whenever a quantum system interacts with an ''environment''. In other words, this effect can be said to pervade open systems due to the difficulty of creating a truly closed, macroscopic quantum system. Along with its ubiquity, it is also known to be a practically irreversible process, since the loss of quantum correlations in the system is accompanied by an increase in entropy.  Early studies of the classicalization of primordial perturbations focussed on intrinsic properties of the system (see, for example \\cite{Guth:1985ya},\\cite{Polarski:1995jg}). This was made possible by the application of ideas of quantum optics to the theory of cosmological perturbations. Primordial density fluctuations (the scalars as well as the tensors) evolve into a peculiar quantum state - a \\textit{squeezed} vacuum state \\cite{Grishchuk:1990bj},\\cite{Albrecht:1992kf}. By studying the large squeezing limit of these states, it was found that quantum perturbations become indistinguishable from a classical stochastic process. In other words, quantum expectation values in a highly squeezed state are identical to classical averages calculated from a stochastic distribution, up to corrections which vanish in the limit of infinite squeezing.\\ The authors of \\cite{Polarski:1995jg} refer to this as ''decoherence without decoherence'' while \\cite{Kiefer:1998qe} endows the phenomenon with the more technical epithet ''quantum non-demolition measurement''. We emphasize that these works focussed on the classical properties of the states and not on the coherence properties of the system.  As is well understood, in order to study true classicalization, one must consider two distinct aspects of a system. First the quantum states must evolve, in some limit, into a set of states analogous to classical configurations. The second is that these resultant states interfere with each other in a negligible fashion. This last property constitutes decoherence and is equivalent to the vanishing of the off-diagonal elements of the density matrix.  A truly closed gravitational system is a practical impossibility (unless one considers the totality of the universe to constitute the system as in, for example, quantum cosmology). Since the gravitational interaction has infinite range and couples to all sources of energy, interactions with some sort of environment are an inevitability. As such, environmentally induced decoherence must also be present and would play an important role in the classicalization of primordial density fluctuations.  The purpose of the present article is to determine precisely the effects by the ''inflationary environment'' (we will elucidate this notion below) on cosmological perturbations and to study the resultant decoherence. Other authors have also examined this problem (see, for example \\cite{Kiefer:1998qe},\\cite{Sakagami:1987mp},\\cite{Brandenberger:1990bx},\\cite{Calzetta:1995ys},\\cite{Lombardo:2005iz}) - however, we are the first to present an exact analytic expression for the density matrix with a realistic environment-system interaction.  The paper is organized as follows: in the next section, we review some basic properties of decoherence of which we will make use. After reviewing the quantum theory of cosmological perturbations in section III, we make clear our concept of the environment and motivate some realistic interactions in section IV. Subsequently, we develop necessary formalism which, in section VI, we make use of to demonstrate the classical nature of the system and calculate the decoherence time scale. ",
        "conclusions": "In this paper, we have studied the decoherence of cosmological fluctuations during a period of cosmological inflation, taking the effects of squeezing into account. We have determined realistic interactions for our system of perturbations and have found that, at the same order, gravitational interactions and  matter (inflaton) interactions are comparable, depending on the scale of inflation and the slow-roll parameter $\\epsilon$. Furthermore, we have justified the use of Hubble scale as a cutoff.  Having considered the leading order gravitational correction to the action of quantized cosmological perturbations, we find that super-Hubble modes decohere long before the end of inflation. Of course, we have only obtained a lower bound on the decoherence rate - interactions more complicated than the ones considered here will generally lead to much faster decoherence times \\cite{Kiefer:1998jk}."
    },
    "0601/astro-ph0601445_arXiv.txt": {
        "abstract": "The microwave sky shows unexpected features at the largest angular scales, among them the alignments of the dipole, quadrupole and octopole. Motivated by recent X-ray cluster studies, we investigate the possibility that local structures at the 100~$h^{-1}$Mpc scale could be responsible for such correlations. These structures give rise to a local Rees--Sciama contribution to the microwave sky that may amount to $\\Delta T/T \\sim 10^{-5}$ at the largest angular scales. We model local structures by a spherical overdensity (Lema\\^itre--Tolman--Bondi model) and assume that the Local Group is falling toward the centre. We superimpose the local Rees--Sciama effect on a statistically isotropic, gaussian sky. As expected, we find alignments among low multipoles, but a closer look reveals that they do not agree with the type of correlations revealed by the data. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601359_arXiv.txt": {
        "abstract": "The galactic environment of the Sun varies over short timescales as the Sun and interstellar clouds travel through space.  Small variations in the dynamics, ionization, density, and magnetic field strength in the interstellar medium (ISM) surrounding the Sun can lead to pronounced changes in the properties of the heliosphere.  The ISM within $\\sim$30 pc consists of a group of cloudlets that flow through the local standard of rest with a bulk velocity of $\\sim$17--19 \\kms, and an upwind direction suggesting an origin associated with stellar activity in the Scorpius-Centaurus association.  The Sun is situated in the leading edge of this flow, in a partially ionized warm cloud with a density of $\\sim$0.3 \\cc. Radiative transfer models of this tenuous ISM show that the fractional ionization of the ISM, and therefore the boundary conditions of the heliosphere, will change from radiative transfer effects alone as the Sun traverses a tenuous interstellar cloud.  Ionization equilibrium is achieved for a range of ionization levels, depending on the ISM parameters.  Fractional ionization ranges of \\chiH=0.19--0.35 and \\chiHe=0.32--0.52 are found for tenuous clouds in equilibrium. In addition, both temperature and velocity vary between clouds. Cloud densities derived from these models permit primitive estimates of the cloud morphology, and the timeline for the Sun's passage through interstellar clouds for the past and future $\\sim$10$^5$ years.  The most predictable transitions happen when the Sun emerged from the near vacuum of the Local Bubble interior and entered the cluster of local interstellar clouds flowing past the Sun, which occurred sometime within the past 140,000 years, and again when the Sun entered the local interstellar cloud now surrounding and inside of the solar system, which occurred sometime within the past 44,000 years, possibly a 1000 years ago.  Prior to $\\sim$140,000 years ago, no interstellar neutrals would have entered the solar system, so the pickup ion and anomalous cosmic ray populations would have been absent.  The tenuous ISM within 30 pc is similar to low column density observed globally.  In this chapter, we review the factors important to understanding short-term variations in the galactic environment of the Sun.  Most ISM within 40 pc is partially ionized warm material, but an intriguing possibility is that tiny cold structures may be present. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601390_arXiv.txt": {
        "abstract": "{} {The dynamics of large-scale structure formation in the universe by gravitational instability still presents many open issues and is mostly studied through numerical simulations. This motivates the study of simpler models which can be investigated by analytical means in order to understand the main processes at work. Thus, we describe here a one-dimensional self-gravitating system derived from the cosmological context, which leads to an effective external potential in addition to the standard gravitational self-interaction. As a first step we consider small times so that the expansion can be neglected. Then we present a thermodynamical analysis of this system as well as the stability properties of the associated hydrodynamical and collisionless systems.} {We consider the mean field limit (i.e. continuum limit) to perform an analytical study.} {We find a second-order phase transition at $T_{c1}$ from an homogeneous equilibrium at high temperature to a clustered phase (with a density peak at one of the boundaries of the system) at low temperature. There also exists an infinite series of unstable equilibria which appear at lower temperatures $T_{cn}$, reflecting the scale-free nature of the gravitational interaction and the usual Jeans instability. We find that, as for the similar HMF (Hamiltonian mean field) model, all three micro-canonical, canonical and grand-canonical ensembles agree with each other, as well as with the stability properties associated with a hydrodynamical approach. On the other hand, the collisionless dynamics governed by the Vlasov equation yields the same results except that at low $T$ the equilibrium associated with two density peaks (one at each boundary) becomes stable.} {} ",
        "introduction": "\\label{Introduction}   In standard cosmological scenarios the large scale structures we observe in the present universe (such as clusters, filaments or voids) have formed through the amplification by gravitational instability of small primordial perturbations, see Peebles (1980). Moreover, the amplitude of these density fluctuations increases at small scales as in the CDM model (Peebles 1982) which leads to a hierarchical process where smaller scales become non-linear first. Then, at large scales or at early times one can use a perturbative approach to study the evolution of initial fluctuations (Fry 1984; Goroff et al. 1986; Bernardeau 1992; Valageas 2001). Next, the weakly non-linear regime may be investigated through the Zeldovich approximation (Shandarin \\& Zeldovich 1989) or the adhesion model (Gurbatov et al. 1989). However, the highly non-linear regime (which corresponds to collapsed structures such as clusters of galaxies) has mostly remained out of reach of systematic approaches. Thus, one may use the Press-Schechter approximation (Press \\& Schechter 1974) or the excursion-set formalism of Bond et al. (1991) to obtain the statistics of just-virialized halos or the saddle-point approach of Valageas (2002) for rare voids. These approximations focus on specific objects which may be ``recognized'' from individual features in the linear density field itself and do not follow the system as a whole through its non-linear evolution. A systematic method to do so was recently developed in Valageas (2004) but its application to the collisionless dynamics has not been performed yet for the highly non-linear regime. Therefore, the non-linear regime of cosmological structure formation is mostly studied through numerical simulations.  In order to simplify this difficult problem one can investigate one-dimensional (1-D) systems which are easier to study both through numerical and analytical approaches. In fact, 1-D systems such as the system of parallel mass sheets (Camm 1950) have been studied for a long time to explore the evolution of isolated $N-$body systems which only interact through classical gravity. In particular, they have been used to investigate relaxation processes and to test Lynden-Bell's prediction for the final state of violent relaxation (Lynden-Bell 1967). Such 1-D systems can already exhibit complex behaviours. For instance, Luwel \\& Severne (1985) found that collisional effects are not sufficient to relax the stationary waterbag configuration towards thermodynamical equilibrium whereas Rouet \\& Feix (1999) showed that holes in phase-space in the initial distribution function can persist over long times and prevent efficient relaxation. On the other hand, the system can break up into smaller clusters as in Hohl \\& Feix (1967) or develop a fractal structure as in Koyama \\& Konishi (2001).  In this paper, we consider the 1-D gravitational system obtained by studying 1-D density fluctuations in a 3-D cosmological background. This model has already been studied mostly through numerical simulation in Aurell \\& Fanelli (2001), Aurell et al. (2001) and Fanelli \\& Aurell (2002), both with and without cosmological expansion. Here we restrict ourselves to time-scales much smaller than the Hubble time so that the expansion of the universe can be neglected and we focus on a mean field analysis (i.e. a continuum limit) which is relevant in the cosmological context. We study the thermodynamics and stability of this system, similarly to the theoretical analysis of Chavanis et al. (2005) performed for the HMF model (defined by a cosine interaction) which shows a similar behaviour. We introduce this model in sect.~\\ref{cosmological-gravitational-system}, where the effect of the cosmological background is seen to reduce to an effective external potential $V$ which balances the gravitational self-interaction $\\Phi$ so that the homogeneous state is an equilibrium solution. Next, we perform a thermodynamical analysis of this system in sect.~\\ref{Thermodynamical-analysis}, for the micro-canonical, canonical and grand-canonical ensembles. We study both the possible equilibrium states and their stability properties. Then, in sect.~\\ref{Hydrodynamical-model} we investigate the stability of such equilibria within an hydrodynamical framework whereas we consider the collisionless case as described by the Vlasov dynamics in sect.~\\ref{Vlasov}. Finally, we present our conclusions in sect.~\\ref{Conclusion}. ",
        "conclusions": "\\label{Conclusion}   We have studied in this article the thermodynamical properties of a 1-D gravitational system built from the problem of large-scale structure formation in cosmology. We have considered time-scales much smaller than the Hubble time so that the expansion of the universe can be neglected and the system can be restricted to a finite size $L$. Then, the cosmological background merely yields an effective external potential $V$ in addition to the gravitational self-interaction $\\Phi$ so that an homogeneous equilibrium (corresponding to the Hubble flow) exists at all temperature $T$. We have found that all three micro-canonical, canonical and grand-canonical ensembles give the same results. There exists a series of critical temperatures $T_{cn}$ ($T_{c1}>T_{c2}>..$) below which two new equilibria $\\pm n$ appear, reflecting the usual Jeans instability mechanism. Above $T_{c1}$ the homogeneous state is thermodynamically stable whereas below $T_{c1}$ it is unstable and only equilibria $\\pm 1$ (corresponding to a density peak at the boundary $x=0$ or $L$) are stable while other states $\\pm n$ with $n\\geq 2$ are unstable. The system shows a second-order phase transition at $T_{c1}$ where the specific heat is discontinuous.  Next, we have studied the hydrodynamical properties of these equilibria. We have found that both the non-linear stability and the dynamical linear stability coincide with the results of the thermodynamical analysis. The evolution of the system can be understood in a simple manner from the translation and diffusion of the various density peaks associated with states $\\pm n$ which would relax towards the stable equilibria $\\pm 1$. Finally, we have studied the collisionless dynamics governed by the Vlasov equation. The equilibria $\\pm 1$ are again both linearly and non-linearly stable while the homogeneous state is unstable below $T_{c1}$. We found that the equilibrium $n=-2$ (one density peak at $x=L/2$) is again unstable but the equilibrium $n=2$ (two density peaks located at both boundaries) shows a new behaviour. Although it is unstable just below $T_{c2}$ as in the thermodynamical and hydrodynamical approaches it now turns stable at low temperatures. This departure from the results of the thermodynamical and hydrodynamical analysis is due to the dynamical constraints associated with the equations of motion of a collisionless system. On the other hand, other equilibria $\\pm n$ with $n\\geq 3$ are again unstable. Therefore, depending on the initial conditions the system would relax towards equilibria $\\pm 1$ or $2$ (i.e. a density peak at either one or both boundaries).  Thus, the behaviour of the 1-D gravitational system (OSC) studied in this article is quite similar to the HMF model, described for instance in Chavanis et al. (2005). We obtain the same second-order phase transition from a high-temperature homogeneous equilibrium to a clustered equilibrium $\\pm 1$ at low temperature. The stability properties of these two configurations are also similar, for all three thermodynamical, hydrodynamical and collisionless analysis. However, in the case of the OSC system the scale-free nature of the gravitational interaction leads to an additional series of equilibria $\\pm n$ ($n\\geq 2$). These new states are unstable, except for the equilibrium $n=2$ which turns linearly stable for a collisionless dynamics at low $T$, but remains unstable for both the thermodynamical and hydrodynamical approaches. This provides an interesting case where the collisionless dynamics exhibits a specific behaviour and leads to a relaxation which clearly depends on initial conditions. On the other hand, it would be interesting to study the non-linear evolution of the system and the possible merging of several density peaks in both the hydrodynamical and collisionless cases. This might be studied through $N-$body simulations as in Hohl \\& Feix (1967) or through a direct computation of the Vlasov dynamics as in Alard \\& Colombi (2005). Finally, one should include the expansion of the universe to study the non-equilibrium dynamics associated with cosmological hierarchical scenarios. This is left for future studies but we can expect this 1-D gravitational system to provide a useful tool to investigate such processes."
    },
    "0601/astro-ph0601673_arXiv.txt": {
        "abstract": "{ We present a measurement of the sky distribution of positronium (Ps) annihilation continuum emission obtained with the SPI spectrometer on board ESA's INTEGRAL observatory. The only sky region from which significant Ps continuum emission is detected is the Galactic bulge. The Ps continuum emission is circularly symmetric about the Galactic centre, with an extension of about $8^\\circ$ FWHM. Within measurement uncertainties, the sky distribution of the Ps continuum emission is consistent with that found by us for the 511~keV electron-positron annihilation line using SPI. Assuming that 511~keV line and Ps continuum emission follow the same spatial distribution, we derive a Ps fraction of $0.92\\pm0.09$. These results strengthen our conclusions regarding the origin of positrons in our Galaxy based on observations of the 511~keV line.  In particular, they suggest that the main source of Galactic positrons is associated with an old stellar population, such as Type~Ia supernovae, classical novae, or low-mass X-ray binaries. Light dark matter is a possible alternative source of positrons. ",
        "introduction": "Introduction}  The annihilation of positrons with electrons gives rise to two characteristic emissions at gamma-ray energies: the hallmark line at 511~keV, and the unique three-photon positronium (Ps) continuum emission \\citep[cf.][]{Guessoum2005}. Direct annihilation of positrons with electrons, and their annihilation via the formation of para-Ps (with the spins of electron and positron being anti-parallel), result in the emission of two 511~keV photons. Annihilation via the formation of ortho-Ps (with the spins of electron and positron being parallel) produces three photons and gives rise to the Ps continuum emission, which is roughly saw-tooth shaped with a peak at the maximum energy of 511~keV \\citep{Ore_Powell49}.  Cosmic positron annihilation radiation was first detected from the Galactic centre (GC) direction in balloon observations during the 1970s and has been the focus of intense scrutiny by a large number of balloon and satellite borne experiments ever since \\citep[see e.g.\\ the reviews by][]{Tueller92,Harris97}. Despite tremendous observational and theoretical efforts, the origin of the positrons is still poorly understood. A large variety of positron sources and production mechanisms have been proposed over the years \\citep[e.g.][]{Chan_Lingenfelter93, Dermer_Murphy01}. Among the more promising source candidates are radioactive nucleosynthesis products from supernovae. More recently, hypernovae/GRBs \\citep{Casse04, Parizot05} and light dark matter \\citep[e.g.][]{Boehm04} have been proposed as possible candidates. Compact objects comprise another potential candidate source of positrons.  Investigations of the sky distribution of the annihilation radiation promise to provide clues to the identification of the source(s) of positrons in our Galaxy, despite the fact that positrons may travel from their birth places before annihilating. First maps of the annihilation radiation, limited to the inner regions of our Galaxy, were obtained using the OSSE instrument on board the Compton Gamma-Ray Observatory \\citep{Johnson93} in the 511~keV line and in Ps continuum emission \\citep[e.g.][]{Purcell97, Chen97, Milne00, Milne01a, Milne01b, Milne02}. Furthermore, the OSSE instrument allowed \\citet{Kinzer99, Kinzer01} to study the one-dimensional distribution in longitude and in latitude of diffuse emission, including annihilation radiation, from the inner Galaxy. With the commissioning of the imaging spectrometer SPI on board ESA's INTEGRAL observatory, high spectral resolution mapping with improved angular resolution has become feasible \\citep{Jean03a, Weidenspointner04, Knoedlseder05}. The 511~keV line emission is found to be dominated by the Galactic bulge and/or halo; emission from the Galactic disk is much fainter, implying that positron annihilation is concentrated in the central regions of our Galaxy \\citep{Knoedlseder05}.  A first SPI measurement of the flux in Ps continuum emission, and of the Ps fraction $f_{Ps}$ \\citep[the fraction of positrons that annihilate through Ps formation;][]{Brown_Leventhal87}, has been presented by \\citet{Churazov05}. The value of $f_{Ps}$, as well as the detailed shape of the 511~keV annihilation line, depend on the physical properties of the annihilation media; therefore detailed spectroscopy of the positron annihilation can provide unique information on the annihilation media and processes \\citep{Guessoum2005, Jean06}.         In this publication, we present results concerning the Galactic distribution of Ps continuum emission using observations of most of the celestial sphere with the spectrometer SPI/INTEGRAL. The analysis of the SPI observations is presented in Sec.~\\ref{data_analysis}. Our mapping, model fitting, and spectral results are given in Sec.~\\ref{results}. A summary and our conclusions can be found in Sec.~\\ref{summary}. ",
        "conclusions": ""
    },
    "0601/astro-ph0601200.txt": {
        "abstract": "We investigate the relative distribution of the gaseous contents of the Universe (as traced by a sample of Lyman $\\alpha$ (\\lya\\ ) absorbers), and the luminous baryonic matter (as traced by a redshift survey of galaxies in the same volume searched for \\lya\\ absorbers), along 16 lines-of-sight (LOS) between redshifts 0 and 1. Our galaxy redshift survey was made with the Multi-Object Spectrograph (MOS) on Canada-France-Hawaii Telescope (CFHT) and, when combined with galaxies from the literature in the same LOS, gives us a galaxy sample of 636 objects. By combining this with an absorption line sample of 406 absorbing systems drawn from published works, we are able to study the relationship between gas and galaxies over the latter half of the age of the Universe. A correlation between absorbers and galaxies is detected out to separation of 1.5 Mpc. This correlation is weaker than the galaxy-galaxy correlation. There is also some evidence that the absorbing systems seen in CIV are more closely related to galaxies, although this correlation could be with column density rather than metallicity. The above results are all consistent with the absorbing gas and the galaxies co-existing in dark matter filaments and knots as predicted by current models where the column density of the absorbing gas is correlated with the underlying matter density. ",
        "introduction": "\\label{sec-intro}  This paper is part of our efforts to provide constraints and measurements of the relationship between the distribution of the gaseous contents of the Universe as traced by neutral Hydrogen and the luminous baryonic mater as traced by galaxies.  The manner in which gas collapses gravitationally into dark matter potential wells to form stars (and hence galaxies), and the way in which these stars then affect the gas is a topic of great interest at present. This is being pursued both observationally and theoretically at redshifts from 7 to zero (i.e. when the universe was approximately 5\\% of its current age, to the present day). In this paper we investigate this process observationally during the second half of the evolution of the universe.  There is a continuing debate about the relationship between low redshift Lyman~$\\alpha$ (\\lya) absorbers and galaxies, where here `low redshift' is taken to mean redshifts less than 1. A simplified (strawman) version of the two sides is (a) that all low redshift \\lya\\ absorbers are part of physically-distinct luminous-galaxy halos, or (b) that they are all part of the filamentary structure seen in recent SPH/Mesh structure formation models, and are only related to galaxies by the fact that both are following the underlying dark matter distribution.  In practice, almost all authors acknowledge that the universe includes a mix of the above two populations (and indeed others), and the debate is more about which population dominates a particular set of observations.  Both of these positions have been vigorously defended in the literature.  The reference list below includes all relevant papers listed on the ADS abstract server from the time of the first available high quality UV spectroscopy from HST to 15 August 2005.  Papers supporting (a) include (in chronological order): \\citet{1994MNRAS.269L..49M}; \\citet{1995ApJ...442..538L}; \\citet{1996ApJ...456L..17L}; \\citet{1998Ap&SS.263...75B}; \\citet{1998ApJ...495..637L}; \\citet{1998ApJ...498...77C}; \\citet{1999ApJ...523...72O}; \\citet{2000ApJ...529..644L}; \\citet{2001ApJ...556..158C}; \\citet{2001ApJ...559..654C}; \\citet{2002ApJ...570..526S}; \\citet{2003ApJ...589..111C}; \\citet{2004ApJ...606..196Z}; \\citet{2004ApJ...609...94S}; \\citet{2004MNRAS.354L..25B}; \\citet{2005A&A...429L...5D}; \\citet{2005ApJ...620...95T}; \\citet{2005ApJ...622..267K}; \\citet{2005ApJ...623...57M} and \\citet{2005ApJ...623..767J}.  Papers supporting (b) include (in chronological order): \\citet{1991ApJ...377L..21M}; \\citet{1993ApJ...419..524M}; \\citet{1994ApJ...427..696M}; \\citet{1994MNRAS.269...52M}; \\citet{1995ApJ...438..650W}; \\citet{1995Natur.373..223D}; \\citet{1995ApJ...451...24S}; \\citet{1996AJ....111...72S}; \\citet{1996ApJ...457...19B}; \\citet{1996ApJ...458..518R}; \\citet{1996A&A...306..691L}; \\citet{1996ApJ...464..141B}; \\citet{1996AJ....112.1397V}; \\citet{1997ApJ...491...45D}; \\citet{1998ApJS..118....1J}; \\citet{1998ApJ...505..506G}; \\citet{1998ApJ...506....1W}; \\citet{1998ApJ...508..200T}; \\citet{1999ApJS..122..355V}; \\citet{1999ApJ...524..536I}; \\citet{2000ApJ...544..150P}; \\citet{2000ApJS..130..121P}; \\citet{2002ApJ...565..720P}; \\citet{2002ApJ...575..697T}; \\citet{2002ApJ...575..697T}; \\citet{2002ApJ...574L.115M}; \\citet{2002ApJ...574..599M}; \\citet{2002ApJ...574L.115M}; \\citet{2002ApJ...580..169B}; \\citet{2003ApJ...591..677R}; \\citet{2003ApJ...591...79M}; \\citet{2004ApJS..152...29P}; \\citet{2004ApJ...614...31B}; \\citet{2004ApJS..155..351S}; \\citet{2005ApJ...618..178C}; \\citet{2005ApJ...624..555D} and \\citet{2005ApJ...629L..25C}.   Finally to complete the reference list, some of the theoretical papers from the last ten years which are relevant for this debate are: \\citet{1996ApJ...457L..51H}; \\citet{1997ApJ...485..496Z}; \\citet{1998ApJ...496..577C}; \\citet{1998MNRAS.297L..49T}; \\citet{1999ApJ...511..521D}; \\citet{1999ApJ...514....1C}; \\citet{2001ApJ...553..528D}; \\citet{2001ApJ...559..507S}; \\citet{2002ApJ...580...42M}; \\citet{2002ApJ...574..590S}; \\citet{2002MNRAS.336..685V}; \\citet{2002ApJ...578L...5T}; \\citet{2002ApJ...580...42M}; \\citet{2004MNRAS.348..421N}; \\citet{2004MNRAS.348..435N}; \\citet{2004ApJ...613..159F}; \\citet{2005ApJ...620L..13A}; \\citet{2005ApJ...623L..97T}; \\citet{2005ApJ...624L...1S}; \\citet{2005MNRAS.359..295F}; \\citet{2005MNRAS.360.1110V} and \\citet{2005MNRAS.361...70J}.  Regardless of whether either hypothesis is correct, we emphasize that even if one (naively) believes a clear answer can be obtained (either (a) or (b) above), this answer will be a function of the neutral Hydrogen column density of the absorbers.  For \\nhi$>10^{21}$ cm$^{-2}$, one is probing a column density of material comparable to that of the disk of our galaxy, and hence all galaxies, to varying degrees, should be expected to contribute to the population of such absorbers. At column densities \\nhi$\\sim10^{12}$ cm$^{-2}$, one is getting close to the neutral hydrogen content of the expected fluctuations within voids, and we expect that a more heterogeneous set of causes could produce such absorption.  The increasingly sophisticated models of galaxy formation and evolution now available (see references above) suggest that if we could make a perfect census of the gaseous and luminous constituents of a large volume of the low redshift Universe, we would find that both the gas (quasar absorption line systems) and stars (galaxies) trace the same fundamental structures (whose course of formation and distribution was set by the early perturbations of the distribution of dark matter in the Universe). These same models, however, also indicate that a wide range of detailed relationships between the gas and galaxies should be observed.  The details of the structures will depend on an equally wide range of interesting astrophysics, including the process of star formation and ``galaxy feedback'' and we hope that the approach of this paper can lead to a better understanding of this astrophysics.  \\begin{figure} \\includegraphics[height=70mm, angle=0]{q0107_image_labelled_gv_new.eps} \\caption{Illustrative image of the two lines-of-sight (LOS) to the QSOs LBQS 0107-025A and B. The region shown is 7.1'$\\times$6.2', with North up and East to the left. The two QSOs discussed in this paper are labelled, as is a third QSO (labelled C) to the North \\citep{2001ApJ...549...76Y}. A few of the measured galaxy redshifts in the field are labelled (drawn so the decimal point is approximately above the relevant galaxy) and the projected spatial offset and velocity difference to detected absorbers in either of the LOS are marked for two example galaxies. See the text for a discussion. The image shown was made from a 90 second exposure without a filter using the Palomar Observatory 5m telescope and COSMIC imager/multi-object spectrograph on September 22, 1995UT.} \\label{fig-q0107_lab} \\end{figure}  Observationally comparing the distribution of galaxies and quasar absorption line systems shows some of the complexity of the phenomena at work. This can be illustrated by considering Figure~\\ref{fig-q0107_lab}.  In this figure we show an image of the lines-of-sight (LOS) towards the QSOs LBQS 0107-025A and B. As discussed later in the paper, for this field we have samples of galaxies and \\lya\\ absorption line systems whose locations in this common volume of space can be compared.  In the figure we have selectively labelled illustrative cases of `associated' galaxies and absorption lines. For two of the galaxies labelled with a redshift, arrows indicate where there is a detected absorption line in both of the QSO spectra. For those absorber-galaxy pairs, the projected spatial separation and velocity difference between the galaxy and the absorber are also provided.  We consider a few of these cases below: \\begin{enumerate} \\item A bright, multi-component object lies between the two QSOs at a redshift of 0.2272. There is absorbing gas seen at nearly this same redshift in both QSO LOS, and in one of them, CIV absorption is also detected. This then could potentially be interpreted as a straightforward case of a large gaseous halo around a bright galaxy. The complex morphology of the galaxy could indicate an ongoing merger. \\item At a distance of 350 kpc to the south and east of the two QSO LOS is another bright morphologically-complex galaxy at a redshift of 0.1145. This galaxy is also close in velocity to absorption seen in both LOS. It is instructive though to compare the small visible extent of the stars in the (large) galaxy with the distance to the two QSO LOS. Smooth, undisturbed gaseous halos of this size seem improbable, although it is one of the goals of this paper to investigate this statistically. \\item Generally to the north of the two QSOs are a collection of 5 galaxies marked with redshifts near 0.20. All 5 of these galaxies are within 500 \\kms\\  (and also within a projected distance of 500 kpc) of an absorber seen in the LOS to LBQS 0107-025B. No absorption is detected in the other QSO LOS. This collection illustrates the common difficulty of picking out one individual galaxy to tie to one absorber. Although there is indeed a single `closest' galaxy, it seems rather arbitrary to claim that it is the sole cause of any absorption. \\item Finally, a collection of six galaxies, all within 500 km s$^{-1}$ of each other at a redshift around 0.24 are marked surrounding the LOS to LBQS 0107-025A. Despite their proximity to the LOS, and the case above where a group of galaxies can be associated with an absorber, no absorption is seen at this redshift in either of the two QSO LOS. While one could speculate that low column density absorbers might be fragile entities that are destroyed in a dense environment, there is no good independent evidence at present that this galaxy `group' has crossed such a threshold. \\end{enumerate}  These examples demonstrate the great variety of appearances the underlying relationship between the absorbers and galaxies presents, along even one LOS. This richness should be kept in mind when interpreting the results derived from the ensemble of data we consider in this paper.  The outline of the paper is as follows: in \\S~\\ref{sec-data} we describe the various data sets used in this paper. In \\S~\\ref{sec-anal} we analyse the data, and in \\S~\\ref{sec-conc} we present our conclusions. ",
        "conclusions": "\\label{sec-conc}  A goal of our work is an improved understanding of how gas is transformed into galaxies, and how those galaxies in turn influence the gas around them (e.g. by winds, jets or radiation).  Studies of absorbers and galaxies at high redshift have already yielded evidence for such mechanisms playing a role in the formation and evolution of galaxies (e.g. \\cite{2003ApJ...584...45A}; \\citet{2002ApJ...580..634C}; \\citet{2003A&A...407..473F}; \\citet{2003MNRAS.343L..41B} and \\citet{2003ApJ...594...75K}).  We have focussed on the second half of the history of the Universe, i.e. at redshifts less than one. This is a period in which the star formation density of the Universe is thought to be dropping away from its peak, the rapid evolution of the properties of the \\lya\\ absorbers has slowed, and when the Hubble sequence of galaxies is well established. Despite this apparent middle-aged placidity, there should remain clear evidence of the wild excesses of youth, and that is what we claim to have measured. An advantage of focussing on this (large) fraction of the history of the Universe is that the galaxy population can be studied in detail, and with confidence that a representative sample has been obtained.  There is evidence from the KP data set alone for clustering of some fraction of low redshift ($z<1$)  \\lya\\ absorbers on velocity scales $<$ 300 km/sec. This evidence is that approximately 10\\% of the KP \\lya\\ absorbers are ``resolved'' in the FOS spectra \\citep{1998ApJS..118....1J}.  At higher column densities, the Lyman limit systems contained in the KP sample all have extensive associated metal line systems.  For all the extensive metal line systems (multiple absorption lines from strong resonance line producing species, e.g. Si, O, N, etc.) for which we could check for the presence of an associated Lyman Limit system, we detect such a system. In other words, the absorption line systems that are unambiguously higher column density systems (at low redshift) are unquestionably more like the gas producing the ISM of a galaxy than like the gas one might expect in a void. These same absorbers also often have associated OVI, or NV, or other evidence of more highly ionized components i.e. multi-phase gas again, like a galaxy.  Our conclusions for the more moderate column density systems that dominate our current sample can be summarized as:  \\begin{itemize} \\item A correlation between absorbers and galaxies has been detected out to impact parameters of at least 1.5 Mpc.  \\item The strength of the absorber-galaxy correlation is weaker than the correlation between galaxies and galaxies.  \\item The velocity differences seen between galaxies and absorbers with detected CIV are typically smaller than the velocity differences seen between galaxies and absorbers seen in \\lya\\ only.  \\item The above is consistent with absorbers being a mixed population containing some \u0091pristine\u0092\\footnote{Given the results of e.g. \\citet{2001ApJ...561L.153S} on the high-z Universe that show that all gas has CIV by z of 5.5, it is debatable whether anything we see in absorption is truly `pristine'} material probably infalling for the first time along the filaments predicted by current models, and some \u0091contaminated\u0092 material produced by outflows which in general lies closer to galaxies in velocity.  \\item The above qualitative picture needs to be fleshed out by comparing the numerical strengths of the observed correlations and relationships with SPH or AMR modelling. \\end{itemize}  While we have so far only made a fairly superficial analysis of the (no-doubt complex) relationship between gas and galaxies at redshifts less than one, we feel the above results are very encouraging, and that they suggest that there is considerable room for further observation and modelling of this crucial interaction."
    },
    "0601/astro-ph0601503_arXiv.txt": {
        "abstract": "IGR J17091--3624 was discovered  in 2003 April by {\\it INTEGRAL}/IBIS during its Galactic Centre Deep Exposure programme. The source was initially detectable only in the 40--100 keV range, but after two days was also detected in the 15--40 keV range. Its flux had by then increased to 40 mCrab and 25 mCrab in the 15--40 keV and 40--100 keV  bands respectively. {\\it RXTE\\/} observed the source simultaneously on 2003 April 20, with an effective exposure of 2 ksec. We report here the spectral and temporal evolution of the source, which shows a transition between the hard and soft states. We analyse in detail the {\\it RXTE}/{\\it INTEGRAL\\/} Comptonised spectrum of the hard state as well as the JEM-X detection of a blackbody component during the source softening. Even though the source spectral behavior and time variability show a similarity with the outburst of the black-hole candidate IGR J17464--3213 (= H1743--322), observed by {\\it INTEGRAL\\/} in 2003, the nature of its compact object (BH vs.\\ NS) remains controversial. ",
        "introduction": "A new source was discovered by {\\it INTEGRAL}/IBIS during a pointed observation on 2003  April 14--15. It was designated IGR J17091--3624 based on its position of R.A.\\ (2000) $=17^{\\rm h}09.1^{\\rm m}$, Dec.\\ $=-36\\degr 24.5\\arcmin$, with an error box of $3\\arcmin$ at 90\\% confidence \\cite{Atel1}.  Initially, its flux was $\\sim$20 mCrab in the 40--100 keV energy band exhibiting a hard spectrum, while it was not detected in the 15--40 keV band with an upper limit of $\\sim$10 mCrab. A week earlier, the source did not appear in either energy band, with upper limits of $\\simeq$10 mCrab \\cite{Atel1}. During subsequent observations of the Galactic Center Deep Exposure (GCDE) on 2003 April 15--16, the source flux  increased to $\\sim$25 mCrab in the 40--100 keV band. The source was also included in the {\\it INTEGRAL\\/} IBIS/ISGRI Galactic plane survey catalogue \\cite{Bird} and classified as unknown. Observations with the Very Large Array revealed the presence of a possible radio counterpart with a steeply falling spectrum typical of synchrotron emitters \\cite{Atel3}, this was also confirmed by a subsequent observation with the Giant Metrewave Radio Telescope (Pandey et al.\\ 2006).  Immediately after the {\\it INTEGRAL\\/} discovery, an {\\it RXTE\\/} observation was performed. The analysis of the {\\it RXTE\\/} data showed a relatively stable source flux without features associated with quasi-periodic oscillations (QPOs) in the power spectrum. Based on the energy and power spectra, Lutovinov \\& Revnivtsev (2003) suggested that IGR J17091--3624 is an X-ray binary in the low/hard state and probably a black-hole system. IGR J17091--3624 was then searched for in the X-ray   catalogs and detected in the archival data of both the TTM telescope on board the KVANT module of the {\\it Mir\\/} orbital station \\cite{Atel2}, and in the {\\it BeppoSAX\\/} WFC \\cite{Atel4}. Following this, Lutovinov et al.\\ (2005) performed a preliminary study of the IBIS/ISGRI spectrum of the source, showing that its flux changed by a factor of $\\sim$2 between 2003 April and 2003 August--September. During 2003, the source spectrum was described by a simple power law with a photon index of $\\Gamma\\simeq 2.2$ in the 20--150 keV energy band, i.e., softer than the {\\it RXTE\\/} spectrum. As the source flux in the hard X-rays increased in 2004, the spectrum became harder with $\\Gamma\\simeq 1.6$. From the above investigations, IGR J17091--3624 appears as a moderately bright variable (probably a transient) source with flaring activity in 1994 October ({\\it Mir}/KVANT/TTM), 1996 September ({\\it BeppoSAX}/WFC), 2001 September ({\\it BeppoSAX}/WFC, in 't Zand et al.\\ 2003), and 2003 April ({\\it INTEGRAL}/IBIS, Kuulkers et al.\\ 2003).  \\subsection{Improvements over previous analysis}  We have performed a more accurate spectral analysis of this source compared to those previously reported, analysing all the {\\it INTEGRAL\\/} data available while the source was detectable (until the end of 2004 April). The new set of data allowed us to fit a joint {\\it RXTE}/{\\it INTEGRAL\\/} spectrum of IGR J17091--3624 in the low/hard state. The improved version of the {\\it INTEGRAL\\/} data analysis software (OSA 4.2, Goldwurm et al.\\ 2003) has allowed us to obtain {\\it INTEGRAL\\/} spectra up to 150 keV. An initial test with the newly delivered OSA 5 software has shown a substantial agreement with the former OSA 4.2 results.  We have fitted the data with detailed models, thus obtaining new information on the geometry of the source and on the accretion disk optical depth and temperature. We have also extracted a JEM-X spectrum that confirms the previous hypothesis of a transition to the soft state \\cite{Lut} and furthermore we have found suggestions of a hysteresis-like behaviour in the source spectral evolution, with the hard-to-soft state transition occurring at higher luminosity than the soft-to-hard transition. We then speculate on the source nature, even if the final picture is still not evident. ",
        "conclusions": "We have followed the evolution of an outburst of  IGR J17091--3624. Within the course of less than one year, the source went from the hard to the soft state and then back to the hard state.  Although Markoff et al.\\ (2005) have recently suggested that the hard state emission could be due to synchrotron self-Compton emission from the base of the jet, we have considered here the standard hard-state model with the dominant radiative process being thermal Comptonization of disk blackbody photons. An intriguing feature of our hard-state spectra is the rather low electron temperature, $\\sim$20 keV, while the range of temperatures typical for black-hole binaries is $\\sim$50--100 keV (e.g., Zdziarski \\& Gierli\\'nski 2004). This low temperature is obtained using both the {\\sc compps} and {\\sc comptt} models, and for both occurrences of the hard state, in the rise and decline. Such a low temperature may be a characteristic of the hard state of neutron-star low-mass X-ray binaries, although at present it is rather poorly constrained (e.g., Gierli\\'nski \\& Done 2002). On the other hand, we might have caught the source during transitionary phases, with enhanced Compton cooling due to the increasing flux of disk blackbody emission and the correspondingly lower electron temperature.  During the soft state in epoch 2, the peak of the $EF(E)$ emission is at several keV, as compared to the $\\sim$100 keV peak in the hard state. There is a high-energy tail, but its weakness does not permit us to put constraints on its origin.  In Section \\ref{spectral}, we have also compared the evolution of IGR J17091--3624 with that of another transient source observed by {\\it INTEGRAL}, IGR J17464--3213. This is a black-hole candidate discovered by {\\it HEAO 1\\/} in 1977, named then H1743--322. An outburst of that source was observed during the GCDE in 2003 \\cite{Cap}. Both sources show similar spectral evolution during outbursts of similar duration. However, as shown in Figure\\ \\ref{comparison}, the hard-state spectrum of IGR J17464--3213 is substantially softer than that of IGR J17091--3624, although both sources share the relatively low electron temperature of $\\sim$20 keV. Furthermore, the relative normalisation of the observed hard and soft state spectra is different. Curiously, the high-energy parts ($\\ga$30 keV) of both spectra of IGR J17464--3213, are virtually identical. The various amplitudes of the hard and soft state spectra are often seen in X-ray transients (e.g., Zdziarski et al.\\ 2004), and are due to the hysteretic behaviour of the accreting systems.  IGR J17091-3624 is located $10\\degr$ from  the Galactic Centre  region, where the highest concentration of LMXRBs in the Galaxy is present \\cite{White}. The model unabsorbed bolometric (0.01--500 keV) luminosities in the hard (revs.\\ 61--63) and soft state (revs.\\ 100--119), assuming a distance of 8.5 kpc and isotropy, are $L\\simeq 1\\times 10^{37}$ erg s$^{-1}$ and $L\\simeq 2\\times 10^{37}$ erg s$^{-1}$, respectively. These values are not far from those characteristic to the hard and soft states of black-hole binaries (e.g., Zdziarski \\& Gierli\\'nski 2004), although the neutron-star nature of the source cannot be excluded. On the other hand,  the {\\it RXTE\\/} power spectrum, without QPO--like features and with band-limited noise \\cite{RXTE}, commonly observed in black-hole binaries in the low/hard spectral state, as well as the probable correlation with a radio source \\cite{Pand,Atel3} may also argue for the black-hole nature of this source."
    },
    "0601/astro-ph0601053_arXiv.txt": {
        "abstract": "{}{Two inter-related peaks occur in high-frequency power spectra of X-ray lightcurves of several black-hole candidates. We further explore the idea that a non-linear resonance mechanism, operating in strong-gravity regime, is responsible for these quasi-periodic oscillations (QPOs).}{By extending the multiple-scales analysis of Rebusco, we construct two-dimensional phase-space sections, which enable us to identify different topologies governing the system and to follow evolutionary tracks of the twin peaks. This suggests that the original (Abramowicz \\& Klu\\'zniak) parametric-resonance scheme can be viewed as an ingenuous account of the QPOs model with an internal resonance.}{We show an example of internal resonance in a system with up to two critical points, and we describe a general technique that permits to treat other cases in a systematical manner. A separatrix divides the phase-space sections into regions of different topology: inside the libration region the evolutionary tracks bring the observed twin-peak frequencies to an exact rational ratio, whereas in the circulation region the observed frequencies remain off resonance. Our scheme predicts the power should cyclically be exchanged between the two oscillations. Likewise the high-frequency QPOs in neutron-star binaries, also in black-hole sources one expects, as a general property of the non-linear model, that slight detuning pushes the twin-peak frequencies out of sharp resonance.}{} ",
        "introduction": "Twin peaks occur in high-frequency ($\\sim50$--$450$~Hz) power spectra of X-ray ($\\sim2$--$60$~keV) lightcurves of several black-hole candidates (see van der Klis \\cite{kli06}; McClintock \\& Remillard \\cite{mcc06} for recent reviews of observational properties and theoretical interpretations). This transient phenomenon seems to be connected with the kilohertz quasi-periodic oscillations (QPOs) in neutron star sources, of which more examples are known (about the tens at present). The nature of black hole high-frequency QPOs remains puzzling despite variety of models proposed in the literature. In neutron-star low-mass X-ray binaries these twin QPOs are known to occur often simultaneously and they can be highly coherent ($Q\\gtrsim10^2$; e.g.\\ Barret et al.\\ \\cite{bar05}), slowly wandering in frequencies between different observations, whereas in black-hole candidates the QPO coherency appears to be lower ($Q\\sim2$--$10$) and the presence of a pair pops up only when a collection of observations is carefully analyzed.  In Abramowicz \\& Klu\\'zniak (\\cite{abr01}) and Klu\\'zniak \\& Abramowicz (\\cite{klu01}) an idea of accretion disc resonance was proposed, which naturally incorporates pairs of frequencies occurring in a ratio of small integer numbers. This scheme  predicts the observed frequency ratios in black-hole QPO sources should prefer the $3$:$2$ ratio; it also  suggests this could be understood if a non-linear coupling mechanism operates in a black-hole accretion disc, where strong-gravity effects are essential. Indeed, especially in those black-hole candidates where the high-frequency QPOs have been reported, they occur very close to the ratio of small integer numbers, $3$:$2$ in particular (Miller et al.\\ \\cite{mil01}; Strohmayer \\cite{str01}; Homan et al.\\ \\cite{hom05}; Remillard et al.\\ \\cite{rem05}; Maccarone \\& Schnittman \\cite{mac05}). Nowadays, the original account can be viewed as a na\\\"{\\i}ve model with the internal resonance. Various realizations of this scheme have been examined in terms of accretion disc/torus oscillations (e.g.\\ Abramowicz et al.\\ \\cite{abr03}; Bursa et al.\\ \\cite{bur04}; Kato \\cite{kat04}; Li \\& Narayan \\cite{li_04}; Schnittman \\& Rezzolla \\cite{sch05}; Zanotti et al.\\ \\cite{zan05}).  So far the ``right'' model has not yet been identified. However, it has been recognized that fruitful knowledge about common properties of high-frequency QPOs can be gained by investigating a very general resonance scheme, which likely governs matter near a compact accreting body. To this aim, Abramowicz et al.\\ (\\cite{abr03}) examined the epicyclic resonances in a nearly-geodesic motion in strong gravity. Rebusco (\\cite{reb04}) and Hor\\'ak (\\cite{hor04}), by employing the method of multiple scales (Nayfeh \\& Mook \\cite{nay79}), have demonstrated that the $3$:$2$ resonance is indeed the most prominent one near horizon of a central black hole. Only certain frequency combinations are allowed, depending on symmetries which the system exhibits, and only some of the allowed combinations have chance to give rise to a strong resonance.  In the present paper we pursue this approach further and we find tracks that an axially symmetric system with two degrees of freedom, near resonance, should follow in the plane of energy (of the oscillations) versus radius (where the oscillations take place). We show different topologies of the phase space in the way that closely resembles the method of disturbing function, familiar from the studies of the evolution of mean orbital elements in celestial mechanics (Kozai \\cite{koz62}; Lidov \\cite{lid62}). The analogy is very illuminating and it provides a systematic way of distinguishing topologically different states of the system. In particular, one can discriminate regions of phase space where the observed frequency ratio fluctuates around an exact rational number from those regions where this ratio remains always outside the resonance. Our model suggests that even black-hole twin QPOs should vary in frequency and they should not stay at a firmly fixed frequency ratio, albeit the expected variation is very small -- certainly less than what has been frequently reported in neutron star binaries and what can be tested with the data available at present. ",
        "conclusions": "We have discussed the resonance scheme for high-frequency QPOs via multiple-scales analysis, assuming an axisymmetric conservative system with two degrees of freedom. This approach provides a useful insight into general properties that are common to different conceivable mechanisms driving the oscillations, although it does not address the question how the observed signal is actually formed and modulated.  In our scenario, amplitudes and phases of the oscillations are mutually connected and they follow tracks in the phase space with distinct topologies. The particular form was assumed to couple the oscillation modes via the non-spherical terms in the gravitational field of a ring. We consider this to be a toy-model for rather general behaviour, which should take place in any system governed by equations of type (\\ref{eq:ms-generalfrh})--(\\ref{eq:ms-generalfth}).  We assumed the Newtonian (or the pseudo-Newtonian) description of the central gravitational field with a perturbation by an aligned ring as an example. The adopted form is not essential for general conclusions. In fact, equations (\\ref{eq:ms-generalfrh})--(\\ref{eq:ms-generalfth}) cover also the nearly-geodesic motion around a Schwarzschild black hole. Compared with the pseudo-Newtonian case, general relativity does not bring qualitatively new features, as long as the system is conservative (additional terms will arise in the expansions, which then translate to slightly different value of the resonance radius and to different duration of time intervals in physical units). A natural question arises whether the gravitational field of a rotating black hole could provide the perturbation required for the internal resonance in a surrounding disc. We considered this possibility, however, it is  unlikely that Kerr metric could by itself suffice: in the weak-field limit non-spherical terms seem to be incapable of creating separatrices in the phase-space sections discussed above, whereas in the full (exact, vacuum) Kerr metric the special mathematical properties of the spacetime ensure the integrability of the geodesic motion, and hence prevent the occurrence of internal resonances. Therefore, the problem of a specific mechanism launching and maintaining the oscillations remains unanswered.  Various options for the generalisation of our scheme could be motivated by papers of other authors who proposed specific models including non-gravitational forces (see P\\'etri \\cite{pet05} for a recent exposition of the problem and for references). As a next step towards an astrophysically realistic scheme, one should take dissipative and non-potential forces into account, as well as non-axisymmetric perturbations. These will allow our system to migrate across contours in the phase-plane and to undergo transitions when crossing separatrices. Such additional terms could also supplement the influence of external forcing and initiate the oscillations of the system. The internal resonance would then define the actual frequencies that are excited; this way the  strong gravity unmasks itself."
    },
    "0601/astro-ph0601579_arXiv.txt": {
        "abstract": "The number of  $z \\sim 1$ damped Lyman alpha systems (DLAs, log \\nhi\\ $\\ge$ 20.3) per unit redshift is approximately 0.1, making them relatively rare objects. Large, blind QSO surveys for low redshift DLAs are therefore an expensive prospect for space-borne UV telescopes.  Increasing the efficiency of these surveys by pre-selecting DLA candidates based on the equivalent widths of metal absorption lines has previously been a successful strategy.  However, the success rate of DLA identification is still only $\\sim$ 35\\% when simple equivalent width cut-offs are applied, the majority of systems having $19.0 <$ log \\nhi\\ $<20.3$. Here we propose a new way to pre-select DLA candidates. Our technique requires high-to-moderate resolution spectroscopy of the Mg~II $\\lambda$ 2796 transition, which is easily accessible from the ground for $0.2 \\sol z \\sol 2.4$.  We define the \\dind, the ratio of the line's equivalent width to velocity spread and measure this quantity for 19 DLAs and 8 sub-DLAs in archival spectra obtained with echelle spectrographs. For the majority of absorbers, there is a clear distinction between the \\dind\\ of DLAs compared with sub-DLAs (Kolmogorov-Smirnov probability = 0.8 \\%).  Based on this pilot data sample, we find that the \\dind\\ can select DLAs with a success rate of up to 90\\%, an increase in selection efficiency by a factor of 2.5 compared with a simple equivalent width cut.  We test the applicability of the \\dind\\ at lower resolution and find that it remains a good discriminant of DLAs for FWHM $\\sol$ 1.5 \\AA. However, the recommended \\dind\\ cut-off between DLAs and sub-DLAs decreases with poorer resolution and we tabulate the appropriate \\dind\\ values that should be used with spectra of different resolutions. ",
        "introduction": "The compilation of large spectroscopic datasets, such as the Sloan Digital Sky Survey (SDSS), has proven a major boon for the study of QSO absorption lines.  These surveys have yielded many hundreds of damped Lyman alpha systems (DLAs, absorption systems with \\nhi\\ $ \\ge2 \\times 10^{20}$ \\cm) which can be used to study the evolution of the galactic neutral gas content at $z>1.6$ (e.g. Prochaska, Herbert-Fort \\& Wolfe 2005).  In the local universe, several wide field 21cm surveys have characterised the $z=0$ distribution of neutral hydrogen gas (e.g. Zwaan et al. 2005; Rosenberg \\& Schneider 2003; Ryan-Weber, Webster \\& Staveley-Smith 2003). Bridging these extremes, over a redshift range that is emerging as a crucial time in galaxy formation, is a high priority for absorption line studies.  In practice, however, the low redshift universe has been challenging to survey. The \\lya\\ line is shifted into the UV at $z<1.5$ where space telescopes are required.  Given the limited resources, large-scale, blind surveys for low redshift DLAs are expensive and largely unpalatable to time assignment committees.  For this reason, pre-selection of low $z$ DLA candidates based on low resolution, ground-based spectra has become an important first step in UV surveys (Rao \\& Turnshek 2000; Rao, Turnshek \\& Nestor 2006).  The equivalent widths (EWs) of resonance lines with red rest wavelengths, such as Mg~II $\\lambda \\lambda$ 2796, 2803 and Fe~II $\\lambda$ 2600, are easily measured from the ground in systems with \\nhi\\ $\\sog 10^{19}$ \\cm\\ and $0.2 \\sol z \\sol 2.4$. Rao et al. (2006) showed that $\\sim$ 35\\% of absorbers with rest frame EW(Mg~II $\\lambda$ 2796, Fe~II $\\lambda$ 2600) $>$ 0.5 \\AA\\ have column densities above the canonical DLA limit, \\nhi $\\ge 2 \\times 10^{20}$ \\cm.  Although absorbers below this limit (sub-DLAs) make a non-negligible contribution to the total neutral gas content of the universe at high redshifts (P\\'eroux et al. 2005; Prochaska et al. 2005), they are usually excluded from statistical samples.  Therefore, although EW pre-selection has been a pivotal strategem of recent Hubble Space Telescope (HST) surveys, the 35\\% `success rate' of metal line pre-selection could be considered low. Moreover, whilst a confirmation probability can be applied in a statistical sense, there is no way of knowing which individual Mg~II absorber is most likely to be a DLA, since there is no correlation between Mg~II EW and \\nhi.  For example, a DLA confirmation rate of 35\\% applied to the $0.6 < z < 1.7$ Mg~II system survey of Ellison et al. (2004a) yields a DLA number density of $n(z)=0.11^{+0.08}_{-0.05}$ (Ellison 2005).  However, UV follow-up of the 14 high EW Mg~II systems in the Ellison et al. sample is still required in order to confirm the identity of the expected 5 DLAs.  In this paper, we investigate whether ground-based spectra of metal lines at high resolution can provide additional information to help discriminate between DLAs and sub-DLAs. Our objective is to identify a quantitative measure that will efficiently separate DLAs and sub-DLAs with a success rate $>>$ 35\\%.  If the inclusion of sub-DLAs (i.e. false positives) can be minimised, the obvious benefit will be that UV follow-up for DLA confirmation and \\nhi\\ determination will be far more efficient. Moreover, if the pre-selection is complete for high \\nhi\\ DLAs, any quantity weighted by \\nhi, such as the total neutral gas content, will be robust, even if a few lower \\nhi\\ systems are missed in the pre-selection process. Finally, a high success rate DLA discriminant could be applied to surveys of Mg~II absorption systems to directly infer the DLA number density to a relatively high degree of accuracy, with no need for UV observations. ",
        "conclusions": "We have presented a new technique for discriminating between low and high \\nhi\\ absorbers based solely on high resolution spectroscopy of the Mg~II $\\lambda$ 2796 line.  Qualitatively, DLAs tend to be characterised by broad, heavily saturated Mg~II, whereas sub-DLAs often have considerable residual flux over their extent.  Quantitatively, we define the \\dind, the ratio of Mg~II $\\lambda$ 2796 EW to velocity spread.  DLAs tend to have high $D$-indices, with a cut of $D > 6.3 $ yielding a success rate of 90\\%. Previously, identification of DLA candidates based on EW cuts from low resolution spectra had a DLA confirmation rate of only 35\\%. A larger sample of absorbers which is more representative of the column density distribution function (i.e. with a higher fraction of sub-DLAs than DLAs), may have a lower success rate depending on the performance of the \\dind\\ at low \\nhi.  This is difficult to judge from the current sample since only 8 of our absorbers have \\nhi\\ $<$ 20.3.  Both of the sub-DLAs with high $D$ can be marked out as `unusual' in some way; Q0058+019 is known to have a very small impact parameter to its parent galaxy and is rare in exhibiting solar metallicity (Pettini et al. 2000) and Q0957+561B is a lensed QSO. Whilst we do not suggest these factors as exemptions or explanations, it once again highlights the need for better sub-DLA statistics. In addition, all DLAs with \\nhi\\ $>$ 20.5 in our pilot sample have $D > 6.3$, so that pre-selection with the \\dind\\ includes the systems which dominate the neutral gas content.  Although the column density distribution function evolves as a function of redshift (e.g. Prochaska, Herbert-Fort \\& Wolfe 2005; Rao, Turnshek \\& Nestor 2006; Zwaan et al. 2005) it can be approximated to single power law $f(N) \\propto N^{\\alpha}$. For an index $\\alpha = - 1.5$, less than 10\\% of the HI gas is in systems with log \\nhi\\ $\\le$ 20.5.  Therefore, although using the \\dind\\ may miss some DLAs with lower column densities, for the purposes of pre-selecting DLAs to estimate the total neutral gas in DLAs, it is quite robust.  Similarly, for any property that is weighted by \\nhi, such as is common practice with metallicity (e.g. Pettini et al. 1999; Kulkarni et al 2005), the DLAs missed by \\dind\\ pre-selection will have a relatively minor effect.  In order for the \\dind\\ to be a successful discriminant of DLAs from sub-DLAs, a spectral resolution of at least 1.5 \\AA\\ FWHM (at the wavelength of Mg~II $\\lambda$ 2796) is required.  Based on simulations which degrade the spectral resolution of our archival data, we have tabulated the \\dind\\ which yields the highest DLA return rate for a given instrumental FWHM.  The distinction between DLAs and sub-DLAs encapsulated by the \\dind\\ is one of spectral morphology. The qualitative tendency for sub-DLAs to have more extended velocities for their EWs with a larger number of distinct components was previously noted by Churchill et al. (2000). Rao \\& Turnshek (2000) also discussed the difference in spectral morphology between high and low \\nhi\\ absorbers noting that `the DLA systems generally appear to exhibit somewhat greater saturation.' The physical reason for this difference, which leads to the difference in $D$-indices in the two populations, is still unclear, but we speculate that it may be driven by impact parameter (see also Rao \\& Turnshek 2000). In their study of a gravitationally lensed QSO, Ellison et al. (2004b) found that the coherence scale of weak (EW $<$ 0.3 \\AA) Mg~II absorbers is $\\sim$ 2 \\hkpc\\ ($\\Omega_{\\Lambda}$ = 0.7, $\\Omega_{\\rm M}$ = 0.3), whilst stronger absorbers have much larger sizes. Ellison et al. discussed a scenario in which galactic halos were composed of numerous Mg~II absorbing `clouds' whose filling factor was larger near the centres of galaxies.  In this picture, a sightline intersecting at small impact parameters would show absorption from many Mg~II regions, but within a relatively narrow velocity range. Conversely, sightlines at larger impact parameters would exhibit more disjointed absorption complexes and may have much wider velocity spreads (depending on the kinematics of the individual galaxy). This would  lead to an anti-correlation between \\dind\\ and impact parameter.  We can test this prediction with the galaxy--Mg~II absorber sample of Churchill et al. (2000) for which galaxies have been identified for 16 Mg~II absorbers (albeit with EWs typically lower than those considered in this paper). Although Churchill et al. (2000) define velocity spread differently from us, the ratio with EW (the analog to our \\dind) does anti-correlate with impact parameter. If the HI column density also falls off with impact parameter, this would lead to a correlation between \\dind\\ and \\nhi\\ and qualitatively explain the distribution of points in Figure \\ref{dindex}.  We conclude that, based on these preliminary investigations, a combination of EW and velocity spread may be a powerful way to pre-select DLA candidates. Although more data are required to test the robustness and selection biases of this technique it could potentially increase the yield of pre-selected DLA candidates by a factor of 2--3 over current methods. With more data available, it will also be possible to extend the investigation of the \\dind\\ to other transitions which may be suitable for DLA identification, such as Fe~II transitions longwards of $\\sim$ 2300 \\AA."
    },
    "0601/astro-ph0601323_arXiv.txt": {
        "abstract": "We present an analysis of ``bistability'' in gas-phase chemical models of dark interstellar clouds. We identify the chemical mechanisms that allow high- and low-ionization solutions to the chemical rate-equations to coexist. We derive simple analytic scaling relations for the gas densities and ionization rates for which the chemistry becomes bistable. We explain why bistability is sensitive to the H$_3^+$ dissociative recombination rate coefficient, and why it is damped by gas-grain neutralization. ",
        "introduction": "\\label{introduction}  In this paper we reexamine and analyze the phenomenon of ``bistability'' that occurs in chemical models of dark interstellar clouds. Bistability was first discussed by Le Bourlot et al.~(1993), who found that multiple solutions to the chemical rate-equations sometimes appear in numerical computations of the gas-phase abundances of atomic and molecular species for steady-state dark cloud conditions. When bistability occurs, three sets, rather than a single unique set, of chemical abundances are predicted for identical input cloud parameters such as the total gas density, the cosmic-ray (or X-ray) ionization rate, and the gas-phase supply of heavy elements. Under such circumstances, the solution that is converged to in a numerical simulation, and perhaps, the actual steady-state of a real physical system, depend on the initial conditions and history of the gas.  The subject of bistability begins with Pineau des Forets et al.~(1992) who studied how the computed chemical states of dark clouds depend on the gas density and ionization rate. They found, as would be expected, that for a given ionization rate the fractional ionization of the gas decreases roughly as the square-root of the gas density. However, they also found that the ionization fraction and associated chemical composition changes abruptly, or quasi-discontinuously, at a specific critical transition density. Below the critical density the clouds are in what Pineau des Forets et al.~referred to as a ``high-ionization'' phase (HIP), characterized by high abundances of atomic ions. For densities greater than the critical density the clouds are in a ``low-ionization'' phase (LIP), with high abundances of molecular ions. Pineau des Forets et al.~also identified a chemical instability that, they argued, is responsible for the abrupt transition that they encountered in their calculations. Bistability was subsequently discovered by Le Bourlot et al.~(1993) who found that instead of a phase change at a specific critical density, the LIP and HIP ``branches'' can sometimes overlap for a narrow range of densities. A third (unstable) branch also appears, intermediate between the LIP and HIP solutions.  Further investigations have explored the effects of the assumed gas phase elemental abundances, ionization rates, gas temperatures, chemical reaction rate coefficients, and gas-grain interactions, on the existence and character of the multiple solutions (Le Bourlot et al.~1995ab; Shalabiea \\& Greenberg 1995; Lee et al.~1998; Pineau des Forets \\& Roueff 2000; Viti et al.~2001; Charnley \\& Markwick 2003).  These studies have shown that bistability is indeed related to the transition from the HIP to LIP (or vice versa), and that it is sensitive to quantities such as the gas-phase heavy-element abundances (Le Bourlot et al.~1995; Viti et al.~2001), and the adopted chemical network (Lee et al.~1998), including specifically the dissociative recombination rate coefficient of H$_3^+$ (Pineau des Forets \\& Roueff 2000). In this paper we explain why the HIP and LIP can overlap, and we identify the chemical mechanisms that allow the possibility of multiple solutions.  We will show that bistability can be simply understood and analyzed using just a handful of reactions. In a nutshell, when bistability occurs atomic ions (usually S$^+$ ions) are the dominant positive charge carriers for both HIP and LIP conditions. Two modes of neutralization are then available to the gas. In the HIP, neutralization proceeds by slow radiative recombination. In the LIP, neutralization occurs by rapid dissociative recombination, via the formation of intermediate molecular ions in reactions with O$_2$ molecules (e.g., ${\\rm S^+ + O_2 \\rightarrow SO^+ + O}$). Bistability occurs when a low O$_2$ abundance in the HIP can be maintained up to maximal gas densities that are {\\it greater} than the minimal densities above which a large O$_2$ density is maintained in the LIP. We will show by means of a simple analysis why this is possible. The explanation involves a chemical instability, of a type first introduced by Pineau des Forets et al.~(1992), but not given sufficient attention in subsequent studies.  Bistability may be of primarily theoretical interest since, as we will discuss, it is a purely gas-phase chemical effect. If realistic gas-dust neutralization processes are included, the multiple solutions are damped or eliminated entirely (cf.~Shalabiea \\& Greenberg 1995). We will illustrate this by considering the effect of grain-assisted ion-electron recombination (Weingartner \\& Draine 2001) on the chemistry. Nevertheless, because bistability has been the subject of considerable discussion in the literature, we feel that a fresh analysis of how, why, and when it occurs is warranted.  The plan of our paper is as follows. In \\S 2 we present results of an illustrative numerical computation in which bistability occurs. In \\S 3 we discuss the key chemical instability that drives the steady-state solutions to either the LIP or HIP. In \\S 4 we derive simple analytic formulae that show how the range of gas densities and ionization rates for which the gas becomes bistable depends on the elemental abundances and several important chemical rate coefficients. In \\S 5 we analyze the sensitivity of bistability to the H$_3^+$ dissociative recombination rate coefficient. In \\S 6 we explain why gas-grain neutralization processes damp and eliminate bistability. We summarize in \\S 7. ",
        "conclusions": "The non-linearity of equations (1)-(4) is a mathematical prerequisite for bistability and the appearance of multiple solutions (Le Bourlot et al.~1993, 1995; Lee et al.~1998; Pineau des Forets \\& Roueff 2000; Charnley \\& Markwick 2003). However, the purpose of our paper has been to identify and analyze the chemical mechanisms underlying the bistability phenomenon. We have shown that a simple chemical instability, involving atomic ions (typically S$^+$), O$_2$ molecules, and H$_3^+$ ions, allows multiple solutions to appear near the transition points from low-density high-ionization (HIP) to high-density low-ionization (LIP) conditions in the gas-phase chemistry.  The instability drives the O$_2$ density to either small (HIP) or large (LIP) values. Bistability occurs when a low O$_2$ abundance in the HIP can be maintained up to maximal densities that are greater than the minimal densities above which a large O$_2$ density is maintained in the LIP. Two modes of neutralization are then simultaneously available to the gas. The slow (HIP) mode is radiative recombination (usually ${\\rm S^+ + e \\rightarrow S + \\nu}$). The fast (LIP) mode is formation of molecular ions followed by dissociative recombination (usually ${\\rm S^+ + O_2 \\rightarrow SO^+ + O}$; ${\\rm SO^+ + e \\rightarrow S + O}$).  We have derived simple analytic scaling relations for the range of gas densities and ionization rates for which bistability occurs. Our formulae show how the bistable range depends on the assumed gas-phase abundances, and several reaction rate-coefficients, including the dissociative recombination rate coefficient of H$_3^+$. Bistability is eliminated when the two modes of gas-phase neutralization are superseded by gas-grain processes such as grain assisted recombination of the atomic ions. This likely happens for the range of electron and O$_2$ densities where bistability typically occurs in purely gas-phase models. We conclude that the bistability phenomenon probably does not occur in realistic dusty interstellar clouds, although it remains of theoretical interest."
    },
    "0601/astro-ph0601609_arXiv.txt": {
        "abstract": "The binary star system $\\gamma$ Cephei is unusual in that it harbours a stable giant planet around the larger star at a distance only about a tenth of that of the stellar separation. Numerical simulations are carried out into the stability of test particles in the system. This provides possible locations for additional planets and asteroids. To this end, the region interior to the planet is investigated in detail and found to permit structured belts of particles. The region between the planet and the secondary star however shows almost no stability. The existence of an Edgeworth-Kuiper belt analogue is found to be a possibility beyond $65$ au from the barycentre of the system, although it shows almost no structural features. Finally, the region around the secondary star is studied for the first time. Here, a zone of stability is seen out to $1.5$ au for a range of inclinations. In addition, a ten Jupiter-mass planet is shown to remain stable about this smaller star, with the habitability and observational properties of such an object being discussed. ",
        "introduction": "$\\gamma$ Cephei is one of the closest separation binary systems that contains a planet. The primary and secondary stars are separated by 18.5 au. The planet (designated $\\gamma$ Cep Ab) has a minimum mass of 1.7 $\\MJ$ and follows an eccentric orbit about the larger star alone. The observational parameters of $\\gamma$ Cephei, as determined by \\citet{Ha03}, are summarised in~Table~\\ref{tab:orbels}.  Although the data are reasonably reliable, there is some uncertainty regarding the masses of the components. The inclination to the line of sight of a spectroscopic binary is usually indeterminate. This means that the mass of the primary is generally reckoned from spectral type. Then, the mass function (e.g., Smart 1977) permits a minimum mass to be assigned to the secondary, albeit crudely (Griffin, Carquillat \\& Ginestet 2002).  For $\\gamma$ Cephei, the mass of the primary is $\\approx 1.57 M_{\\odot}$ from photospheric modelling \\citep{Fu04}. The most recent determination from stellar evolution models is $\\approx 1.7 M_{\\odot}$~\\citep{Af05}. The mass of the secondary is given by \\citet{Dv03} as $\\approx 0.4 M_{\\odot}$. However, assuming the original value of the primary's mass of 1.57 $M_{\\odot}$, a minimum mass can be derived as $\\approx 0.34$ $M_{\\odot}$ from the velocity amplitude fitted by \\citet{Ha03}.  There have been a number of studies of the dynamics of the $\\gamma$ Cephei system to date. Foremost is the numerical investigation of \\citet{Dv03}, which does use an earlier, and slightly different, set of orbital parameters (see Table \\ref{tab:altorbels}). They used Burlisch-Stoer and Fast Liapunov Indicator methods to show that the $\\gamma$ Cephei system is stable over Myr time-scales. They carried out test particle integrations as a guide to the possible existence of further planets.  The results show a stable inner region between 0.5 and 0.8 au and then an additional stable zone at low inclination around 1 au. \\citet{Dv03} noted that this coincides with the 3:1 mean motion resonance. This resonance is stable in the $\\gamma$ Cephei system, but is unstable in the Solar system. \\citet{Dv03} conclude that Earth-mass planets (up to 90 $\\ME$) can exist in this region which is fortunately just on the edge of the habitable zone \\footnote{The habitable zone has boundaries which depend on assumptions regarding stellar luminosity and effective temperature, as well as the climate model adopted for the hypothetical habitable planet.}. An extension to this work by \\citet{Dv04} showed that the planet's eccentricity is less important than that of the two stars for the stability of a second planet.  \\citet{Ha05} also carried out numerical studies of the system's stability using a Burlisch-Stoer integrator, for a range of possible configurations of the planet's and binary's semimajor axes, eccentricities and inclinations, confirming and extending the results of \\citet{Dv03}.  Here, we use both numerical simulations and analytic calculations to study zones of stability as possible locations of additional planetary companions to either star and asteroid or Edgeworth-Kuiper belt analogues.  Edgeworth-Kuiper belts are of particular interest in binary systems, as they may be being observed (indirectly) in exosystems such as $\\tau$ Ceti and $\\eta$ Corvi (Greaves et al. 2004; Wyatt et al. 2005). The results fall into three categories: possible planetary companions and asteroid belts centred on the primary star (\\S 2), planetary companions around the secondary star (\\S 3) and finally possible Edgeworth-Kuiper belt about both stars (\\S 4).  \\begin{table*} \\centerline{ \\scriptsize \\begin{tabular}{|l|r@{ $\\pm$ }l|r@{ $\\pm$ }l|r@{ $\\pm$ }l|} \\hline Name                                 &\\multicolumn{2}{|c|}{$\\gamma$ Cep A}      & \\multicolumn{2}{|c|}{$\\gamma$ Cep B} & \\multicolumn{2}{|c|}{$\\gamma$ Cep Ab } \\\\ \\hline Class                                &\\multicolumn{2}{|c|}{K1IV sub-giant star} &  \\multicolumn{2}{|c|}{M dwarf star } & \\multicolumn{2}{c}{Planet}\\\\ Mass                                 & 1.59 & 0.12 $M_{\\odot}$ &  \\multicolumn{2}{|c|}{0.4 $M_{\\odot}$} & 1.7        & 0.4 $M_{J}$\\\\ Period (days)                        &\\multicolumn{2}{|c|}{--} & 20750.6579 & 1568.6 & 905.574    & 3.08\\\\ Semimajor axis (au)                 &\\multicolumn{2}{|c|}{--} & 18.5       & 1.1 & 2.13       & 0.05\\\\ Eccentricity                         &\\multicolumn{2}{|c|}{--}                  & 0.361      & 0.023 & 0.12       & 0.05\\\\ Longitude of periastron ($^{\\circ}$) &\\multicolumn{2}{|c|}{--} & 158.76     & 1.2  & 49.6       & 25.6\\\\ Time of periastron passage (JD)      &\\multicolumn{2}{|c|}{--} & 2448429.03  & 27.0 & 2453121.925 & 66.9\\\\ \\hline \\end{tabular}} \\large \\caption{Best fit orbital parameters for the $\\gamma$ Cephei system from \\citet{Ha03}. Mass of star B is from \\citet{Dv03}.} \\label{tab:orbels} \\end{table*}  \\begin{table} \\centerline{ \\scriptsize \\begin{tabular}{|l|c|c|c|} \\hline Name                 & $\\gamma$ Cep A      & $\\gamma$ Cep B & $\\gamma$ Cep Ab\\\\ \\hline Class                & K1IV sub-giant star & M dwarf star   & Planet\\\\ Mass                 & 1.6$M_{\\odot}$      & 0.4$M_{\\odot}$ & 1.76$M_J$\\\\ Period (days)        & --                  & 25567.5        & 902.2 \\\\ Semimajor axis (au) & --                  & 21.36          & 2.15\\\\ Eccentricity         & --                  & 0.44           & 0.209\\\\ \\hline \\end{tabular}} \\large \\caption{Orbital parameters for the $\\gamma$ Cephei system used by \\citet{Dv03}.} \\label{tab:altorbels} \\end{table} ",
        "conclusions": "In this paper, we have carried out a suite of test particle integrations for the $\\gamma$ Cephei system, as a guide to possible locations for additional planets. For test particles in the plane of the binary and planet, there are three zones of stability for 1 Myr timespans at least. These are [1] the region interior to the planet, which is stable within the bands $1.2$ and $1.3$ au, $1.05$ to $1.15$ au, and $0.75$ to at least $0.5$ au, [2] the region around star B from $1.5$ au into at least $0.5$ au and [3] the region around both stars extending out from about $65$ au. These can be used to constrain possible locations of additional planetary companions.  The region interior to the planet has a complicated structure.  Low order resonances, such as the 4:1, 3:1 and 5:2, mark transitions between stable and unstable regimes. There is a secular resonance located at $0.8$ au from the primary that also plays a role in the dynamics of the test particles.  The results for this region match up well with the previous work of Dvorak et al. (2003) despite the differences in the parameters of the system and the methods used. This implies that the slight improvements in the observations of the system have not significantly altered its dynamical characteristics.  One difference though is that, unlike Dvorak et al. (2003), we find that test particles close to the 3:1 resonance are unstable, just as for the asteroid belt in the Solar System. In agreement with Dvorak et al. (2003), we find that small planets interior to $\\gamma$ Cephei Ab are long-lived.  Another as yet unconsidered possibility seen from the results here is the existence of an asteroid belt interior to the giant planet. However, the stability seen for test particles in this region disappears when $\\gamma$ Cephei Ab is inclined in the same plane as the test particles. Now, the test particles are rapidly driven to inclinations which are subject to the Kozai instability.  The region interior to the binary and exterior to the planet is unpromising. It seems unlikely that any prograde planet or asteroid can remain stable between $\\gamma$ Cephei Ab and the binary, although retrograde test particles in this region are more long-lived.  The region around the secondary is stable out to $1.5$ au for test particles. We have shown that planets up to 10 Jupiter-masses in circular orbits at $0.5$ au can survive for at least 1 Myr.  This seems to be a promising place for additional planets to reside, in a similar manner to satellites about a planet around a single star.  In the final region, that exterior to the binary, test particles are also stable. Although planets would be able to reside out here, the existence of an Edgeworth-Kuiper belt structure is also a possibility. Once again the retrograde particles are more stable than those that are prograde. Whilst in every region studied this is true, such objects are perhaps unlikely, if the origin of the system was a common disk.  If particles are captured from elsewhere, then this may become a possibility.  Retrograde asteroids and comets certainly exist in the Solar system.  Since the survival of additional planets in the system has been shown to be a possibility, it is interesting to consider their habitability. There are a number of estimates of the habitable zone around star A in the literature. For example, Dvorak et al. (2003) place it at $1.0$ to $2.2$ au, in which case habitable planets could exist at the very edge of the zone.  However, Haghighipour (2005) places the habitable zone at $3.1$ to $3.8$ au, while Jones et al. (2005) place it at $2.07$ to $4.17$ au.  These different locations reflect differences in the criteria for the location of the habitable zone or the assumed stellar luminosity and effective temperature.  The zone $2.07$ to $4.17$ au is unstable for all except retrograde test particles. So, this would permit the possibility of a habitable planet only if retrograde. The detection of any such small planet is challenging, given the small (of the order of a few ms$^{-1}$) radial velocity signatures that they cause. In addition, a retrograde planet appears in radial velocity curves merely as a prograde one with a $180^\\circ$ phase difference. Although retrograde objects -- especially planets -- are thought to be unlikely, it has been shown that giant planets in binary systems can end up with retrograde orbits after close encounters within a planetary system (Marzari et al. 2005).  A more interesting possibility is that of a habitable planet around star B, since it has been shown that planets can survive here. Theories of planetary formation do not preclude this possibility (e.g., Armitage, Clarke \\& Palla 2003). The habitable zone for an M dwarf extends out to about 0.3 au (Kasting, Whitmire \\& Reynolds 1993), which is within the stable zone found here. The large primary star nearby might also act as a `shield' from comets and asteroids similar to Jupiter for the Earth. As seen in the results in Sections~\\ref{sec:b} and \\ref{sec:kb}, very few test particles collide with the secondary star once the region interior to the binary has been rapidly cleared. Edgeworth-Kuiper belt objects that are perturbed into the inner regions of the system are also likely to suffer encounters with the primary, thus leaving the secondary and its environs largely unscathed.  The detection of such a habitable planet around star B has, unfortunately, been shown to be virtually impossible from the radial velocity signature of star A alone. Should the secondary be resolved this would change, and any companion giant planet would be easily detectable.  The region around star B, none the less, remains as a promising place for the existence of a habitable planet in this system."
    },
    "0601/astro-ph0601101_arXiv.txt": {
        "abstract": "{We  report  about a  short  flare from  the blazar  NRAO  530 occurred on  $17$ February $2004$ and detected  serendipitously by the IBIS/ISGRI  detector  on  board  \\emph{INTEGRAL}. In  the  $20-40$~keV energy  range, the  source,  that is  otherwise  below the  detection limit, is detected at a level of $\\approx 2\\times 10^{-10}$~erg~cm$^{-2}$~s$^{-1}$ during a time interval of less than $2000$~s, which is about a factor 2 above the detection threshold.  At  other wavelengths, only nearly-simultaneous radio data are available (1 observation  at $2$~cm on $11$ February $2004$), indicating  a moderate  increase  of the  polarization. This  appears to  be the shortest  time variability episode ever detected in a high luminosity blazar  at hard X-rays, unless the blazar is contaminated by the presence of an unknown unresolved rapidly varying source. ",
        "introduction": "Blazars are, among active galactic nuclei (AGN), the most luminous and most dramatically  variable.  These extreme properties  are likely due to relativistic  aberration in a  kilo-parsec jet oriented at  a small angle with  respect to our line  of sight, where plasma  moving with a Lorentz  factor  of $\\sim$10-20  causes  radiation  boosting and  time foreshortening (see reviews by  Urry \\& Padovani 1995,  Wagner \\& Witzel 1995, Ulrich et al. 1997).  In these sources, the maximum power output  and  the  largest   variability  amplitudes  are  observed  at high-energies,  from  X-   to  $\\gamma-$rays,  therefore  blazars  are interesting targets  for monitoring with  satellites like \\emph{INTEGRAL} (Pian  et  al.   1999).  During  its first  three  years  of  activity \\emph{INTEGRAL}  detected two blazars  in outburst,  with observations triggered  by ground  telescopes  (S5 $0716+714$  Pian  et al.   2005, 3C$454.3$ Foschini et al.  2005, Pian et al., in preparation), and one in  high state  (S5 $0836+710$),  serendipitously detected  during the observation of  S5 $0716+714$ (Pian  et al. 2005).  Indeed,  thanks to the large field-of-view of  the IBIS instrument onboard INTEGRAL, many blazars, possibly in active  state, may be detected serendipitously by the satellite  during its  long pointings.  In  order to  exploit this advantage, we have  systematically searched the \\emph{INTEGRAL} public archive  for possible detection  of flaring  blazars.  We  report here about the interesting detection  of a rapid flare apparently associated with the blazar NRAO~$530$.  NRAO~$530$ ($z=0.902$) is a  blazar belonging to the optically violent variable  (OVV) quasar subclass,  with emission  lines in  the optical spectrum, particularly H$\\beta$ with  a rest-frame equivalent width of $16\\AA$, smaller  than  usually observed in OVVs  (Junkkarinen 1984; cf also  Veron-Cetty \\&  Veron 2000).   It exhibits  strong  outbursts in radio (Bower  et al. 1997), optical  (up to $3$  magnitudes in $1977$, Pollock  et  al. 1979,  Webb  et  al.  1988), X-ray (\\emph{HEAO-1}, Marscher et al. 1979), and  $\\gamma-$ray  bands (EGRET/CGRO, Mukherjee  et al. 1997), with  significant emission above $1$~GeV  (Lamb \\&  Macomb  1997). The flare detected by the IBIS/ISGRI instrument on board \\emph{INTEGRAL} occurred  in February~$2004$ with a timescale ($\\approx 2000$~s) much shorter than any variability reported before.  In  the following,  we present  in detail the  \\emph{INTEGRAL} data analysis and results (Section 2),  data obtained quasisimultaneously in the radio (Section 3) and a discussion of the relevance of the detected flare for blazars.  \\begin{figure*}[!ht] \\centering \\includegraphics[scale=0.45,angle=270]{4804_f1.ps} \\caption{IBIS/ISGRI significance map in the $20-40$~keV energy band for the pointing before, during, and after the detection of NRAO $530$.} \\label{isgrimap} \\end{figure*} ",
        "conclusions": "Hard X-ray emission from the blazar NRAO~$530$ was detected serendipitously by \\emph{INTEGRAL} with IBIS/ISGRI during a long exposure of the Galactic Center. During a time interval lasting $\\approx 2000$~s the source flux in the $20-40$ keV range rose above the detection threshold by about a factor of $\\sim 2$ and then faded again.  The  spectral  energy  distribution  of  NRAO~$530$  constructed  with historical data and with  the present \\emph{INTEGRAL} flux is reported in Fig.~\\ref{fig2}. Interestingly, although very high, the IBIS/ISGRI point is consistent with the extrapolation to lower energies of the average spectrum observed in the EGRET energy range. The variability amplitude measured by EGRET was a factor $6$ and rather erratic, still  on  timescales of weeks (Mukherjee et al. 1997). On the other hand, timescales as short as implied here would not have been accessible for EGRET due to the low counting rate. The hard X-ray flux measured during the flare for any reasonable spectral shape would imply a much larger X-ray flux than observed previously. Large variations have been observed from this object also in the optical (more than a factor 10 over timescales of years), thus overall the amplitude of variability is not unprecedented.  What is truly exceptional in this event is the short timescale, less than an hour in an object of extremely high luminosity. This may cast doubts on the association with NRAO~$530$. The event could then be attributed to some still unknown galactic source inside the $3'$ radius circle of positional uncertainty in the ISGRI detection. In this case, only future satellites, like SIMBOL-X (Ferrando et al. 2005), will have the necessary spatial resolution at hard X-rays to confirm the association with NRAO~$530$. For the moment, with the present data, instrument performances, and catalogs, it is interesting to consider the possibility that this highly significant variation is associated with the blazar.  Rapid flares, with timescales of thousands of seconds in the X-rays, are often observed in BL Lac sources, especially in the TeV sources (e.g. Brinkmann et al. 2005), but they are rarely seen in blazars belonging to the quasar class, where typical timescales, even in hard X-ray--$\\gamma$-ray band, are in the range of several hours--days (e.g. Wagner \\& Witzel 1995). It is worth noting that FSRQ are generally faint in X-rays and therefore the above assertion could be biased in the past by the lack of high-energy astrophysics satellites with the necessary sensitivity. The present and the next generation of X- and $\\gamma-$ray satellites should remove this bias.  The interpretation of such a rapid flare in the context of the standard synchrotron-inverse Compton models for blazars (e.g. Ghisellini et al. 1998) requires quite extreme conditions for the emitting source. A flare lasting less than $\\Delta t=3600$~s must be emitted by a region with a size $R < 1.1\\times 10^{14} \\delta$ cm, where $\\delta$ is the Doppler factor. Adopting the value of $\\delta=15$, suggested by VLBI observations (Bower \\& Backer 1998), the size of the source is $R< 1.6 \\times 10^{15}$~cm, an order of magnitude less than dimensions typically estimated for blazars ($10^{16}-10^{17}$~cm). In the case of a mass of the central supermassive black hole of $10^{8}M_{\\odot}$, the above distance would correspond to only about $50$ times the radius of the innermost stable orbit.  A direct alternative is to admit a large value of the Doppler factor, $\\delta \\sim 100$. Such large values of $\\delta$ are sometimes invoked to explain the rapid intra day variability observed in the radio band (e.g. Wagner \\& Witzel 1995). In this case, it would be possible to explain the short flare in terms of unsteadyness of the jet flow, due to a single non stationary shock (Hughes et al. 1985) possibly induced by a collision of two relativistic plasma shells in the jet (internal shock, Spada et al. 2001). The event discussed here could also result from an internal shock developping very close to the origin of the jet with less extreme values of $\\delta$. The moderate polarization observed at radio wavelengths is consistent with these scenarios, since it arises from regions much further out in the jet.  Although the data available are few and sparse,  it is possible that the present observation represents the ``tip of an iceberg'' leading the way to discover more short time scale events. We plan to continue our search of variability of blazars in the  \\emph{INTEGRAL} field of view."
    },
    "0601/astro-ph0601271_arXiv.txt": {
        "abstract": "We develop a covariant formalism to  study nonlinear perturbations of dissipative and interacting relativistic fluids. We derive nonlinear evolution equations for various covectors defined as linear combinations of the spatial gradients of the local number of e-folds and of some scalar quantities characterizing the fluid, such  as the energy density or the particle number density. For interacting fluids we decompose perturbations into adiabatic and entropy components and derive their coupled evolution equations, recovering and extending the results obtained in the context of the linear theory. For non-dissipative and noninteracting fluids, these evolution equations reduce to the conservation equations that we have obtained in recent works. We also illustrate geometrically the meaning of the covectors that we have introduced. ",
        "introduction": "Recently, we have developed \\cite{lv05a,lv05b} a new  formalism for nonlinear relativistic  perturbations, based on a covariant approach \\cite{Hawking:1966qi,ellis}  and inspired by the work of Ellis and Bruni \\cite{Ellis:1989jt}.  A key quantity in our formalism is the covector $\\zeta_a$, a linear combination of the spatial gradients of the energy density $\\rho$ and of $\\alpha$, the local number of e-folds, or integrated expansion along each worldline of the fluid.  As we showed in \\cite{lv05a,lv05b}, $\\zeta_a$, for a barotropic fluid, obeys a remarkably simple {\\it conservation equation}, which is {\\it exact and fully nonlinear}. In the linear approximation, this conservation equation reduces  to the usual conservation law for the linear curvature perturbation on uniform energy density hypersurfaces, usually denoted $\\zeta$. Thus, our covector $\\zeta_a$ can be seen as a covariant and nonlinear generalization of the usual $\\zeta$.  It must be stressed that, although our initial motivations were related to  cosmology, our formalism applies, in fact, to any relativistic fluid, whatever the underlying geometry. Thus, in the following we will not assume any particular type of spacetime, unless specified.  The purpose of the present work is to clarify the geometrical meaning of $\\zeta_a$, and to extend our formalism in two new directions. Whereas our formalism has been developed so far for a {\\it perfect} fluid, we consider here its extension to the case of dissipative relativistic fluids. Moreover, we consider the possibility of having several different {\\em interacting} fluids.     The covariant formalism of Ellis and Bruni \\cite{Ellis:1989jt} was first extended to a dissipative multifluid system in  \\cite{Hwang:1990am,Dunsby:1991vv}, where both the exact and the linear form of perturbation equations were given for noninteracting fluids. It was later generalized to the case of interacting fluids in \\cite{Dunsby:1991xk} for {\\em linear} perturbations.   The extension to interacting fluids is particularly useful in cosmology, where several types of matter coexist. As a typical example, one can mention the analysis of the cosmic microwave background anisotropies, which has been for instance considered within the covariant formalism and {\\em at linear order} in \\cite{Challinor}.   In this paper we work in the framework of our nonlinear formalism \\cite{lv05a,lv05b}, based on covectors such as $\\zeta_a$. As in the single perfect fluid case, we show that it is possible to derive simple, exact, and fully nonlinear evolution equations for various covectors, constructed as linear combinations of the spatial gradients of scalar quantities that characterize the fluids, such as the energy density or the number particle, and of the local number of e-folds $\\alpha$.  The  equations we obtain represent nonlinear generalizations of linear perturbation equations, some of which  have been already studied in the context of linear cosmological perturbations (see, in particular, \\cite{Malik:2002jb,Malik:2004tf}). However, our equations are covariant, i.e., independent of the coordinate system, and nonlinear, and it is straightforward to linearize or expand them up to  second, or higher, order in the perturbations. Furthermore, our covariant approach allows to identify easily which properties belong specifically to the linear or second order expansion, and which properties remain valid  at higher orders.    This work is organized as follows. In Sec.~\\ref{sec:covariant} we introduce the covariant formalism for dissipative fluids while in Sec.~\\ref{sec:covariantpert} we discuss the geometrical meaning of the nonlinear covector $\\zeta_a$. In Sec.~\\ref{sec:identity} we derive an identity that can be used to construct conservation equations for various nonlinear covectors, representing the perturbations of scalar quantities of a dissipative fluid. The conservation equations for these quantities are derived in Sec.~\\ref{sec:nonconservation}, while in Sec.~\\ref{sec:interacting} we extend our formalism to interacting fluids.  Finally, in Sec.~\\ref{sec:conclusion}, we conclude. ",
        "conclusions": "\\label{sec:conclusion}   In the present work, we have extended our covariant formulation for nonlinear perturbations to the case of dissipative relativistic  fluids, allowing for interactions. This extension could be used to tackle more sophisticated physical situations. Cosmology is such an example since the matter content of the universe is made of  several fluids which can interact.   In our approach, the important quantities are the covectors $\\zeta_a^{(f)}$, defined as linear combinations of the spatial gradients of  the number of e-folds $\\alpha$, and of some scalar fluid quantities $f$, for one or several fluids. These covectors are fully nonlinear and generalize the curvature perturbations on hypersurfaces of constant $f$ defined in the context of linear cosmological perturbation theory.  The non-conservation equation for $\\zeta_a$ associated with the energy density, that we obtain when the fluid is unperfect, can be derived and written in a form analogous to that found in our initial formalism with a single perfect fluid, provided that the pressure is modified such as to include dissipative effects.  For several interacting fluids, we can also define fully nonlinear quantities that generalize analogous quantities introduced in the linear theory. Remarkably, as in our initial formalism with a single perfect fluid, the evolution equations that govern these quantities ``mimic'' those found in the linear theory, with the advantage that they are covariant, i.e., independent of any coordinate system. Thus, it is straightforward to linearize them or expand them at second, or higher,  order in the perturbations.   Moreover, our fully nonlinear approach allows to identify which property is specific to linear  (or second) order and which property will remain valid at all orders.  For example, in the multifluid case, there are a priori several  nonlinear generalizations of the curvature perturbation, which depend on the choice of the reference frame.  In the cosmological context, they all reduce to the same quantity at linear order.      \\appendix"
    },
    "0601/astro-ph0601047_arXiv.txt": {
        "abstract": "The employment of a large area Phase Fresnel Lens (PFL) in a gamma-ray telescope offers the potential to image astrophysical phenomena with micro-arcsecond (\\muas) angular resolution \\cite{Skinner1}. In order to assess the feasibility of this concept, two detailed studies have been conducted of formation flying missions in which a Fresnel lens capable of focussing gamma-rays and the associated detector are carried on two spacecraft separated by up to 10$^6$ km. These studies were performed at the NASA Goddard Space Flight Center Integrated Mission Design Center (IMDC) which developed spacecraft, orbital dynamics, and mission profiles. The results of the studies indicated that the missions are challenging but could be accomplished with technologies available currently or in the near term.  The findings of the original studies have been updated taking account of recent advances in ion thruster propulsion technology. ",
        "introduction": "A unique feature of missions using Fresnel lenses for gamma-ray astronomy is the large distances over which precise station-keeping of two spacecraft is required. The concept of a mission using a Fresnel lens for ultra-high angular resolution gamma-ray astronomy has been studied by the NASA GSFC Integrated Mission Design Center (IMDC) \\cite{IMDC}. The IMDC is a resource dedicated to developing space mission concepts into advanced mission designs by providing system engineering analysis of all the flight systems and subsystems.   The IMDC process involves the mission scientists working with engineers of all the major engineering disciplines, e.g. mechanical, electrical, propulsion, flight dynamics, communications, mission operations, etc., in a highly interactive environment which naturally allows for inter-discipline communications and trade studies.  For the Fresnel mission studies, an additional formation-flying flight dynamics group also participated.  Two configurations of the mission were studied,  a definitive Fresnel mission with a $10^6$ km spacecraft separation and a smaller, pathfinder mission with a $10^5$ km focal length.  Each mission configuration completed a dedicated, one week IMDC study. After discussing the requirements and assumptions used as input for the IMDC studies, the mission profiles and key findings are summarized along with a re-analysis of the propulsion requirement to include the implications of recent advances in ion thruster technology. ",
        "conclusions": "Two gamma-ray astronomy missions employing Fresnel lenses were developed at the NASA GSFC Integrated Mission Design Center; a pathfinder mission with a $10^5$ km spacecraft separation and 10 micro-arcsecond (\\muas)  imaging ability and a definitive mission with a $10^6$ km separation and with a 1 \\muas~  angular resolution.  The development of the spacecraft was determined to be straightforward. While the studies determined that the flight dynamics and propulsion requirements are challenging, they can be accomplished with current ion propulsion technology if a re-pointing time scale of a week is used.  Furthermore, recent advances in ion thrusters have relaxed the propulsion requirements potentially allowing the incorporation of an additional detector-craft into a mission profile, thus doubling the number of viewed gamma-ray sources in a 5 year mission lifetime.  This work was supported by a grant from the NASA Revolutionary Aerospace Systems Concepts (RASC) program."
    },
    "0601/astro-ph0601667_arXiv.txt": {
        "abstract": "We qualitatively  study the  effects of gravitational  microlensing on our view of unresolved  extragalactic star formation regions.  Using a general gravitational microlensing  configuration, we perform a number of simulations that reveal that  specific imprints of the star forming region are imprinted, both photometrically and spectroscopically, upon observations.   Such observations  have  the potential  to reveal  the nature and size  of these star forming regions,  through the degree of variability observed  in a monitoring campaign, and  hence resolve the star formation regions  in distant galaxies which are  too small to be probed via more standard techniques. ",
        "introduction": "Gravitational microlensing is now a well established technique for the investigation  of the  distribution of  compact (dark)  matter  in the universe.   Furthermore, it  also provides  a powerful  tool  to study unresolved  sources, such as  in the  case of  the structure  of QSOs, through  temporal   differential  magnification  (e.g.    Yonehara  et al. 1998).  From an  observer's point of  view, gravitational microlensing  can be naturally  divided in  two different  regimes.  In  the  case of Galactic microlensing, the optical  depth is low and a  single star microlenses another star within the Galactic halo or in one of the galaxies in the Local  Group (Paczy\\'nski  1986a).   With Extragalactic  microlensing, where the light from a  distant quasar shines through a closer galaxy, the optical  depth is roughly unity  and many stars  contribute to the overall  microlensing effect (Paczy\\'nski  1986b).   This paper considers this latter regime, were the source region is populated by a number of hot, young stars in a star forming region.  Such a situation will occur  in strongly  lensed, multiply imaged  systems, such  as the multiple images  seen in galaxy  clusters (Mellier 1999), or  the case where a  isolated galaxy gravitationally lenses a  more distant galaxy (i.e. Warren et al 1996).  In a similar vein, Lewis \\&  Ibata (2001) investigated the effect of a cosmological distribution of compact objects on the surface brightness distributions   of  galaxies   at  $z$$<$0.5,   considering   a  small microlensing  optical  depth  ($\\leq$0.04)  and they  determined  that low-level  fluctuations  in  surface  brightness of  $\\sim$2\\%  should result.   Lewis  et al.   (2000)  extended  that  analysis to  distant galaxies observed through galaxy  clusters, assuming dark matter to be composed of compact objects.  Focusing upon Abell~370 as a case study, concluding  that  for  low-luminosity ($\\sim$10$^4$L$_\\odot$)  stellar populations would  show rapid fluctuations exceeding 10\\%  of the mean in the highest cases.  In  this contribution  we address  the question  of  what microlensing signatures should be apparent in the case of part of a galaxy which is lensed by another  galaxy.  In particular, if the  lensed parts of the source  galaxy are  regions  of star  formation,  highly dominated  by young, massive stars.  Such a  situation was recent presented by Smith et  al.    (2005)  who  reported   the  discovery  of  a   new  strong gravitationally lensed system, with an elliptical galaxy acting as the lens.  The  lens galaxy in this  system is at  redshift $z=0.0345$ and the source, proposed to be a star formation region, is at $z\\sim0.45$, with the arcs formed by the gravitational mirage showing `knots' of an extreme  blue  color  of  $B-I_c=1.1$  (extinction  corrected).   This discovery poses the  idea that microlensing in the  multiple images of these systems  might be able to  distinguish the type  of source stars involved in the mirage and help in the interpretation of its nature.   Within the  context of gravitational lensing, a  star formation region would appear as  a non-uniform source, composed of  a number of bright points in a more  extended background. Hence, the microlensing imprint of such  a source  should show quite  a different  variability imprint from  the  uniform   sources  typically  considered  in  gravitational microlensing experiments. The nature of this imprint is the basis of this current contribution. ",
        "conclusions": "We  described   in  this  contribution  how   to  apply  gravitational microlensing  to  the observations  of  unresolved extragalactic  star forming  regions.   The discussion  shows  that  due  to the  caustics configuration  in   the  magnification  maps  of   the  region,  rapid monitoring campaigns, both  photometric or spectroscopic, would reveal high  variability  fluctuations  due   to  the  number  of  early-type stars. The specific amount of variability will depend on the number of stars and  their distribution in the  region, as well as  on the exact configuration  of the microlenses  in the  lensing galaxy.   Thus, the study of a particular system requires the knowledge of a lens model to perform the right  simulations and the analysis of  the results should be based on  a statistical approach.  The advantage  of the method, if these  circumstances  take  effect,  is  that  we  might  be  able  to investigate star formation regions which are difficult to analyse with more traditional techniques."
    },
    "0601/astro-ph0601384_arXiv.txt": {
        "abstract": "{Cross sections for the rotational (de)excitation of CO by ground state para- and ortho-H$_2$ are obtained using quantum scattering calculations for collision energies between 1 and 520~cm$^{-1}$. A new CO$-$H$_2$ potential energy surface is employed and its quality is assessed by comparison with explicitly correlated CCSD(T)-R12 calculations.  Rate constants for rotational levels of CO up to 5 and temperatures in the range 5$-$70~K are deduced. The new potential is found to have a strong influence on the resonance structure of the cross sections at very low collision energies. As a result, the present rates at 10~K differ by up to 50$\\%$ with those obtained by \\citet{flower01} on a previous, less accurate, potential energy surface. ",
        "introduction": "Since its discovery in interstellar space more than 30 years ago \\citep{wilson70}, carbon monoxide (CO) has been extensively observed both in our own and in other galaxies. CO is indeed the second most abundant molecule in the Universe after H$_2$ and, thanks to its dipole moment, it is the principal molecule used to map the molecular gas in galactic and extragalactic sources. CO is also currently the only molecule to have been observed at redshifts as high as $z=6.42$ \\citep{walter03}. Collisional excitation of rotational levels of CO occurs in a great variety of physical conditions and emission from levels with rotational quantum numbers $J$ up to 39 have been identified in circumstellar environments \\citep{pepe96}.  Hydrogen molecules are generally the most abundant colliding partners in the interstellar medium, although collisions with H, He and electrons can also play important roles, for example in diffuse interstellar clouds. Rate coefficients for collisions with H$_2$ based on the potential energy surface (PES) of Jankowski \\& Szalewicz (1998, hereafter JS98) have been calculated recently \\citep{mengel01,flower01}. These rates were found to be in reasonable agreement (typically within a factor of 2 at 10~K) with those computed on older and less accurate PES (e.g. \\citet{schinke85}). However, in contrast to previous works, \\citet{mengel01} and \\citet{flower01} found the inelastic rates with $\\Delta J=1$ larger than those with $\\Delta J=2$ for para-H$_2$. This result reflects the crucial importance of the PES, whose inaccuracies are one of the main sources of error in collisional rate calculations, especially at low temperatures (see, e.g., \\citet{bala02}). The JS98 surface has been checked against experiment by a couple of authors. \\citet{gottfried01}, in particular, computed second virial coefficients below 300~K and showed that the attractive well of the PES was slightly too deep. A modification which consists of multiplying the PES by a factor of 0.93 was therefore suggested by these authors. This modification was applied by \\citet{mengel01} and it was indeed found to produce better agreement with measurements of pressure broadening cross sections. On the other hand, the original PES was found to give a very good agreement with experimental state-to-state rotational cross sections measured near 1000~cm$^{-1}$ \\citep{antonova00}. At these energies, however, only the short-range repulsive region of the PES is probed. Finally, \\citet{flower01} showed that the scaling of 0.93 has only minor influence on the rotational rates and he employed the original PES.  Very recently, a new PES has been computed by \\citet{jankowski05} (hereafter JS05). In contrast to the JS98 surface, which was calculated within the symmetry-adapted perturbation theory (SAPT), the JS05 surface has been obtained using the supermolecular approximation based on the CCSD(T) method using correlation-consistent basis sets complemented by bonding functions and basis set extrapolation. It turned that the major inaccuracy of JS98 was due to the use of a correction for the effects of the third and higher orders extracted in an approximate way from supermolecular SCF calculations. It recently became clear that this correction should not be used for nonpolar or nearly nonpolar monomers. If this term is removed from JS98, the agreement of the resulting surface with JS05 is excellent. \\citet{jankowski05} have employed their new PES to calculate rovibrational energy levels of the para-H$_2-$CO complex and they obtained a very good agreement (better than 0.1~cm$^{-1}$ compared to 1~cm$^{-1}$ for the JS98 surface) with the measurements by \\citet{mckellar98}. The computed values of the second virial coefficients also agree very well with the experimental ones. Full details can be found in \\citet{jankowski05}. In the present paper, we present new rate constants for the rotational excitation of CO by H$_2$ based on the JS05 surface. Our results are restricted to low temperatures (T$<$70~K) because the new PES calculations are found to affect the rotational cross sections only at low collision energy (E$_{\\rm coll}<$ 60~cm$^{-1}$). As shown below, differences with the results of \\citet{flower01} range from a few percent to 50~$\\%$. Such effects, if not dramatic, are significant in view of the special importance of the CO$-$H$_2$ system in astrophysics. Furthermore, the present results are of particular interest for comparisons with future measurements planned in Rennes by Sims and collaborators \\citep{sims05}. ",
        "conclusions": "Cross sections for the rotational (de)excitation of CO by ground state para- and ortho-H$_2$ have been computed using quantum scattering calculations for collision energies in the range 1$-$520~cm$^{-1}$. A new, highly accurate, CO$-$H$_2$ potential energy surface has been employed and it has been shown to strongly influence the resonance structure of the cross sections in the very low collision energy regime ($E_{\\rm coll} \\lesssim 60$~cm$^{-1}$). Conversely, at higher energies, the effect of the new potential was found to be only minor. As a result, the present rate constants are found to differ significantly from those of \\citet{flower01}, obtained on a previous and less accurate potential, only for temperatures lower than 70~K. Transitions among all levels up to $j_1=5$ only were computed as higher states are generally not populated at the low temperatures investigated here.  This work illustrates the relationship between the accuracy of the potential energy surface and the accuracy of the corresponding inelastic cross sections for a simple system of astrophysical relevance. Even at low temperatures, moderate inaccuracies in the surface (in the 10-20\\% range) result in semi-quantitative inelastic rates with typical errors in the 20-50\\% range with no dramatic amplification of errors. However, this optimistic conclusion presents also a harder counterpart, as we show that in order to obtain accurate inelastic rates one has to satisfy all three conditions, i.e., use an accurate surface, a properly converged angular expansion, and a detailed description of the cross-section resonances. Consequently, the improvement of the accuracy of inelastic rates, especially at higher temperatures when the number of channels is large, represents a considerable computational effort. In this respect, the work by \\citet{flower01} represented an excellent compromise to achieve a typical 20\\% accuracy below 100~K.  This work extends the initial objectives of the ``Molecular Universe'' European FP6 network (2005-2008) and participates to the current international efforts to improve the description of the underlying microscopic processes in preparation of the next generation of molecular observatories."
    },
    "0601/astro-ph0601498_arXiv.txt": {
        "abstract": "Two series of models and their yields are presented in this paper. The  first series consists of 20 $M_\\odot$ models with varying initial metallicity (solar down to $Z=10^{-8}$) and rotation ($\\upsilon_{\\rm ini}=0-600$\\,km\\,s$^{-1}$). The second one consists of models with an initial metallicity of $Z=10^{-8}$, masses between 20 and 85 $M_\\odot$ and average rotation velocities at these metallicities ($\\upsilon_{\\rm ini}=600-800$\\,km\\,s$^{-1}$). The most interesting models are the models with $Z=10^{-8}$ ([Fe/H]$\\sim-6.6$). In the course of helium burning, carbon and oxygen are mixed into the hydrogen burning shell. This boosts the importance of the shell and causes a reduction of the size of the CO core. Later in the evolution, the hydrogen shell deepens and produces large amount of primary nitrogen. For the most massive models ($M\\gtrsim 60$\\,$M_\\odot$), significant mass loss occurs during the red supergiant stage. This mass loss is due to the surface enrichment in CNO elements via rotational and convective mixing.  The yields of the fast rotating 20 $M_\\odot$ models can best reproduce (within our study) the observed abundances at the surface of extremely metal poor (EMP) stars. The wind of the massive models can reproduce the CNO abundances of the carbon--rich UMPs, in particular for the most metal poor star known to date, HE1327-2326. ",
        "introduction": "Precise measurements of abundances of extremely metal poor (EMP) stars have recently been obtained by \\citet{FS5,FS6,IER04}, \\ldots . These provide new constraints for the stellar evolution models \\citep[see][]{CMB05,Fr04,Pr04}. The most striking constraint is the need for primary $^{14}$N production in very low metallicity massive stars. Other possible constraints are an upturn of the C/O ratio with a [C/Fe] about constant or slightly decreasing (with increasing metallicity) at very low metallicities, which requires an increase (with increasing metallicity) of oxygen yields below [Fe/H]$\\sim$ -3. About one quarter of EMP stars are carbon rich (C-rich EMP, CEMP stars). \\citet{RANB05} propose a classification for these stars. They find two categories: about three quarter are main s-process enriched (Ba-rich) stars and one quarter are enriched with a weak component of s-process (Ba-normal). The two most metal poor stars known to date, HE1327-2326 \\citep{Fr05,Ao05} and HE 0107-5240 \\citep{Ch04} are both CEMP stars. These stars are believed to have been enriched by only one to several stars and we can therefore compare our yields to their observed abundances without the filter of a galactic chemical evolution model (GCE). In an attempt to explain the origin of the abundances observed as well as the metallicity trends, I computed pre-supernova evolution models of rotating single stars with metallicities ranging from solar metallicity down to $Z=10^{-8}$ following the work of \\citet{MEM05}. ",
        "conclusions": ""
    },
    "0601/astro-ph0601451_arXiv.txt": {
        "abstract": "We report the discovery of up to 35 new supernova remnants (SNRs) from a $42\\arcsec$ resolution 90cm multi-configuration Very Large Array survey of the Galactic plane covering $4.5\\arcdeg<\\ell <22.0\\arcdeg$ and $\\mid b\\mid < 1.25\\arcdeg$. Archival 20cm, 11cm, and 8 \\mum\\/ data have also been used to identify the SNRs and constrain their properties. The 90cm image is sensitive to SNRs with diameters $2.5\\arcmin$ to $50\\arcmin$ and down to a surface brightness limit of $\\sim 10^{-21}$ W~m$^{-2}$~Hz$^{-1}$~sr$^{-1}$.  This survey has nearly tripled the number of SNRs known in this part of the Galaxy, and represents an overall 15\\% increase in the total number of Galactic SNRs.  These results suggest that further deep low frequency surveys of the inner Galaxy will solve the discrepancy between the expected number of Galactic SNRs and the significantly smaller number of currently known SNRs. ",
        "introduction": "Statistical studies of supernova (SN) rates, based on OB star counts, pulsar birth rates, Fe abundance, and the SN rate in other Local Group galaxies, suggest that there should be many more supernova remnants (SNRs) in our Galaxy \\citep[$\\ga 1000$;][]{Li1991,Tammann1994} than are currently known \\citep[$\\sim 231$;][]{Green2004}.  This deficit is likely the result of selection effects acting against the discovery of old, faint, large remnants, as well as young, small remnants in previous low resolution and/or poor sensitivity Galactic radio surveys (e.g.\\ Green 1991). For example, past single dish surveys with $\\lambda\\ge 11$cm have resolutions $>4\\arcmin$ \\citep[e.g.][]{Haslam1982,Reich1990}.  The 20cm NRAO Very Large Array (VLA) Sky Survey (NVSS) has high resolution ($\\sim 45\\arcsec$) but poor surface brightness sensitivity \\citep[][]{Condon1998}. The missing remnants are likely concentrated toward the inner Galaxy where the diffuse Galactic plane synchrotron emission and thermal \\HII\\/ regions cause the most confusion.  This premise is supported by the large number of new SNRs discovered in the 4th quadrant by \\citet{Whiteoak1996} using the Molongolo Observatory Synthesis Telescope (MOST) at 35cm with $\\sim 45\\arcsec$ resolution and comparatively high sensitivity. Despite the success of the MOST survey, a significant shortfall compared to predictions remains -- most likely due to limited dynamic range.  Thus, more sensitive, high resolution surveys of the inner Galaxy at low radio frequencies are the key to determining whether the ``missing'' remnants exist or if our understanding of SN rates is significantly flawed. From multi-configuration VLA 90cm (330 MHz) observations of just 1~deg${}^2$ centered on SNR G11.2-0.3, Brogan et al. (2004) discovered three new SNRs, demonstrating the power of such observations.  This Letter presents the discovery of 35 SNRs (including the three described above) from a more extensive 90cm VLA survey. ",
        "conclusions": "Generally, the SNRs discovered in this survey are smaller and fainter than those previously known in this region. The mean and median diameters of the new SNRs are $\\sim 12\\arcmin$ and $8\\arcmin$. The SNR candidates have 1~GHz surface brightnesses in the range $\\Sigma_{\\rm 1~GHz}=(1-15)\\times 10^{-21}$ W~m$^{-2}$~Hz$^{-1}$~sr$^{-1}$. This result suggests that the Bonn 11cm completeness limit estimated by \\citet{Green2004} of $\\sim 10^{-20}$ W~m$^{-2}$~Hz$^{-1}$~sr$^{-1}$ is accurate and that the current survey is $\\sim 10$ times deeper.  This limited 90cm survey of 42.5~deg${}^2$ has produced a $\\sim 15\\%$ increase in the total number of known Galactic SNRs. The number of identified remnants within the survey boundaries (previously 19) has increased by nearly a factor of 3 to 54.  Based on completeness considerations and empirical evidence, respectively, \\citet{Helfand1989} and \\citet{Brogan2004} estimate that there should be between 65 to 85 SNRs in this region, in reasonable agreement with the observed number, given that we are insensitive to the smallest ($\\la 2.5\\arcmin$) and largest ($\\ga 0.8\\arcdeg$) size scale remnants, as well as Crab like remnants lacking radio shells.  To estimate the implications of these new SNR discoveries on the total number of Galactic SNRs we assume (1) the current SNR catalog \\citep{Green2004} is essentially complete for $\\mid b\\mid > 1.25\\arcdeg$ and for $\\mid\\ell\\mid >50\\arcdeg$; (2) that the Bonn 11cm survey in the 1st quadrant and the MOST 35cm survey in the 4th quadrant have similar levels of completeness \\citep[e.g.][]{Green2004}; and (3) the number of known SNRs between $\\mid b\\mid< 1.25\\arcdeg$ and $\\mid\\ell\\mid <50\\arcdeg$ (115) should be scaled by a factor of three (i.e.\\ the increase achieved in our limited survey). The $\\ell=50\\arcdeg$ cutoff is guided by the longitude extent of the inner Galactic disk. These assumptions result in a prediction of 460 for the number of Galactic SNRs, still a factor of two less than the expected $\\sim 1,000$ SNRs (see \\S 1).  However, five of the new SNRs are in close proximity to (or coincident with) the two largest previously known SNRs in the survey region: W28 and W30 (Figs.~1a,b), but appear distinct from them. Although we do not currently know if the new SNRs are physically close to the two known SNRs (along the line of sight), such a result would be unsurprising since high mass stars form in clustered environments. A new SNR has also been proposed to lie along the line of sight to the Vela SNR \\citep[i.e.\\ G266.2--1.2][]{Aschenbach1998}. Such superpositions throughout the plane may also explain some fraction of the ``missing'' remnants. The assumption of completeness for $\\mid\\ell\\mid >50\\arcdeg$ is also probably flawed as a few SNRs continue to be discovered in the outer Galaxy \\citep[e.g.][]{Kothes2001}.  Overall, these results suggest that the ``missing SNRs'' problem can be attributed to selection effects and not our understanding of SN rates.  Future instruments like the EVLA, LWA, ATA, and SKA will allow even deeper low frequency, high dynamic range surveys which will likely discover the remaining shortfall of Galactic SNRs."
    },
    "0601/astro-ph0601294_arXiv.txt": {
        "abstract": "We present new photometry for the Galactic thick disk globular cluster Palomar 11 extending well past the main sequence turn-off in the V and I bands.  This photometry shows noticeable, but depleted red giant and subgiant branches.  The difference in magnitude between the red horizontal branch (red clump) and the subgiant branch is used to determine that Palomar 11 has an age of $10.4\\pm 0.5 $ Gyr.  The red clump is used to derive a distance $\\mathrm {d}_\\sun=14.3\\pm0.4 $ kpc, and a mean cluster reddening of $\\mathrm{ E(V-I)}=0.40\\pm0.03$.  There is differential reddening across the cluster, of order $\\delta\\, \\mathrm{ E(V-I)} \\sim 0.07$.  The colour magnitude diagram of Palomar 11 is virtually identically to that of the thick disk globular cluster NGC 5927, implying that these two clusters have a similar age and metallicity.  Palomar 11 has a slightly redder red giant branch than 47 Tuc, implying that Palomar 11 is 0.15 dex more metal-rich, or 1 Gyr older than 47 Tuc. Ca II triplet observations \\citep{RHS97} favour the hypothesis that Palomar 11 is the same age as 47 Tuc, but slightly more metal-rich. ",
        "introduction": "The globular cluster Palomar 11 (Pal 11) is located at $\\rm \\alpha_{2000}=19^h45^m14.4^s $, $\\rm \\delta_{2000}=-8\\arcdeg 0\\arcmin 26\\arcsec $ ($b = -8 00$; $\\ell = 31.81$) in the constellation Aquila. It sits roughly 13 kpc from the Sun and 8 kpc from the galactic center with an estimated metallicity of $\\mathrm{[Fe/H]}=-0.4$ \\citep[][hereafter H96]{H96}.  It belongs to the population of metal-rich clusters with $R_{GC} \\ga 3\\,$kpc often associated with the thick disk.  However, Pal 11 is located well outside the plane of the disk at $\\mathrm{Z}= -3.5\\,$kpc \\citep{CER93}.  There has been only one published color magnitude diagram (CMD) for this cluster \\citep[][hereafter OBB01]{OBB01}.  OBB01 did note two other unpublished CMDs from literature abstracts which derived some cluster properties for Pal 11; \\cite{CER93} and \\cite{CS84}.  While these sources derive some cluster parameters (see Table \\ref{tbl1} for a list of published values), none of them report a value for the main sequence turnoff.  In this paper a new V,V-I CMD for Pal 11 with photometry reaching well past the main sequence turnoff is presented and used to determine the age of Pal 11.  There are currently only 6 other thick disk clusters with age determinations \\citep{SW02,DA05}.   In Section 2 the details of the observations and photometry are discussed. In \\S 3 we present the CMD and briefly discuss the blue straggler population, \\S 4 determines the cluster parameters, such as distance, reddening, metallicity, and age.  Finally, \\S 5 discusses these findings. ",
        "conclusions": "A new CMD for Pal 11 has been presented which reaches roughly 4 magnitudes past the main sequence turnoff. Using the magnitude of the RC, and the method of \\cite{AL02} leads to a distance of $\\mathrm{d}_{\\sun}=14.3 \\pm 0.4\\,$ kpc and a mean reddening of $\\mathrm{ E(V-I)}=0.40\\pm0.03$. These values are in general agreement with previously published values. There is differential reddening across cluster, such that stars in in the NorthEast of the cluster are approximately 0.07 mag redder in V -- I than stars in the SouthWest part of the cluster. Following the method of \\cite{SJ94}, we find that Pal 11 has a similar metallicity to the metal rich clusters 47 Tuc and NGC 5927.  Using the difference in magnitude between the SGB and RC, an age of $10.4\\pm0.5 $ Gyr is determined for Pal 11.  This is the first age estimate for Pal 11.  Very similar ages are found for 47 Tuc and NGC 5927.  After shifting to a common reddening and distance modulus, the CMDs of Pal 11 and NGC 5927 are nearly identical, which implies that these clusters have the same age and metallicity.  In contrast, the RGB of 47 Tuc is slighter bluer than the Pal 11 RGB.  This implies that Pal 11 and 47 Tuc must differ in age by 1 Gyr or in metallicity by 0.15 dex.  Given that the CMDs of NGC 5927 and Pal 11 are very similar and NGC 5927 has a higher Ca II triplet equivalent width than 47 Tuc, it is likely that Pal 11 has a similar age to 47 Tuc, but is slightly more metal-rich.  In agreement with previous studies \\citep{SW02,DA05} we find no evidence for an age dispersion among the thick disk clusters in the Milky Way.  This suggests that the stars in the thick disk clusters formed over a relatively short time-scale ($\\la 500\\,$Myr)."
    },
    "0601/astro-ph0601541_arXiv.txt": {
        "abstract": "The value of accurately knowing the absolute calibration of the polarizing elements in the Near Infrared Camera and Multi-Object Spectrometer (NICMOS) becomes especially important when conducting studies which require measuring degrees of polarization of close to 1\\% in the near infrared. We present a comprehensive study of all previously observed polarimetric standards using the NIC2 camera on NICMOS. Considering both pre- and post-NICMOS Cooling System observations we find variations in the polarimetry consistent with the effects of sub-pixel mis-alignments and the point spread function. We also measure non-zero results from unpolarized standards indicating an instrumental polarization of $p\\approx1.2\\%, \\theta\\approx88\\degr$. The lack of polarized and unpolarized standard stars with which to perform a comprehensive calibration study means we cannot be confident that the current calibration will be effective for a number of recent large NICMOS GO programs. Further observations of polarimetric standards are needed in order to fully characterize the behavior of NICMOS at around $p=1\\%$. ",
        "introduction": "While photometry and spectroscopy provide information on spatial distributions, chemical compositions, and dynamics, polarimetry adds a valuable extra data dimension that can be used to determine the geometry of astronomical objects and the properties of interstellar particles \\citep{aandm85,hranda94,kas03}. Polarized light may intrinsically originate from the emission process or it may be induced via interactions with a diffuse medium \\citep{coll93,land02}. In either case, the preferred orientation of the electromagnetic vector, as well a quantitative measure of that preference, can be determined by placing polarizing elements in the light paths of detectors and deriving the Stokes parameters \\citep{chan60}.  From the ground some care is needed in determining the Stokes parameters as the atmosphere introduces photometric variations with amplitudes similar to those expected from the polarized light. Space based polarimeters are not exposed to these anomalies and can use multiple polarizing filters to gather the required information. However, multichannel polarimeters \\citep{smandf75}, which are yet to be employed on space based polarimeters, do reduce some of the complexities of the data reduction through synchronous recovery of the Stokes parameters. In fact, with enough care, ground-based polarimeters can measure fractional polarizations of $10^{-6}$ or less \\citep{hough05}.  Dust frequently enshrouds many interesting astronomical objects and inhibits the transmission of optical wavelengths. This effect declines as the wavelength of the light increases, thus many studies are carried out in the infrared. It is therefore of great advantage to be able to carry out polarimetric studies with the Near Infrared Camera and Multi-Object Spectrometer (NICMOS) on board the {\\it Hubble Space Telescope} ({\\it HST}).  \\begin{deluxetable*}{lccccccc} \\tablecaption{Details of Observations \\label{tab:sample}} \\tablewidth{0pt} \\tablehead{ \\colhead{Target} &\\colhead{Type} &\\colhead{RA (J2000)} & \\colhead{Dec (J2000)} & \\colhead{Epoch} & \\colhead{Orientation (\\degr)} &\\colhead{PID} } \\startdata CHA-DC-F7 & B5V (Pol.)     & 10 56 12.91 & -76 35 54.2 & 1997-09-13 &  -36.48 & 7692 \\\\ &                & 10 56 12.03 & -76 35 52.0 & 1998-04-20 & -179.00 & 7958 \\\\ &                & 10 56 12.03 & -76 35 52.0 & 1998-07-05 & -106.49 & 7692 \\\\ &                & 10 56 12.13 & -76 35 53.2 & 2002-09-02 &  -50.83 & 9644 \\\\ &                & 10 56 12.15 & -76 35 53.3 & 2003-05-20 & -160.83 & 9644 \\\\ G191B2B   & DAw (Unpol.)   & 05 05 30.80 & +52 49 52.9 & 1997-12-24 &  -74.32 & 7904 \\\\ HD64299   & A1V (Unpol.)   & 07 52 25.61 & -23 17 50.3 & 1997-09-01 &   12.71 & 7692 \\\\ HD283812  & A2V (Pol.)     & 04 44 25.16 & +25 31 43.1 & 1997-09-28 &   14.52 & 7692 \\\\ &                & 04 44 25.16 & +25 31 42.4 & 1997-12-04 &  -47.48 & 7692 \\\\ HD331891  & A4III (Unpol.) & 20 12  2.29 & +32 47 45.2 & 1997-09-01 &  -93.57 & 7692 \\\\ &                & 20 12  2.08 & +32 47 42.2 & 1998-04-26 &   32.55 & 7958 \\\\ &                & 20 12  2.11 & +32 47 43.5 & 2002-09-09 & -109.63 & 9644 \\\\ &                & 20 12  2.11 & +32 47 43.5 & 2003-06-08 &    0.57 & 9644 \\\\ \\hline \\enddata \\tablecomments{Types taken from SIMBAD (http://simbad.u-strasbg.fr/Simbad), all other data are as they appear in the {\\it HST} archive (http://archive.stsci.edu/hst/). (Pol.) and (Unpol.) after the  stellar type refers to polarized and unpolarized standards respectively. } \\end{deluxetable*}  As with many instruments, NICMOS executes programs that cover a wide range of required accuracies and sensitivities. It is therefore essential to probe the polarimetric capabilities of NICMOS, especially in low polarization targets, and assess any signatures that may be introduced by the instrument itself. We present here such an assessment based on the previous efforts of, for example, \\citet{hsands00} and \\citet{hines02}. By examining the on-line data archive for NICMOS we can directly compare the results from polarimetric standards whose degree and orientation of polarization have been determined during other investigations. The study is limited to include data obtained with the NIC2 ($\\Delta\\lambda=1.9-2.1\\micron$) camera as the number of previous NIC1 ($\\Delta\\lambda=0.8-1.3\\micron$) studies, excluding the calibration data, is low. In addition, these NIC1 programs generally probe polarizations of $\\sim40\\%$, i.e., \\#7264, where high accuracies are not necessarily needed. However, future works similar to those carried out here may also be prudent for NIC1.  In \\S~\\ref{obs+dr} we describe the data and its subsequent reduction. The methods used in determining the degree and orientation of polarization are outlined in \\S~\\ref{ptheta}. The results are presented in \\S~\\ref{results} before being discussed and concluded in \\S~\\ref{discussion} and \\S~\\ref{cons} respectively. ",
        "conclusions": "\\label{cons}  We have conducted polarimetric analyses on all (un)polarized standards available in the NICMOS data archive for the NIC2 camera. The principle aim was to determine the behavior of the instrument at low levels of polarization. We have thoroughly tested our routine for deriving aperture polarimetry of objects and found it to be robust. It has also demonstrated that observed small radius variation in the polarimetric curves of growth can be attributed to the effects of sub-pixel mis-alignments between polarizing elements and the point spread function. We have found no evidence for persistence.  \\begin{figure}[t] \\vspace{1.0cm} \\plotone{f4.eps} \\caption[The Parallel Transmission Co-efficients]{ The values of $t_k$ required to null the unpolarized standard HD331891. The value of $t_3$ is held constant at 0.9667. }\\label{fig:coeffs} \\end{figure}  Our findings also indicate that there is a measured polarization of $p\\approx1.2\\%$ in unpolarized targets which may indicate an intrinsic instrumental polarization and the need for further tweaking of the on-orbit transmission co-efficients. Assuming we are detecting an instrumental effect we have corrected data for polarized targets in the ($Q,U$)-plane and found the dispersion around previous (ground based) polarization estimates to have been reduced, but not totally removed. We have also derived averaged values of $t_k$ that null the unpolarized standard HD331891 across all apertures, and found $t_1$ and $t_2$ to be 0.8717 and 0.8341, respectively, when fixing $t_3=0.9667$. In addition, we have attempted to investigate the possibility of an observationally dependent instrumental polarization, but are inhibited from any conclusions by an insignificant sample.  While such levels of residual intrinsic polarization may not hamper studies of highly polarized targets, these levels will have a detrimental effect on studies attempting to measure $p<5\\%$. It is clear that a more comprehensive calibration study of NICMOS is critical in order for a number of {\\it HST} programs to be carried out successfully. The current post-NCS calibration archive (one polarized and one unpolarized standard) is insufficient for this to occur effectively. An array of polarized and unpolarized standards should be observed, at many roll angles and in all quadrants of the instrument, in order for the low level polarization characteristics of NICMOS to be properly and fully investigated, understood and removed."
    },
    "0601/astro-ph0601631_arXiv.txt": {
        "abstract": "We review the structural properties of giant extragalactic HII regions and HII galaxies based on  two dimensional hydrodynamic calculations, and propose an evolutionary sequence that accounts for their observed detailed structure. The model assumes a massive and young stellar cluster surrounded by a large collection of clouds. These are thus exposed to the  most important star-formation feedback mechanisms:  photoionization and the cluster wind. The models show how the two feedback mechanisms compete with each other in the disruption of clouds and lead to two different hydrodynamic solutions: The storage of clouds into a long lasting ragged shell that inhibits the expansion of the thermalized wind, and the steady filtering of the shocked wind gas through channels carved within the cloud stratum. Both solutions are here claimed to be concurrently at work in giant HII regions and HII galaxies, causing their detailed inner structure. This includes multiple large-scale shells, filled with an X-ray emitting gas, that evolve to finally merge with each other, giving the appearance of shells within shells. The models also show how  the  inner  filamentary structure of the giant superbubbles is largely enhanced with matter ablated from clouds and how cloud ablation proceeds within the original cloud stratum. The calculations point at the initial contrast density between the cloud and the intercloud media as the factor that defines which of the two feedback mechanisms becomes dominant throughout the evolution. Animated version of the models presented can be found at http://www.iaa.csic.es/\\~{}eperez/ssc/ssc.html. ",
        "introduction": "Multiple studies during the last decades have addressed the impact of photoionization, stellar winds and supernova explosions on interstellar matter (ISM). Disruptive events that re-structure the birth place of massive stars and their surroundings, while leading to multiple  phase transitions in the ISM (see e.g. Comeron 1997; Yorke et al. 1989; Garc\\'i{}a-Segura et al. 2004; Hosokawa \\& Inutsuka 2005 and  Tenorio-Tagle \\& Bodenheimer 1988 and references therein). On the other hand, giant molecular clouds, the sites of ongoing star formation, present a hierarchy of clumps and filaments of different scales whose volume filling factor varies between  10$\\%$ to 0.1$\\%$ (see McLow \\& Klessen, 2004, and references therein). As the efficiency of star formation in molecular clouds is estimated to be $\\leq $ 10$\\%$ (Larson 1988; Franco et al. 1994) the implication is that the bulk of the cloud structure remains after a star forming episode. All of these studies are relevant within the fields of interstellar matter and star formation and, in particular, regarding  the physics of feedback, a major ingredient in the evolution of galaxies and thus in cosmology.  From the observations, we know that the energetics of major bursts of star formation allow us to track galaxies up to very high redshifts, however,  by the time these sources are found star formation is already in a well advanced stage of its evolution, and thus we do not know for certain, not even for the closest sources, which is the state of the matter left over from a major massive  burst of star formation. We do know that giant extragalactic HII regions and HII galaxies are excellent examples of the impact of massive stars on the ISM. They belong to the same class of objects because  they are powered by young  massive bursts of star formation. Also because of their similar physical size, morphology and inner structure. Detailed studies of 30 Doradus (see Chu \\& Kennicutt 1994; Melnick et al. 1999), NGC 604 (Sabalisck et al. 1995; Yang et al 1996) and several giant extragalactic HII regions and HII galaxies (Mu\\~noz-Tu\\~n\\'on et al. 1996; Telles et al. 2001; Ma\\'\\i z-Apell\\'aniz et al. 1999) have been designed with the aim of unveiling the inner structure and dynamics of the nearest examples. All of them present a collection of nested shells that enclose an X-ray emmiting gas and that may extend up to kpc scales. Some of the largest shells have stalled while others present expansion speeds of up to several tens of km s$^{-1}$. Detailed HST images have also confirmed these issues  for HII galaxies as well (Martin et al. 2002; http://hubblesite.org/newscenter/newsdesk/archive/releases/2004/35/image/a). All of these sources, as in the studies of Telles et al (2001) and Cair\\'os et al. (2001), present a central  bright condensation coincident with the massive burst of stellar formation. In giant HII regions and in some HII galaxies (cf. Figure 10) this is resolved as the brightest filament or a broken ragged shell sitting very close to the exciting cluster. The central bright condensations are  much smaller than the sizes of the  ionized volumes, which become more and more diluted at large radii to finally merge with the background galactic interstellar medium.  The energy powering these giant volumes comes from the recently found unit of massive star formation: super star clusters (Ho 1997; see also Lamers et al. 2004 and references therein). Massive concentrations of young stars within the range of 10$^5$ M$_\\odot$ to several $10^6$ M$_\\odot$, all within a small volume $\\sim$ 3 -- 10 pc. Multiple young super star clusters are found within the most luminous HII galaxies and starburst galaxies, as in M82 where a collection of two hundred of them have been found within the central 150 pc of the dwarf galaxy (Melo et al. 2005), while the most compact HII galaxies and particularly giant HII regions, are structured by the action of only one or  a few of these exciting clusters.  However, despite multiple efforts we still lack a detailed model. None of the numerical simulations in the literature, that have considered  either the powerful winds or the impact of the ionizing radiation on the surrounding ISM, or both, have reproduced the kpc-scale morphology and inner structure of giant HII regions and HII galaxies. There are thus a number of central questions that remain in this field. And thus apart from the fact that we ignore the initial condition, the distribution of matter around new massive bursts of star formation, it is crucial to understand how the central brightest condensations or ragged shells survive or withstand the full power of the exciting clusters, and how the multiple nested shells develop within the large-scale ionized volume and acquire their inner filamentary structure.  Here we address all these issues by means of two dimensional (2-D) hydrodynamics. As in all other studies, our initial conditions are arbitrary and perhaps much more so in this case, as we have surrounded the central powerful SSC by a large concentration of spherical, static, cloudlets, reflecting perhaps the simple fact that we do not know the initial distribution of interstellar matter around a new major burst(s) of star formation. Our initial and boundary conditions are also arbitrary, in the sense that we approach the energetics from stellar clusters following the frequently used mode of instantaneous star formation (see Leitherer \\& Heckman 1995), although a continuous star formation spread over a few Myr may be a more reasonable approximation. In section 2 we give a detailed description of the boundary and initial conditions used in our calculations.  Section 3 describes the results of four different cases in which we explore the impact that energetic winds or photoionization, or both, may have on  the selected cloudlet distribution and its surrounding gas. Our conclusions together with a full discussion of the results are given in section 4. ",
        "conclusions": "We have shown here the effects of feedback from a  massive stellar cluster into a selected cloudlet distribution, which although arbitrary, as it could have had a different extent or it could have considered clouds of different sizes, separations and locations, it has allowed us to explore a wide range of possibilities. These however, have lead to the two main possible physical solutions to be expected: partial pressure confinement of the shocked wind, and a stable and long-lasting filtering of the thermalized wind through the cloudlet stratum. Similar solutions would have been found for different  cloud densities and cloud distributions and/or for different SSC masses ($M_{SC}$) or energetics. Thus despite the arbitrary and simple boundary and initial conditions here assumed, the structures that develop within the flow resemble  the structure of giant HII regions and HII galaxies (see Figure 10). In particular we refer to the giant and multiple well structured shells evident at optical wavelengths, often referred to as nested shells (see Chu \\& Kenniccutt 1994), that enclose a hot X-ray emitting gas, as in 30 Dor (Wang 1999) and NGC 604 (Ma\\'\\i{}z-Apell\\'aniz et al. 2004). The calculations also lead to the slowly expanding brightest filament (or ragged shell) present in all sources close and around the exciting stars, and to the elongated, although much fainter,  filaments on either side of the open channels that reseamble the ends of elongated columns into the giant superbubbles.  None of these structural features appear in calculations that assume a constant density ISM, nor in those that allow for a constant density molecular cloud as birth place of the exciting sources, or in calculations where the sources are embedded in a plane stratified background atmosphere. For all of these features to appear, a clumpy circumstellar medium seems to be a necessary requirement. A medium that would refrain the wind from an immediate exit into the surrounding gas and that would also allow  for the build up of multiple channels through which the wind energy would flow in a less unimpedded manner. And thus giant HII regions and HII galaxies, both powered by massive star formation events producing  an ample supply of UV photons and a powerful wind mechanical energy, all of them seem to process their energy into a stratum of dense cloudlets sitting in the immediate vicinity of the star formation event.    \\begin{figure}[htbp] \\vspace{17.0cm} \\caption{The observational evidence. HST H$\\alpha $ images of the IIZw 40, the central regions of NGC 4214, and NGC 604 (in M33), and an ESO NTT image of 30 Doradus, all of them ploted in a linear parsec scale. For IIZw 40 the inset shows the most central region. HST images were retrieved from the HST archive at STScI. The NTT image courtesy of Jes\\'us Ma\\'\\i z-Apell\\'aniz and Rodolfo Barb\\'a. } \\label{fig-ha} \\end{figure}  Our calculations confirm many of the results by Pittard et al. (2005) for the rapid sequence of the hydrodynamical events that emerge from the wind-cloud early interaction. In particular, a global bow shock is to engulf two or more cloudlets if the separation between them is smaller than their diameter while  individual bow shocks are set around each condensation when they are further appart from each other, as shown in all small-scale Figures 2, 4 and 6. Also, the elongated tails of ablated matter behind bow shocks, are not necessarily aligned with the incoming stream velocity vector. The deviation in the alignment of tails is due to pressure gradients on either side of each cloud. As shown also by Pittard et al. (2005) and in our small-scale figures, tails that show deviations from the alignment with the incoming wind end up pointing away from the cloudlets, and when the flow is subsonic the tails end up pointing towards each other.  Here we have shown that the ablated matter that composes the various tails, ends up merging with other cloudlets further downstream to eventually compose the almost even density, slow moving ragged shell.  Given the large number of possibilities regarding SSCs, we consider  that cases A and B also seem relevant. One can imagine situations  in which the stellar photon output is more developed and dominant than the winds, as in the initial stages of the evolution of coeval clusters, and thus leading to results similar to those of case B. Also cases in which the photon flux has dropped sufficiently, as predicted for coeval clusters after 3 Myr of evolution, leading (as in case A) to a long-lasting SSC wind as the dominant feedback mechanism.    In the calculations here shown, as the selected cloudlet stratum covers most of the sky above the central SSC, all calculated cases follow a similar evolutionary track in which two competing hydrodynamical processes define the time evolution. On the one hand, photoionization that aims at establishing a constant density medium everywhere and, on the other, the wind mechanical energy that looks for all possible exits out of the cloudlet distribution. In all cases that consider both effects, the ensemble of clouds ends up being driven to compose an almost constant density, low velocity ragged shell, that acts as a barrier to most of the shocked wind.  In this way the wind matter becomes subsonic upon crossing the reverse shock and trapped behind the ragged shell. On the other hand, channels through the cloudlet stratum allow for the rapid streaming of the shocked wind gas and lead to the build up of large-scale superbubbles with a pronounced  inner filamentary structure. The latter is promoted by the ablation of clouds from each of the channels inner edges. The calculations also show that for individual giant bubbles to be easily recognized they would have to be fed through channels in the cloud stratum, carved well further apart from each other. Neighboring channels would rapidly end up bursting their carried  shocked wind energy into a single shell.  Photoionization tends to equalize the density everywhere, causing an evolution towards a higher even pressure, however its power to slow the effects of the wind, its power to make the wind less dominant strongly depend on the initial contrast density between the cloud and intercloud medium. And thus, the larger the initial contrast density, the less dominant and longer lasting the impact of the wind would be."
    },
    "0601/astro-ph0601407_arXiv.txt": {
        "abstract": "A generic prediction of hierarchical clustering models is that the mass function of dark haloes in dense regions in the Universe should be top-heavy.  We provide a novel test of this prediction using a sample of galaxies drawn from the Sloan Digital Sky Survey. To perform the test, we compare measurements of galaxy clustering in dense and underdense regions.  We find that galaxies in dense regions cluster significantly more strongly than those in less dense regions. This is true over the entire 0.1--30~Mpc pair separation range for which we can make accurate measurements. We make similar measurements in realistic mock catalogs in which the only environmental effects are those which arise from the predicted correlation between halo mass and environment. We also provide an analytic halo-model based calculation of the effect.   Both the mock catalogs and the analytic calculation provide rather good descriptions of the SDSS measurements. Thus, our results provide strong support for hierarchical models. They suggest that, unless care is taken to study galaxies at fixed mass, correlations between galaxy properties and the surrounding environment are almost entirely due to more fundamental correlations between galaxy properties and host halo mass, and between halo mass and environment. ",
        "introduction": "The correlation between galaxy properties (morphology, star formation rates, luminosity, color etc.) and the surrounding environment has been the subject of extensive studies in the last few decades:  dense environments are preferentially occupied by elliptical, red, luminous galaxies, whereas star formation rates are higher in less dense regions (Dressler 1980; Butcher \\& Oemler 1984; Norberg et al. 2001, 2002; Balogh et al. 2002; Gomez et al. 2003; Hogg et al. 2004; Kauffmann et al. 2004; Berlind et al. 2005; Croton et al. 2005). In hierarchical models, this behaviour is expected to be a consequence of the fact that galaxies are surrounded by dark matter halos, and the properties of halos (mass, formation time, concentration, internal angular momentum, etc.) are correlated with their environments (Mo \\& White 1996; Sheth \\& Tormen 1999, 2002, 2004; Lemson \\& Kauffmann 1999; Gottloeber et al. 2001; Avila-Reese et al. 2005; Gao et al. 2005; Harker et al. 2006; Wechsler et al. 2006).  Recently, we described how the clustering of galaxies can be used to test the assumption that the correlations between galaxy properties and their environments are {\\em entirely} a consequence of the correlations between haloes and their environments (Abbas \\& Sheth 2005). This is a strong assumption which significantly simplifies interpretation of the observed luminosity dependence of galaxy clustering (e.g. Zehavi et al. 2005).  It is also a standard assumption in current halo-model descriptions of galaxy clustering (see Cooray \\& Sheth 2002 for a review). The main goal of this paper is to perform this test.  This paper is arranged as follows: in Section 2 we show how galaxy clustering depends on environment in the SDSS (Adelman-McCarthy et al. 2006).  In particular, we measure the pair correlation function in redshift space, $\\xi(s|\\delta_s)$, for a range of environments $\\delta_s$, as well as the projected quantity, $w_p(r_p|\\delta_s)$; the latter is free of redshift-space distortions. These measurements are compared with similar measurements in carefully constructed mock catalogs, and from an analytic calculation based on the halo-model.  In both the mocks and the analytic calculation, correlations between galaxy properties and environment are entirely a consequence of the correlation between galaxy properties and halo masses, and between halo mass and environment.  We summarize our results in Section 3, where we also discuss some implications. An Appendix provides details of the analytical model, which generalizes our earlier (real-space) work so that it can be used to model redshift space measurements as well. ",
        "conclusions": "One of the luxuries of the latest generation of large-scale sky surveys is that they contain sufficiently many objects that one can study subsamples of galaxies divided up in various ways. Here, we have focused on the clustering of galaxies in a volume limited sample drawn from the SDSS, and studied how the clustering of these galaxies depends on environment. We find that galaxies in dense regions are considerably more strongly clustered than those in less dense regions (Figures~\\ref{zvls} and~\\ref{wprp}).  This is perhaps not so surprising---after all, a dense region is one in which many galaxies are crowded together. What is more surprising is that this dependence on environment is very well reproduced by numerical (Section~\\ref{mocks}) and analytic (Appendix~A) models in which the entire effect is due to the fact that galaxy properties correlate with the masses of their parent halos, and massive halos preferentially populate dense regions. Hierarchical models make quantitative predictions for this correlation between halo mass and environment, and so the agreement between our models and the measurements provides strong support for such models. In this respect, our results are consistent with those of Mo et al. (2004), Kauffmann et al. (2004), Berlind et al. (2005), Blanton et al. (2006) and Skibba et al. (2006); this is reassuring, since our methods are very different.  Our test of environmental effects is particularly interesting in view of recent work showing that, at fixed mass, haloes in dense regions form earlier (Sheth \\& Tormen 2004), and that this effect is stronger for low mass haloes (Gao et al. 2005; Harker et al. 2006; Wechsler et al. 2006).  Such a correlation is not part of our analytic model, nor is it included in our mock catalogs.  Presumably, the good agreement with the SDSS measurements is due to the fact that we have concentrated on luminous galaxies, and these populate the more massive haloes.  It will be interesting to see if this agreement persists at lower luminosities.  The agreement between our models and the measurements has an important consequence:  Unless care is taken to study a population at fixed halo mass, our results indicate that observed correlations between gastrophysical effects (e.g. ram pressure stripping, strangulation, harrassment) and environment are dominated by the fact that these effects also correlate with halo mass, and halo mass correlates with environment.  Larger samples will allow us to study if these trends persist to fainter, presumably less massive galaxies. And more distant samples will allow us to study if these trends evolve."
    },
    "0601/astro-ph0601227_arXiv.txt": {
        "abstract": "We present the results from a spectroscopic survey for intervening \\ion{Mg}{2} absorption in the spectra of 381 background QSOs conducted at the Multiple Mirror Telescope.  This survey complements our earlier SDSS EDR \\ion{Mg}{2} survey, extending our results to lower redshift ($z \\simeq 0.15$) and weaker \\ion{Mg}{2} $\\lambda2796$ rest equivalent width ($W_0^{\\lambda2796} \\simeq 0.1$\\AA).  We confirm two major results from that survey: the transition in the $W_0^{\\lambda2796}$ distribution at $W_0^{\\lambda2796} \\approx 0.3$\\AA, and the $W_0^{\\lambda2796}$-dependent evolution of the incidence of systems. The nature of $\\partial^2N/\\partial z \\partial W_0^{\\lambda2796}$ is consistent with the idea that multiple physically-distinct components/processes contribute to the incidence of \\ion{Mg}{2} absorption systems in a $W_0$-dependent manner and evolve at different rates.  A significant decrease in the total proper absorption cross section is detected in our MMT data for systems as weak as 1.0 \\AA\\ $\\le W_0^{\\lambda2796} < 1.5$\\AA\\ at $z\\lesssim 0.4$.  We discuss this $W_0$-dependent evolution in the context of the evolution of galaxy structures, processes including superwinds and interactions,  and damped-Ly$\\alpha$ absorbers.  We also consider the possibility that the observed redshift and $W_0^{\\lambda2796}$ dependence of the incidence of absorption in spectroscopic surveys for low-ion/neutral gas results from the effects of dust-induced extinction. ",
        "introduction": "Over the past 25 years, spectroscopic QSO surveys have established the statistics of intervening metal-line absorption systems, most notably those selected via \\ion{C}{4} and \\ion{Mg}{2} resonance transitions (e.g., Weyman et al. 1979; Lanzetta, Turnshek, \\& Wolfe 1987; Tytler et al. 1987; Sargent, Boksenberg \\& Steidel 1988; Steidel \\& Sargent 1992, hereafter SS92; Churchill, Rigby, Charlton, \\& Vogt 1999, hereafter CRCV99; Nestor, Turnshek, \\& Rao 2005a, hereafter NTR05). These systems identify structures based on their gas absorption cross section, and  are widely believed to select galactic structures and extended gaseous galactic halos (e.g., Bergeron \\& Boiss\\'{e}, 1991; Steidel et. al, 2002; Churchill, Kacprzak \\& Steidel 2005).  The results of the early \\ion{Mg}{2} surveys, which we summarize in NTR05, indicated that the distribution of rest equivalent widths of the $\\lambda2796$ line ($W_0^{\\lambda2796}$) can be described by a power law or an exponential, and that the incidence of the strongest systems seems to be decreasing with decreasing redshift.  In NTR05 we presented a large ($\\approx 1,300$ systems) \\ion{Mg}{2} survey using the Sloan Digital Sky Survey (SDSS) Early Data Release (EDR) QSO spectra.  We demonstrated that the $W_0^{\\lambda2796}$ distribution is fit very well by an exponential over the range 0.3\\AA\\ $\\le W_0^{\\lambda2796} \\lesssim 5$\\AA\\ and that previous power law fits over predict the incidence of strong lines.  However, extrapolating the NTR05 exponential to weak (0.02\\AA\\ $\\le W_0^{\\lambda2796}<0.3$\\AA) systems under predicts their incidence when compared to the results of CRCV99.  This apparent transition is consistent with the idea that weak \\ion{Mg}{2} absorbers are in part comprised of a population that  is physically distinct from intermediate/strong \\ion{Mg}{2} systems (e.g., Rigby, Charlton \\& Churchill 2002).  Since the transition occurs at the boundary of two different surveys, however, it is important to investigate the possibility that this is simply an artifact of some unknown systematic difference between the surveys.  NTR05 also investigated the incidence of \\ion{Mg}{2} systems as a function of  redshift for different ranges of $W_0^{\\lambda2796}$. The results were consistent with the total proper cross section for absorption decreasing with decreasing redshift, especially at $z\\lesssim1$, such that the stronger lines evolve faster.  The apparent evolution can be described by a steepening of the $W_0^{\\lambda2796}$ distribution with cosmic time.  However, the data only reach $z\\ge0.37$ and the evolutionary signal has high significance only for very strong ($W_0^{\\lambda2796} \\gtrsim 2.0$\\AA) systems.  Thus, extending the measurement of $\\partial N/\\partial z$ to lower redshift is a major goal of this work.  Furthermore, the evolution of \\ion{Mg}{2} systems over cosmic time has a more general significance, since \\ion{Mg}{2} absorbers trace neutral hydrogen. Extending the statistics of \\ion{Mg}{2} absorption to lower redshift is therefore important for the study of the evolution of the \\ion{H}{1} content of the Universe as well.  In fact, results from this survey have already been used to study the incidence of damped Ly$\\alpha$ systems at low redshift (Rao, Turnshek \\& Nestor, 2005).  Accurate measurements of the incidence of \\ion{Mg}{2} absorption systems as a function of $W_0^{\\lambda2796}$ and redshift for weak to ``ultra-strong''  (4\\AA\\ $\\lesssim W_0^{\\lambda2796}\\lesssim6$\\AA) systems and from high to low redshift, together with kinematic and relative abundance information from high-resolution spectroscopy (see, e.g., Ding, Charlton, \\& Churchill, 2005, and references therein) and composite-spectrum analysis (e.g., Nestor, Rao, Turnshek \\& Vanden Berk, 2003), are allowing the ability to understand the physical nature and evolution of the structures selected by such systems.  Here we present the results of a spectroscopic QSO survey conducted at the 6.5m Multiple Mirror Telescope on Mount Hopkins, AZ.  The goals of the survey are: 1) to provide a survey large enough to detect a significant number of intermediate/strong systems (0.3\\AA $\\le W_0^{\\lambda2796} \\lesssim 2$\\AA) and sensitive enough to measure weak (0.1\\AA\\ $\\le W_0^{\\lambda2796} < 0.3$\\AA) systems in a single survey; and 2) extend the measurement of their incidence to lower redshift ($z\\approx 0.15$).  In \\S2 we describe the observations and our data reduction procedure.  We present the results in \\S3.  In \\S4, we discuss the relevance of our results to the problem of understanding the nature of low-ion metal-line absorbers, and present our conclusions in \\S5. ",
        "conclusions": "We have detected and measured 140 \\ion{Mg}{2} absorption systems in the spectra of 381 QSOs obtained at the MMT Observatory.  These data include systems with $\\lambda2796$ rest equivalent width in the  range 0.1 \\AA\\ $\\le W_0^{\\lambda2796} \\le 3.2$ and extend our determination of the incidence of \\ion{Mg}{2} absorbers down to redshift $z=0.15$.  Coverage of $W_0^{\\lambda2796}$ across the value $W_0^{\\lambda2796}\\simeq 0.3$\\AA\\ is important since the distribution $\\partial N/\\partial W_0^{\\lambda2796}$ is represented very well by a single exponential function above $W_0^{\\lambda2796}\\simeq0.3$ \\AA, but diverges from the fit with an excess of absorbers below this value.  This effect was first noted in NTR05 by comparing our SDSS EDR survey data for $W_0^{\\lambda2796}\\ge 0.3$\\AA\\ with data from CRCV99 having $W_0^{\\lambda2796} < 0.3$\\AA.  Here, we confirm the nature of the $\\partial N/\\partial W_0^{\\lambda2796}$ distribution in a single survey thereby removing the concern that the putative upturn was an artifact of unknown systematic differences in the two independent surveys.  By extending the redshift coverage of our \\ion{Mg}{2} survey to lower redshift, we sample systems down to a look-back time of only $\\approx 2$ Gyrs, compared to $\\approx 4$ Gyrs for SDSS samples, and thereby increase the total look-back time covered by $\\approx 1/3$.  The low-redshift regime is of particular importance, especially for intermediate/strong systems, since the apparent evolution in their incidence is easily manifested only at low redshift.  While NTR05 report evolution in systems with $W_0^{\\lambda2796}\\gtrsim 2.0$\\AA\\ at redshifts $z \\lesssim 1$, we now detect evidence for evolution in systems with $W_0^{\\lambda2796}\\gtrsim 1.0$\\AA\\ at redshifts $z \\lesssim 0.5$.  The evolution apparent in our MMT sample is consistent with the nature of the evolution described in NTR05, whereby the incidence decreases from the no-evolution prediction (i.e., the total proper absorption cross section decreases) at a rate dependent on $W_0^{\\lambda2796}$.  It seems likely that multiple populations and processes contribute to the total cross section of \\ion{Mg}{2} absorption in a $W_0^{\\lambda2796}$-dependent manner, and that the different evolutionary rates of these populations and processes lead to the $W_0^{\\lambda2796}$-dependent evolution in $\\partial N/\\partial z$. It is also possible, however, that biases due to dust-induced extinction give rise to the detected evolution.  It will be necessary to determine the relative contribution of the various physically-distinct components/processes to the incidence of \\ion{Mg}{2} systems, as well as obtain stronger constraints on the evolution of the incidence, in order to better understand low-ion/neutral gas absorption systems.  Determining the relative contribution as a function of $W_0^{\\lambda2796}$ will help us further understand the evolution in time of each type (and, perhaps, vice-versa).  The results of our ongoing studies, such as the imaging of ``ultra-strong'' systems and the analysis of the full SDSS sample, are two such projects that should add to our understanding of the nature and evolution of these systems."
    },
    "0601/astro-ph0601011_arXiv.txt": {
        "abstract": "{ } { Accurate measurement of gravitational shear from images of distant galaxies is one of the most direct ways of studying the distribution of mass in the universe. We describe a new implementation of a technique for measuring shear that is based on the shapelets formalism.  } { The shapelets technique describes PSF and observed images in terms of Gauss-Hermite expansions (Gaussians times polynomials).  It allows the various operations that a galaxy image undergoes before being registered in a camera (gravitational shear, PSF convolution, pixelation) to be modeled in a single formalism, so that intrinsic ellipticities can be derived in a single modeling step.  } { The resulting algorithm, and tests of it on idealized data as well as more realistic simulated images from the STEP project, are described. Results are very promising, with attained calibration accuracy better than four percent (1 percent for round PSFs) and PSF ellipticity correction better than a factor of 20. Residual calibration problems are discussed.  } { }  \\keywords {Gravitational lensing -- Techniques: image processing -- Dark matter} ",
        "introduction": "Weak gravitational lensing is recognized as a profitable way to study the dark matter distribution in the universe, and with a series of ever wider-field astronomical cameras coming online, very large weak lensing surveys are being planned and performed. The inevitable source of noise in weak lensing measurements is shape noise, caused by the diversity of projected galaxy shapes on the sky. To beat down the noise, very large numbers of galaxy ellipticities need to be averaged: thus new lensing surveys such as the CFHT Legacy Survey (Hoekstra et al.\\ \\cite{Hoekstra2005}), the CTIO survey (Jarvis et al.\\ \\cite{Jarvis2003}), the VIRMOS-Descart survey (van Waerbeke, Mellier \\& Hoekstra \\cite{VanWaerbeke2005}) or the recently-approved Kilo-Degree Survey (KIDS) on the VLT Survey Telescope, will contain millions of background sources, which in principle should enable very accurate shear measurements.  This averaging down of the shape noise through large number statistics only makes sense if systematic errors can be controlled: the main one is still the blurring of the source images by atmosphere, telescope and detector pixels. The commonly-used Kaiser, Squires \\& Broadhurst~(\\cite{Kaiser1995}; henceforth KSB) and Luppino \\& Kaiser (\\cite{Luppino1997}) methods provide recipes for correcting these effects, and are very successful. Nevertheless, they are based on an idealized model of the effect of point spread function (PSF) convolution on ellipticity, and it is possible to construct plausible PSFs that it fails to correct properly (e.g., Hoekstra et al.~\\cite{Hoekstra1998}, Appendix~D). Therefore it seems unlikely that the KSB recipes will deliver the factor of 10 to 100 improvement in fidelity that will be required to exploit the new surveys (Erben et al.\\ \\cite{Erben2001}).  A number of different approaches have been put forward to improve PSF correction (Kuijken \\cite{Kuijken1999}; Kaiser \\cite{Kaiser2000}; Rhodes, Refregier \\& Groth \\cite{Rhodes2000}; Bernstein \\& Jarvis \\cite{Bernstein2002}; Refregier \\& Bacon \\cite{RefregierBacon2003}; Mandelbaum et al.\\ \\cite{Mandelbaum2005}). In this paper we present a new technique which combines elements from most of these.  The paper is organized as follows: in section 2 shear and ellipticity are defined, and the effect of one on the other. Section 3 summarizes the shapelets formalism, and describes how a shapelet description of a source and its PSF can be used to generate an ellipticity estimate that is useful for shear estimation. Section 4 substantiates the approach with idealized tests of the algorithm. In section 5 a software pipeline is presented that implements the full processing chain from an astronomical image to a shear estimate. Results of applying the pipeline to test data from the STEP project are shown in Section 6. In Section 7 we compare with other techniques, and Section 8 gives the conclusions. ",
        "conclusions": "Shapelets provide a neat framework in which to describe the transformations that an astronomical source image undergoes until it is registered on a detector. Gravitational shear, convolution with a PSF, and pixelation can all be modeled within the shapelets formalism.  All these elements can be combined into an efficient algorithm for extracting image ellipticities that can be used for accurate gravitational lensing shear measurements.  The implementation of these techniques into a working pipeline is presented in this paper. Tests show that the pipeline is able to recover input gravitational shears with very small calibration error (of the order of a percent) and PSF residual (better than a factor of 30 in PSF ellipticity).  It remains to be seen to what extent this approach can be applied successfully to diffraction-limited PSFs, which cannot be described easily with a shapelet expansion. A different set of basis functions for the expansion might be the answer. Further possible improvements are also under investigation."
    },
    "0601/astro-ph0601606.txt": {
        "abstract": "{ We calculate the broad-band photometric evolution of unresolved star clusters, including the preferential loss of low-mass stars due to mass segregation, in a simplified way. The stellar mass function of a cluster evolves due to three effects: (a) the evolution of the massive stars reduces their number; (b) tidal effects before cluster-wide mass segregation reduce the mass function homogeneously, i.e. independently of the stellar mass; (c) after mass segregation has completed, tidal effects preferentially remove the lowest-mass stars from the cluster. These effects result in a narrowing of the stellar mass range. These three effects are described quantitatively, following the results of $N$-body simulations, and taken into account in the calculation of the photometric history, based on the {\\sc galev} cluster evolution models for solar metallicity and a Salpeter mass function. We find the following results:\\\\ (1) During the first $\\sim$ 40\\% of the lifetime of a cluster its colour evolution is adequately described by the standard {\\sc galev} models (without mass segregation) but the cluster gets fainter due to the loss of stars by tidal effects. During this phase the colour evolution is the same for clusters with and without initial mass segregation.\\\\ (2) Between $\\sim$ 40 and $\\sim$ 80\\% of its lifetime (independent of the total lifetime) the cluster gets bluer due to the loss of low-mass stars. This will result in an underestimate of the age of clusters if standard cluster evolution models are used. The correction increases from 0.15 dex for a cluster with a total lifetime of 3 Gyr to 0.5 dex for clusters with a total lifetime of 30 Gyr.\\\\ (3) After $\\sim$ 80\\% of the total lifetime of a cluster it will rapidly get redder. This is because stars at the low-mass end of the main sequence, which are preferentially lost, are bluer than the AGB stars that dominate the light at long wavelengths. This will result in an overestimate of the age of clusters if standard cluster evolution models are used.\\\\ (4) Clusters with mass segregation and the preferential loss of  low-mass stars evolve along almost the same tracks in colour-colour diagrams as clusters without mass segregation. Therefore it will be difficult to distinguish this effect from that due to the cluster age for unresolved clusters. Only if the total lifetime of clusters can be estimated then the colours can be used to give reliable age estimates.\\\\ (5) The changes in the colour evolution of unresolved clusters due to the preferential loss of low-mass stars will affect the determination of the star formation histories of galaxies if they are derived from clusters that have a total lifetime of less than about 30 Gyr.\\\\ (6) The preferential loss of low-mass stars might explain the presence of old ($\\sim$13 Gyr) clusters in NGC 4365 which are photometrically disguised as intermediate-age clusters (2 -- 5 Gyr), if the expected total lifetime of these clusters is between 16 and 32 Gyr. It may also explain the concentration of these clusters towards the center of NGC 4365.  % ",
        "introduction": "\\label{sec:1}  The preferential loss of low-mass stars due to mass segregation in star clusters is not taken into account in the cluster photometry models available in the literature. However, this effect is of importance for the photometric evolution of (``simple'') stellar population models (e.g. Leitherer et al. 1999; Bruzual \\& Charlot, 1993; Anders \\& Fritze-v. Alvensleben 2003) and hence the luminosity and mass functions of cluster populations that are derived from the analysis of observed spectral energy distributions (SEDs) of unresolved clusters compared to standard models. In this paper we describe the expected effect of this preferential loss of low-mass stars during the evolution of dissolving clusters, using a simplified description for the dynamical mass-loss mechanism. We compare and confirm this approach based on the results of {\\it N}-body simulations.  {\\it N}-body simulations of clusters in tidal fields show that the massive stars tend to concentrate towards the cluster center and the low-mass stars preferentially populate the outer regions. This ``mass segregation'' is commonly observed in open and globular clusters in the Milky Way and in the Magellanic Clouds (e.g., de Grijs et al. 2002a,b,c and references therein).  An immediate and important consequence of mass segregation is that low-mass stars tend to be more easily, and predominantly, ejected from star clusters than high-mass stars (but see Brandl et al. 2001; de Grijs et al. 2002a). Recently, Baumgardt \\& Makino (2003; hereafter BM2003) calculated the changes in the MF of clusters in the tidal field of the Galaxy. They showed that the amplitude and shape of the MF initially decrease almost homologously: the mass distribution of the stars that survive stellar evolution decreases with almost constant shape up to some critical time, after which the mass distribution of the low-mass stars becomes very steep due to the preferential loss of the lowest-mass stars. Their \\nbody\\ simulations of clusters at different galactocentric distances in the Galaxy and with different orbits (circular vs. eccentric) show that mass segregation becomes important when a cluster has lost a certain fraction of its mass (approximately 70\\%), almost independent of the dissolution time-scale. In these simulations the role of binaries is not taken into account, although binaries are important for mass segregation (e.g., Inagaki \\& Saslaw 1985; Bonnell \\& Davies 1998; Portegies Zwart et al. 1999). Therefore, in real clusters mass segregation might occur earlier than predicted by the \\nbody\\ simulations of BM2003. In some clusters mass segregation might even be established right from the beginning (e.g. Hillenbrand 1997). We will refer to this as ``primordial mass segregation''.   In this paper we describe the expected effect of the preferential loss of  low-mass stars, due to mass segeration, on the integrated photometric evolution of clusters. We use a simplified method that enables us to explain and predict the effects for clusters of widely different lifetimes (e.g. as derived by Lamers, Gieles \\& Portegies Zwart 2005a).  The structure of this paper is as follows. In Sect. 2 we provide an overview of the observational evidence for primordial and dynamical mass segregation. We discuss the assumptions of our cluster models in Sect. 3, i.e., we give a description of the decreasing mass of a cluster due to the combination of stellar evolution and evaporation or dissolution. In Sect. 4 we describe the way in which we treat the loss of stars from the cluster due to stellar evolution and dissolution. In Sect. 5 we predict the photometric evolution of clusters, taking the preferential loss of low-mass stars into account. We compare the results with those of the standard {\\sc galev} simple stellar population models (Schulz et al. 2002;, Anders \\& Fritze-v.-Alvernsleben 2003) and describe the expected errors in the age determinations of unresolved clusters. Sect. 6 contains a discussion and a suggestion of the a possible explanation for the presence of an apparent intermediate age cluster population in the giant elliptical galaxy NGC 4365. The conclusions are in Sect. 7.  %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% %%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%%% ",
        "conclusions": "We predicted the photometric evolution of unresolved clusters with different total life times, \\tdis, in the $B,~V,~R,~I$ and $K$ bands of the {\\sl HST}/WFPC2 broad-band filter system, with mass loss due to stellar evolution and evaporation of low mass stars taken into account. The photometric evolution is compared with that of cluster models not including mass segregation nor the preferential loss of low-mass stars. We call these the ``standard'' models. We considered clusters with a total lifetime in the range of 0.3 to 30 Gyr. We obtained the following results.\\\\  \\begin{enumerate} \\item{} During the first part of the lifetime of the cluster, i.e. during the first $\\sim$ 40\\%, irrespective of the total life time \\tdis, the photometric evolution is the same as predicted for the standard models if the decreasing mass is taken into account. The dissolution of the cluster makes it fainter in all bands, but the colours are unaffected. \\item{} Between $\\sim$ 40 and $\\sim$ 80\\% of its total lifetime the cluster is {\\it bluer} than predicted by the standard models. This is due to the fact that the cluster has lost a large fraction of its (red) low-mass stars. This effect is small: $\\Delta (V-I)\\simeq 0.03$ mag for clusters with $\\tdis=1$ Gyr, but increases steeply with increasing \\tdis. It is about 0.1 mag for $\\tdis=10$ Gyr. This implies that the age of the clusters will be {\\it underestimated} when standard models are used. The error in the age estimate is about 0.15 dex in $\\log(t)$ if $\\tdis=3$ Gyr, 0.30 dex if $\\tdis=10$ Gyr and 0.5 dex if $\\tdis=30$ Gyr. Thus, the age of clusters with a total lifetime of 20 Gyr and a real age of 14 Gyr, will erroneously be estimated as about 4 Gyr on the basis of the $V-I$ and $V-K$ photometry. \\item{} Between about $\\sim$ 0.80 and 1.0 \\tdis\\ the clusters are much {\\it redder} than predicted by the standard models. At those ages the AGB stars are the dominant contributors to the photometry at long wavelengths. These AGB stars are redder than the stars at the low-mass end of the main sequence. So the removal of the lowest mass main-sequence stars will make the cluster redder than predicted by standard models. This effect increases from $\\Delta(V-I)\\simeq 0.0$ mag at $t \\simeq 0.8 \\tdis$ to 0.3 mag at 0.95 \\tdis. This reddening will result in an {\\it overestimate} of the age of clusters based on broad-band photometry from $B-V$ to $V-K$ colours if  standard cluster evolution models are used. The effect is large and can grow to an overestimate of a factor $\\sim$ 4 near the end of the cluster's life. \\item{} The changes in colour due to mass segregation and the preferential loss of low-mass stars occurs almost along the same lines in the colour-colour plots as the photometric evolution of the standard models. This makes it difficult to distinguish this effect from reddening due to the age of clusters. \\item{} The predicted photometric history of clusters with inital cluster-wide mass segregation is indistinguishable from that of clusters without initial mass segregation. This is due to the fact that mass segregation will occur quite rapidly (within $\\sim$0.20 \\tdis), even if there were no initial mass segregation. Because both the total lifetime of a cluster and the time for mass segregation depend on the half-mass relaxation time, mass segregation will occur at about a constant fraction of \\tdis. \\end{enumerate}  $N$-body simulations have shown that mass segregation and the preferential loss of low-mass stars will occur in clusters in tidal fields. Even if there is no initial mass segregation, the effects predicted by our models will occur in real clusters.    %****************************************************************************** %******************************************************************************"
    },
    "0601/hep-th0601097_arXiv.txt": {
        "abstract": "\\noindent The dynamical  effects of the tensor modes of the geometry are investigated in the context of curvature bounces. Since the bouncing behaviour implies sharp deviations from a radiation-dominated evolution, significant back-reaction effects of relic gravitons may be expected at short wavelengths. After developing a  general iterative  framework  for the calculation of dynamical back-reaction effects, explicit analytical and numerical examples are investigated for different parametrizations of the energy-momentum pseudo-tensor(s) of the produced gravitons.  The reported results suggest that dynamical  back-reaction effects are a necessary ingredient for a consistent description of bouncing models at late times. ",
        "introduction": "\\setcounter{equation}{0}  Relic gravitons are copiously produced in the early Universe due to the pumping action of the background geometry \\cite{gris1,gris2}. If a quasi-de Sitter phase of expansion  is followed by a radiation-dominated phase, the logarithmic energy spectrum (in units of the critical  energy density), customarily denoted by  $\\Omega_{\\rm gw}(\\nu)$, is quasi-flat  \\cite{star} (see also \\cite{inflsp,GG1,sahni}) for present (physical) frequencies $\\nu$ ranging between $10^{-16} \\,\\,{\\rm Hz}$ and, approximately, $100 \\,\\,{\\rm MHz}$. The transition from the radiation-dominated to the matter epoch produces an infrared branch where $\\Omega_{\\rm gw} \\sim \\nu^{-2}$ between $10^{-18} {\\rm Hz}$ and $10^{-16} {\\rm Hz}$ \\cite{rub,FP,ab}.  In pre-big bang models \\cite{SFD,pbb}, the spectrum of relic gravitons is far from scale-invariant. The spectral slope of $\\Omega_{\\rm gw}(\\nu)$ (for frequencies larger than $10^{-16}$ Hz) is violet, i.e. $ \\Omega_{\\rm gw} \\sim \\nu^{\\gamma}$ with $\\gamma > 1$ \\cite{GG1,GG3} (see also \\cite{occhionero} and, in a complementary perspective, \\cite{norma}). Minimal pre-big bang models are characterized by a slope $\\gamma = 3$ that become progressively less steep as the frequency increases in the kHz and MHz region.  Depending on the features of the cosmological model, the energy density of relic gravitons may affect the dynamics and change the time evolution of the scale factor as well as of the other homogeneous quantities. If the equation of state of the sources of the background geometry is stiffer than radiation (i.e. $ p = w\\rho$ with $w> 1/3$) the relativistic gravitons of short wavelengths can change the dynamical evolution, as it was noticed in the context of the stiff model of Zeldovich \\cite{gris2}.  When a stiff phase follows an inflationary stage of expansion \\cite{mg1} (see also \\cite{zelsta,fordpp}\\footnote{It is relevant to remark that the authors of Refs.  \\cite{zelsta} and \\cite{fordpp} choose to compute the energy and pressure density of the produced particles by means of a perturbative expansion whose small parameter is the deviation from conformal coupling ($|\\xi -1/6|$ in the notations of \\cite{fordpp}). This approach (also discussed in \\cite{birrel,BD}) may be applied, with some caveats, in the minimally coupled case (corresponding to $\\xi \\to 0$) where, it could be argued that the expansion parameter does not exceed $1$. In the present investigation this expansion will not be used  since the coupling of relic gravitons to the background geometry is not close to conformal. Furthermore the expansion in   $|\\xi -1/6|$ may partially break down during the bouncing regime. When needed, numerical techniques will be employed.} ) the back-reaction effects of the produced gravitons set a limit on the possible duration of the stiff (post-inflationary) phase. This aspect is relevant in different situations and, for instance, in the context of quintessential inflationary models \\cite{pv1} where $\\Omega_{\\rm gw} \\propto \\nu$ for frequencies larger than the ${\\rm m Hz}$ \\cite{mg2,mg3,sas} (see also \\cite{maxdan1,li,cruise,picasso} for possible detection strategies). Blue spectral  slopes have been also derived in the context of quintessential models on the brane \\cite{var1,var2} (see also \\cite{sami1}). Constraints arising from the production of relic gravitons in quintessential models have been discussed in \\cite{orito}.  Defining as $\\tau$ the conformal time coordinate and $k$ as the comoving wave-number, two physically different regimes appear naturally in the problem. If $ k\\tau \\ll 1$ (short wavelength limit) the Fourier modes of the tensor fluctuations of the geometry are said to be super-adiabatically amplified. In the opposite regime, i.e. $k\\tau \\gg 1 $ (long wavelength limit) the tensor modes are oscillating. In the short wavelength limit  the  effective equation of state obeyed by the relic gravitons is radiative, i.e. $\\langle p_{\\rm gw} \\rangle = \\langle \\rho_{\\rm gw} \\rangle /3$ where $\\langle p_{\\rm gw} \\rangle$ and $\\langle  \\rho_{\\rm gw} \\rangle$ correspond to the averaged pressure and energy densities. In this case, $\\langle \\rho_{\\rm gw} \\rangle \\sim a^{-4}$, where $a(\\tau)$ is the scale factor. In the long wavelength limit the effective equation of state may depend on the salient features of the background evolution. For instance, during the transition from a quasi-de Sitter stage of expansion to a radiation dominated stage,  the modes of long wavelengths are compatible with an effective equation of state $\\langle p_{\\rm gw} \\rangle =- \\langle \\rho_{\\rm gw} \\rangle /3$ implying that $ \\langle \\rho_{\\rm gw} \\rangle \\sim a^{-2}$. In this limit the system behaves, effectively, as a  generalized fluid dominated by the spatial gradients of the tensor modes of the geometry.  While the conclusions mentioned in the previous paragraph must rely on specific definitions of the energy and pressure densities of the relic gravitons, it is  well known that, in general, it is impossible to assign a localized energy density to the gravitational field \\cite{landau}.  This caveat does not exclude the possibility of adopting consistent frameworks for the analysis of a  gravitational energy-momentum pseudo-tensor.  To be more specific, the tensor modes of the geometry can be characterized by a rank-two tensor $h_{ij}$ defined in the three spatial dimensions (that will taken to be flat), i.e. \\begin{equation} ds^2 = a^2(\\tau) [ d\\tau^2 - (\\delta_{ij} + h_{ij}) dx^{i} dx^{j}], \\qquad h_{i}^{i} = \\partial_{i}h^{i}_{j} =0, \\label{pertmet} \\end{equation} where $\\delta_{ij}$ is Kroeneker symbol and $h_{ij}$, being traceless and divergence-less, carries two independent degrees of freedom that correspond to the two polarizations of a gravitational wave in a conformally flat Friedmann-Robertson-Walker (FRW) background \\footnote{Since $h_{ij}$ is transverse and traceless, the direction of propagation can be chosen to lie along the third axis and, in this case the two physical polarizations of the graviton will be $h_{1}^{1} = - h_{2}^{2} = h_{\\oplus}$ and $ h_{1}^{2} = h_{2}^{1} = h_{\\otimes}$. This will be the nomenclature followed in the present paper.}. The simplest way of estimating the impact of the created gravitons on the background dynamics is by computing the lowest-order nonlinear corrections to the Einstein tensor \\begin{equation} {\\cal G}_{\\mu}^{\\nu} = R_{\\mu}^{\\nu} - \\frac{1}{2} \\delta_{\\mu}^{\\nu} R, \\label{ein} \\end{equation} where $R_{\\mu}^{\\nu}$ and $R$ are, respectively, the Ricci tensor and the Ricci scalar. The nonlinear corrections to the Einstein tensor, will consist, to lowest order,  of quadratic combinations of $h_{ij}$ that can be formally expressed as \\footnote{In the present notations the Planck length will be defined as $\\ell_{\\rm P} = \\sqrt{8\\pi G}$ in units $\\hbar= c =1$. Natural gravitational units $16\\pi G = 2 \\ell_{\\rm p}^2 = 1$ will also be adopted when needed.} \\begin{equation} \\ell_{\\rm P}^2 {\\cal T}_{\\mu}^{\\nu} = - \\delta^{(2)}_{\\rm t} {\\cal G}_{\\mu}^{\\nu}, \\label{ps1} \\end{equation} where the superscript at the right hand side denotes the second-order fluctuation of the corresponding quantity while the subscript refers to the tensor nature of the fluctuations. This procedure is essentially the one described in \\cite{isacson1,isacson2} and has been re-explored, in a related context, by the authors of Refs. \\cite{abramo1,abramo2} mainly in connection with conventional inflationary models where the Universe is always expanding\\footnote{Recently in \\cite{bab} the field theoretical formulation of General Relativity was further developed with the purpose of deriving an energy-momentum tensor for the gravitational field. Our approach, in the present context, is more modest since we simply want to compare the dynamical consequences of different possible definitions of the energy-momentum pseudo-tensor of the relic gravitons.}  In \\cite{ford1} (see also \\cite{ford2}) a complementary perspective was invoked and  it was observed that, by perturbing the Einstein-Hilbert action to second order in the amplitude of the tensor modes of the geometry in a Friedmann-Robertson-Walker background, each single polarization of the graviton obeys, up to total derivatives, the action of a minimally coupled scalar degree of freedom \\begin{equation} \\delta_{\\rm t}^{(2)} S = \\frac{1}{2} \\int d^{4} x\\,\\, a^2(\\tau)\\,\\,\\eta^{\\alpha\\beta} \\partial_{\\alpha} h\\partial_{\\beta} h, \\label{FPdef} \\end{equation} where the dimensionless amplitude $h$ has been defined as \\begin{equation} h = \\frac{h_{\\oplus}}{\\sqrt{2} \\ell_{\\rm P}} =\\frac{h_{\\otimes}}{\\sqrt{2} \\ell_{\\rm P}}. \\label{def1} \\end{equation} Consequently, following \\cite{ford1}, a natural ansatz for describing the back-reaction of the created gravitons on the background space-time would be to use, for each single polarization, the energy-momentum tensor of a minimally coupled scalar field. This approach has been exploited, for instance, in \\cite{sahni} and in \\cite{fordpp} for estimating the energy and pressure densities of the relic gravitons in a multi-stage Universe. A possible weakness of this second approach is that Eq. (\\ref{FPdef}) holds, strictly speaking, in a Friedmann-Roberston-Walker background. In spite of this, it is clearly useful to bear in mind this possibility.  The purpose of the present paper is to scrutinize the dynamical back-reaction effects of the relic gravitons in the extreme situation where the background geometry undergoes a contraction that smoothly evolves, through a bounce, into an expanding phase.  This sequence of events takes place, for instance, in pre big-bang models. The bouncing beahaviour is related to strong deviations from a radiative Universe. Consequently, short wave fluctuations are expected to modify or even restrict the dynamical range of bouncing models since, in this regime, the effective equation of state of relic gravitons is the one of radiation. For long wavelength fluctuations, the qualitative expectation  is more difficult to formulate since, as already mentioned, the effective equation of state is sensitive to the specific features of the background evolution before and after the bounce.  Consider then, as an example, one of the simplest realizations of bouncing  dynamics, i.e. the one provided by minimal pre-big bang models where, in the Einstein frame, the Universe first contracts as $a(\\tau) \\sim \\sqrt{-\\tau/\\tau_{1}}$ and then expands, in the post-big bang phase, as $a(\\tau) \\sim \\sqrt{\\tau/\\tau_{1}}$. An accurate description of the bouncing regime in the region $ -\\tau_{1} <\\tau <\\tau_{1}$ may be achieved, for instance,  by means of a non-local potentials \\cite{mv1,mv2,gm} that depend on the so-called shifted dilaton which is usually introduced in the context of the scale-factor-duality symmetry \\cite{SFD} (or more generally, in the treatment of $O(d,d)$ transformations \\cite{jna1,jna2} acting on the background fields of the low-energy string effective action). The features of the regular bouncing solutions discussed in \\cite{maxb1,maxb2}  are such that the r\\^ole of the dilaton potential is only crucial in the bouncing region while, away from the bounce (i.e. $|\\tau| > \\tau_{1}$) the the geometry is driven by the kinetic energy of the dilaton field.   As soon as the tensor modes reach into the super-adiabatic regime, i.e. $k \\tau \\ll 1$ the effective pressure density of the relic gravitons becomes $\\langle p_{\\rm gw} \\rangle \\simeq \\langle \\rho_{\\rm gw} \\rangle$. Consequently, at least up to the bounce, the energy density of the relic gravitons  is not likely dominate on the energy density of the dilatonic sources since the two components evolve exactly in the same way with the scale factor. While it is plausible that only around the bounce the short wavelength modes become dynamically relevant, it is also clear that the whole qualitative picture should be corroborated by a detailed numerical and analytical treatment.  It will be instructive to conduct the  calculation of the dynamical back-reaction within different ansatz for the energy-momentum pseudo-tensor. Then the results will be compared. A byproduct of the present analysis will indeed be that different forms of the energy-momentum pseudo-tensor lead, quantitatively, to the same back-reaction effects.  The present paper is then organized as follows. In Section 2 a class of string inspired bouncing cosmologies will be introduced. Section 3 is devoted to the analysis of the production of relic gravitons and to the different definitions of their energy-momentum pseudo-tensor in the framework of bouncing solutions. Section 4 deals with the analytical estimates of the effective barotropic indices. Section 5 describes the implementation of a self-consistent (iterative) scheme that allows to compute the back-reaction effects of relic gravitons. Section 6 summarizes the main lessons to be drawn from the present analysis and contains also the concluding remarks. Finally, various technical results that are relevant for the present investigation are reported in the appendix.  \\renewcommand{\\theequation}{2.\\arabic{equation}} ",
        "conclusions": "\\setcounter{equation}{0}  Dynamical back-reaction of relic gravitons is relevant in different contexts, so that the nature of the induced physical effects may vary. A particularly interesting situation is the one of bouncing solutions of pre-big bang type where back-reaction effects are known to be important already  from qualitative estimates. One of the purposes of the present investigation is to propose a theoretical scheme where more quantitative predictions could be derived. The present findings are then useful for implementing a smooth exit from the pre-big bang phase with consequent production of radiation. The bouncing solutions examined here may also be viewed as a practical theoretical laboratory where the methods for the analysis of the back-reaction effects of relic gravitons may be tested. In this respect, the reported results suggest that the r\\^ole of the produced gravitons is crucial for determining the asymptotic state of the solution at late times.  The strategy followed in the present analysis relies on the interplay between analytical methods and numerical calculations. We have been interested in situations where the Universe undergoes an accelerated contraction before the bounce, while, after the bounce, the zeroth-order solution exhibits a decelerated expansion. After quantizing the tensor modes of the geometry, the initial values of the field operators have been set well before the bounce. At that time all the gravitons were of short-wavelengths in comparison with the typical time scale of the background evolution. Then, the field operators have been evolved in the Heisenberg representation and the mixing coefficients have been accurately determined both analytically and numerically. This step allowed  to understand what kind of regularization is required for the energy and pressure densities of the produced gravitons. The second step has been to adopt a consistent ansatz for the gravitational energy-momentum pseudo-tensor. Two complementary definitions  of the energy and pressure densities of the relic gravitons have then been scrutinized and compared both analytically and numerically. On the basis of the reported results, these two different approaches lead to compatible quantitative results.  The self-consistent analysis of back-reaction effects leads naturally to an integro-differential problem. Consequently, an iterative method for its solution  has been proposed and tested in the specific dynamical framework of bouncing solutions that arise in the context of the low-energy string effective action supplemented by a non-local dilaton potential. It turns out that the treatment both of the background and of the tensor fluctuations becomes rather simple by selecting a different parametrization of the time coordinate which does not coincide with the usual conformal or cosmic times. In this time parametrization analytical bouncing solutions can be derived in the Einstein frame metric.  By looking at the structure of the low-energy string effective action, it is possible to infer that not only relic gravitons but also other massless fields may lead to similar effects whose analysis is certainly one of the possible developments of future studies. In this respect, the present investigation dealt with the minimal field content of the model with the aim of achieving a reasonably accurate description of the underlying physical processes. Thus, some of the ideas discussed here will hopefully be useful also for the description of dynamical systems with a richer field content.  \\newpage \\begin{appendix} \\renewcommand{\\theequation}{A.\\arabic{equation}} \\setcounter{equation}{0}"
    },
    "0601/astro-ph0601014.txt": {
        "abstract": " ",
        "introduction": "\\label{intro} Recent distance-redshift surveys \\cite{perl98,riess98,tonry03,snobs,Riess:2004nr,Astier:2005qq} of cosmologically distant Type Ia supernovae (SnIa) have indicated that the universe has recently (at redshift $z\\simeq 0.5$) entered a phase of accelerating expansion. This expansion has been attributed to a dark energy \\cite{Sahni:2004ai} component with negative pressure which can induce repulsive gravity and thus cause accelerated expansion. The evidence for dark energy has been indirectly verified by Cosmic Microwave Background (CMB) \\cite{Spergel03} and large scale structure \\cite{Tegmark04} observations.  The simplest and most obvious candidate for this dark energy is the cosmological constant \\cite{Sahni:1999gb} with equation of state $w=\\frac{p}{\\rho}=-1$. The extremely fine tuned value of the cosmological constant required to induce the observed accelerated expansion has led to a variety of alternative models where the dark energy component varies with time. Many of these models make use of a homogeneous, time dependent minimally coupled scalar field $\\phi$ (quintessence\\cite{Peebles:1987ek,Sahni:1999qe}) whose dynamics is determined by a specially designed potential $V(\\phi)$ inducing the appropriate time dependence of the field equation of state $w(z) = {{p(\\phi)} \\over {\\rho(\\phi)}}$. Given the observed $w(z)$, the quintessence potential can in principle be determined. Other physically motivated models predicting late accelerated expansion include modified %% gravity\\cite{Perrotta:1999am,Torres:2002pe,Nojiri:2003ni}, Chaplygin gas\\cite{Kamenshchik:2001cp}, Cardassian cosmology\\cite{Freese:2002sq}, theories with compactified extra dimensions\\cite{Perivolaropoulos:2002pn,Perivolaropoulos:2003we}, braneworld models\\cite{Sahni:2002dx} etc. Such cosmological models predict specific forms of the Hubble parameter $H(z)$ as a function of redshift $z$. The observational determination of the recent expansion history $H(z)$ is therefore important for the identification of the viable cosmological models.  The most direct and reliable method to observationally determine the recent expansion history of the universe $H(z)$ is to measure the redshift $z$ and the apparent luminosity of cosmological distant indicators (standard candles) whose absolute luminosity is known. The luminosity distance vs. redshift is thus obtained which in turn leads to the Hubble expansion history $H(z)$.  The goal of this review is to present the methods used to construct the recent expansion history $H(z)$ from SnIa data and discuss the most recent observational results and their theoretical implications. In the next section I review the method used to determine $H(z)$ from cosmological distance indicators and discuss SnIa as the most suitable cosmological standard candles. In section 3 I show the most recent observational results for $H(z)$ and discuss their possible interpretations other than accelerating expansion. In section 4 I discuss some of the main theoretical implications of the observed $H(z)$ with emphasis on the various parametrizations of dark energy (the simplest being the cosmological constant). The best fit parametrizations are shown and their common features are pointed out. The physical origin of models predicting the best fit form of $H(z)$ is discussed in section 5 where I distinguish between minimally coupled scalar fields (quintessence) and modified gravity theories. An equation of state of dark energy with $w<-1$ is obtained by a specific type of dark energy called {\\it phantom energy} \\cite{Caldwell:1999ew}. This type of dark energy is faced with theoretical challenges related to the stability of the theories that predict it. Since however the SnIa data are consistent with phantom energy it is interesting to investigate the implications of such an energy. These implications are reviewed in section 6 with emphasis to the Big Rip future singularity implied by such models as the potential death of the universe. Finally, in section 7 I review the future observational and theoretical prospects related to the investigation of the physical origin of dark energy and summarize the main conclusions of this review. ",
        "conclusions": "The question of the physical origin and dynamical evolution properties of dark energy is the central question currently in cosmology. Since the most sensitive and direct probes towards the answer of this question are distance-redshift surveys of SnIa there has been intense activity during the recent years towards designing and implementing such projects using ground based and satellite observatories. Large arrays of CCDs such as MOSAIC camera at Cerro Tololo Inter-American Obsrevatory, the SUPRIME camera at Subaru or the MEGA-CAM at the Canada-France-Hawaii Telescope (CFHT) are some of the best ground based tools for supernova searches. These devices work well in the reddest bands (800-900nm) where the ultraviolet and visible light of redshifted high-z SnIa is detected. Searches from the ground have the advantages of large telescope apertures (Subaru for example has 10 times the collecting area of the Hubble Space Telescope (HST)) and large CCD arrays (the CFHT has a 378-milion pixel camera compared to the Advanced Camera for Surveys on HST which has 16 million pixels). On the other hand the advantage of space satellite observatories like the HST include avoiding the bright and variable night-sky encountered in the near infrared, the potential for much sharper imaging for point sources like supernovae to distinguish them from galaxies in which they reside and better control over the observing conditions which need not factor in weather and moonlight.  The original two SnIa search teams (the Supernova Cosmology Project and the High-z Supernova Search Team) have evolved to a number of ongoing and proposed search projects both satellite and ground based. \\begin{figure} \\centering % Use the relevant command for your figure-insertion program % to insert the figure file. % For example, with the option graphics use \\includegraphics[bb=40 70 550 790,width=8.0cm,height=12.0cm,angle=-90]{fig17.ps} % % If not, use %\\picplace{5cm}{2cm} % Give the correct figure height and width in cm % \\caption{Ongoing and proposed SnIa search projects with the corresponding redshift ranges.} \\label{fig:17}       % Give a unique label \\end{figure} These projects (see Fig. 17) include the following: \\begin{itemize} \\item {\\bf The Higher-z Supernova Search Team (HZT)\\cite{Strolger:2004kk} and the GOODS\\cite{Giavalisco:2003ig} team of HST:} This has originated from the High-z Supernova Search Team and has A. Riess of Space Telescope Sci. Inst. as its team leader. This team is in collaboration with the GOODS program (Great Observatories Origin Deep Survey) using the ACS of the HST to detect and analyze high redshift ($0.5<z<2$) SnIa. Successive GOODS observations are spaced by 45 days providing 5 epochs of data on two fields: the Hubble Deep Field (HDF) north and south. Whereas the GOODS team adds these images to build a superdeep field, the HZT subtracts the accumulated template image from each incoming frame. Thus the HZT has already detected more than 42 supernovae in the above redshift range. \\item {\\bf Equation of State: SupErNovae trace Cosmic Expansion (ESSENCE)\\cite{Sollerman:2005qj}:} This has also originated from the High-z Supernova Search Team and has C. Stubbs of the Univ. of Washington, C. Smith and N. Suntzeff of Cerro Tololo as its team leaders. This ongoing program aims to find and measure 200 SnIa's in the redshift range of $0.15<z<0.7$ where the transition from decelerating to accelerating expansion occurs. Spectroscopic backup to the program comes from the ground based Gemini, Magellan, VLT, Keck and MMT Obsevatory. The ESSENCE project is a five-year endeavor, with the goal of tightly constraining the time average of the equation-of-state parameter $w = p/\\rho$ of the dark energy. To help minimize systematic errors, all of their ground-based photometry is obtained with the same telescope and instrument. In 2003 the highest-redshift subset of ESSENCE supernovae was selected for detailed study with HST. \\item {\\bf The Supenova Legacy Survey (SNLS)\\cite{Palanque-Delabrouille:2005yf}:} The CFHT Legacy Survey aims at detecting and monitoring about 1000 supernovae in the redshift range $0<z<1$ with Megaprime at the Canada-France-Hawaii telescope between 2003 and 2008. High-z spectroscopy of SnIa is being carried on 8m class telescopes (Gemini, VLT, Keck). Team representatives are: C. Pritchet (Univ. Victoria), P. Astier (CNRS/IN2P3), S. Basa (CNRS/INSU) et. al. The SNLS has recently released the first year dataset \\cite{Astier:2005qq}. \\item {\\bf Nearby Supernova Factory (SNF)\\cite{Wood-Vasey:2004pj}:} The Nearby Supernova Factory (SNF) is an international collaboration based at Lawrence Berkeley National Laboratory. Greg Aldering of Berkeley Lab's Physics Division is the principal investigator of the SNF. The goal of the SNF is to discover and carefully study 300 to 600 nearby Type Ia supernovae in the redshift range $0<z<0.3$. \\item{\\bf Carnegie SN Project (CSP)\\cite{Freedman:2004uz}:} The goal of the project is the comprehensive study of both Type Ia and II Supernovae in the local ($z < 0.07$) universe. This is a long-term program with the goal of obtaining exceedingly-well calibrated optical/near-infrared light curves and optical spectroscopy of over 200 Type~Ia and Type~II supernovae. The CSP takes advantage of the unique resources available at the Las Campanas Observatory (LCO). The team leader is R. Carlberg (Univ. of Toronto).\\item {\\bf Supernova Acceleration Probe (SNAP)\\cite{Aldering:2002dp}:} This is a proposed space mission originating from LBNL's Supernova Cosmology Project that would increase the discovery rate for SnIa's to about 2000 per year. The satellite called SNAP (Supernova / Acceleration Probe) would be a space based telescope with a one square degree field of view with 1 billion pixels. The project schedule would take approximately four years to construct and launch SNAP, and another three years of mission observations. SNAP has a 2 meter telescope with a large field of view: 600 times the sky area of the Hubble Space Telescope\u0092s Wide Field Camera. By repeatedly imaging ~15 square degrees of the sky, SNAP will accurately measure the energy spectra and brightness over time for over 2,000 Type Ia supernovae, discovering them just after they explode. \\end{itemize} These projects aim at addressing important questions related to the physical origin and dynamical properties of dark energy. In particular these questions can be structured as follows: \\begin{itemize} \\item Can the accelerating expansion be attributed to a dark energy ideal fluid with negative pressure or is it necessary to implement extensions of GR to understand the origin of the accelerating expansion? \\item Is $w$ evolving with redshift and crossing the PDL? If the crossing of the PDL by $w(z)$ is confirmed then it is quite likely that extensions of GR will be required to explain observations. \\item Is the cosmological constant consistent with data? If it remains consistent with future more detailed data then the theoretical efforts should be focused on resolving the coincidence and the cosmological constant problems which may require anthropic principle arguments. \\end{itemize} The main points of this brief review may be summarized as follows: \\begin{itemize} \\item {\\it Dark energy} with {\\it negative pressure} can explain SnIa cosmological data indicating accelerating expansion of the universe.\\item The existence of a {\\it cosmological constant} is consistent with SnIa data but other {\\it evolving} forms of dark energy {\\it crossing the $w=-1$} line may provide better fits to some of the recent data (Gold dataset). \\item New {\\it observational projects} are underway and are expected to lead to significant progress in the understanding of the properties of {\\it dark energy}. \\end{itemize} %\\begin{theorem} %Theorem text\\footnote{Footnote} goes here. %\\end{theorem} % % % BibTeX users please use %"
    },
    "0601/astro-ph0601233_arXiv.txt": {
        "abstract": "We study the evolution of the mass function of dark matter halos in the concordance $\\Lambda$CDM model at high redshift. We employ overlapping (multiple-realization) numerical simulations to cover a wide range of halo masses, $10^7-10^{15}h^{-1}M_\\odot$, with redshift coverage beginning at $z=20$. The Press-Schechter mass function is significantly discrepant from the simulation results at high redshifts. Of the more recently proposed mass functions, our results are in best agreement with Warren et al. (2005). The statistics of the simulations -- along with good control over systematics -- allow for fits accurate to the level of $20\\%$ at all redshifts. We provide a concise discussion of various issues in defining and computing the halo mass function, and how these are addressed in our simulations. ",
        "introduction": "Dark matter halos occupy a central place in the paradigm of gravitationally-driven structure formation arising from the nonlinear evolution of primordial Gaussian density fluctuations. Gas condensation, resultant star formation, and eventual galaxy formation occurs within halos. Consequently, the halo profile and mass function are central ingredients in phenomenological models of nonlinear clustering of galaxies. The distribution of halo masses -- the halo mass function -- and its time evolution, are also sensitive probes of cosmology.  The halo mass function at the high-mass end (cluster mass scales) is exponentially sensitive to the amplitude of the initial density perturbations, the mean matter density parameter, $\\Omega_{m}$, and to the dark energy controlled late-time evolution of the density field. The last feature, particularly at low redshifts, $z<2$, allows cluster observations to constrain the dark energy content, $\\Omega_{\\Lambda}$, and the equation of state parameter, $w$~\\cite{Holder01}.  The halo mass function is also of considerable interest at high redshift, relating to questions such as predictions for quasar abundance and formation sites~\\cite{HaimanLoeb01}, the formation history of collapsed baryonic halos, and the reionization history of the Universe~\\cite{Furlanetto05}. Recent results from the Wilkinson Microwave Anisotropy Probe (WMAP) (Kogut et al. 2003; Spergel et al. 2003) indicate that reionization could have begun at redshifts as high as $z\\sim 20$. Much of the work on possible reionization scenarios is based on the simple Press-Schechter (PS) mass function~(Press \\& Schechter~1974, Bond et al. 1991) the use of which can lead to imprecise predictions for the reionization history.  Simulations play a dual role in characterization of the halo mass function. If only a few fixed sets of cosmological parameters and a finite dynamic range are required, simulations can produce valuable results. In order to investigate a variety of cosmologies and different scenarios for physical processes, e.g., reionization, it is nevertheless very convenient, if not necessary, to have accurate analytic fitting relations. Simulations can be used to validate these fits over a wide (albeit, discretely sampled) range of parameters.  Various numerical studies of the mass function have been carried out over different mass and redshift ranges. The closest to the present work are Reed et al. (2003) and Springel et al. (2005); in comparison to their results, our halo mass range goes deeper by three orders of magnitude, with good statistics and control of systematics out to $z=20$, substantially higher than in these papers. (We review results from other work below.) Essentially, the earlier results are in very good agreement with the Sheth-Tormen mass function~\\cite{Sheth99}, at redshifts $z\\leq 10$. As we show below, various fitting formulae given in the literature -- most tuned to simulation results at $z=0$ -- can differ substantially in their predictions at high redshifts, by as much as a factor of two. Therefore, it is important to carry out simulations of sufficient dynamic range and accuracy to test these predictions.  In order to extract the mass function from simulations, different questions have to be addressed, such as: How is the mass function to be defined? When do the first halos form in a simulation?  When must the simulation be started in order to capture these halos?  What force and mass resolution is required to capture halos of a certain mass at a specific redshift? We have derived and tested certain criteria to ensure that our simulations capture the halos of interest; details will be given elsewhere~\\cite{lukic06}.  The paper is organized as follows. In Section~\\ref{massf} we discuss popular mass function formulae, previous work, and strategies for determining the halo mass function at high redshifts. The simulations and mass function results are discussed in Section~\\ref{code}. Criteria for mass and force resolution and initial redshift needed to span the desired mass and redshift range are given here. We present our conclusions in Section~\\ref{conclusion}. ",
        "conclusions": "\\label{conclusion} In this paper we have studied the evolution of the mass function starting at a redshift of $z=20$ and covering a halo mass range of $10^7$ to $10^{15}h^{-1}M_\\odot$. Our results incorporate new halo-based N-body error control criteria that are described in more detail in Luki\\'c et al. (2006). We find that the Press-Schechter mass function deviates significantly from our results. More recent mass function fits are in better agreement; in particular, the recently introduced fitting function of Warren et al. (2005) agrees at the $20\\%$ level over the entire redshift range.  The precise agreement of the numerically obtained halo growth function as well as the evolution of the mass function with the (evolved $z=0$) Warren fit demonstrates the remarkable result that the evolution of the mass function is completely controlled by the linear growth of the variance of the linear density field.  In order to find a mass function fit relevant to observations, several hurdles remain to be overcome, including reaching agreement on an appropriate definition of halo mass (White 2001) and improving the precision and accuracy of N-body codes beyond the current state of the art (O'Shea et al. 2005; Heitmann et al.  2005). Depending on the level of precision required, as White (2002) points out, ``it may not be sufficient to use a simple parametrized form'' in constraining cosmological parameters with the mass function.  The error control criteria developed in this work have a natural application in high-resolution AMR simulations in the setting of refinement and error control criteria. Work in this direction is in progress."
    },
    "0601/astro-ph0601143_arXiv.txt": {
        "abstract": "Radiatively-driven flow in a luminous disk is examined in the subrelativistic regime of $(v/c)^1$, taking account of radiation transfer. The flow is assumed to be vertical, and the gravity and gas pressure are ignored. When internal heating is dropped, for a given optical depth and radiation pressure at the flow base (disk ``inside''), where the flow speed is zero, the flow is analytically solved under the appropriate boundary condition at the flow top (disk ``surface''), where the optical depth is zero. The loaded mass and terminal speed of the flow are both determined by the initial conditions; the mass-loss rate increases as the initial radiation pressure increases, while the flow terminal speed increases as the initial radiation pressure and the loaded mass decrease. In particular, when heating is ignored, the radiative flux $F$ is constant, and the radiation pressure $P_0$ at the flow base with optical depth $\\tau_0$ is bound in the range of $2/3 < cP_0/F < 2/3 + \\tau_0$. In this case, in the limit of $cP_0/F = 2/3 + \\tau_0$, the loaded mass diverges and the flow terminal speed becomes zero, while, in the limit of $cP_0/F = 2/3$, the loaded mass becomes zero and the terminal speed approaches $(3/8)c$, which is the terminal speed above the luminous flat disk under an approximation of the order of $(v/c)^1$. We also examine the case where heating exists, and find that the flow properties are qualitatively similar to the case without heating. ",
        "introduction": "Accretion disks are tremendous energy sources in the active universe: in young stellar objects, in cataclysmic variables, in galaxtic X-ray sources and microquasars, and in active galaxies and quasars (see Kato et al. 1998 for a review). In particular, in a supercritical accretion disk, the mass-accretion rate highly exceeds the critical rate, the disk local luminosity exceeds the Eddington one, and the mass loss from the disk surface driven by radiation pressure takes place.  Such a radiatively driven outflow from a luminous accretion disk has been extensively studied in the context of models for astrophysical jets by many researchers (Bisnovatyi-Kogan, Blinnikov 1977; Katz 1980; Icke 1980; Melia, K\\\"onigl 1989; Misra, Melia 1993; Tajima, Fukue 1996, 1998; Watarai, Fukue 1999; Hirai, Fukue 2001; Fukue et al. 2001; Orihara, Fukue 2003), as on-axis jets (Icke 1989; Sikora et al. 1996; Renaud, Henri 1998; Luo, Protheroe 1999; Fukue 2005a), as outflows confined by a gaseous torus (Lynden-Bell 1978; Davidson, McCray 1980; Sikora, Wilson 1981; Fukue 1982), or as jets confined by the outer flow or corona (Sol et al. 1989; Marcowith et al. 1995; Fukue 1999), and as numerical simulations (Eggum et al. 1985, 1988). In almost all of these studies, however, the disk radiation field was treated as an external field, and the radiation transfer was not solved.   The radiation transfer in the disk, on the other hand, has been investigated in relation to the structure of a static disk atmosphere and the spectral energy density from the disk surface (e.g., Meyer, Meyer-Hofmeister 1982; Cannizzo, Wheeler 1984; Shaviv, Wehrse 1986; Adam et al. 1988; Hubeny 1990; Ross et al. 1992; Artemova et al. 1996; Hubeny, Hubeny 1997, 1998; Hubeny et al. 2000, 2001; Davis et al. 2005; Hui et al. 2005; see also Mineshige, Wood 1990). In these studies, however, the vertical movement and mass loss were not considered. Moreover, the relativistic effect for radiation transfer was not considered in the current studies.   In this paper we first examine the radiatively driven vertical flow -- {\\it moving photosphere} -- in a luminous flat disk within the framework of radiation transfer in the subrelativistic regime of $(v/c)^1$, where we incorporate the effect of radiation drag up to the first order of the flow velocity. At the first stage, we ignore the gravity of the central object as well as the gas pressure, although internal heating is considered in a limited way. We then consider analytical solutions for such a radiative flow, and examine the flow properties.  In the next section we describe basic equations in the vertical direction. In section 3 we examine the radiative flow without internal heating, while the radiative flow with internal heating is discussed in section 4. The final section is devoted to concluding remarks. ",
        "conclusions": "In this paper we have examined the radiative flow from a luminous disk, while taking account of the radiative transfer, in the subrelativistic regime of $(v/c)^1$. The flow velocity, the radiation pressure distribution, and other quantities are analytically solved as a function of the optical depth for the cases without and with heating. At the flow base (disk ``inside''), where the flow speed is zero, the initial optical depth is $\\tau_0$ and the initial radiation pressure is $P_0$; in the usual accretion disk these quantities are determined in terms of the central mass, the mass-accretion rate, and the viscous process as a function of the radius. That is, the optical depth is determined by the disk surface density, the radiative flux is given by the central mass and the accretion rate, and the radiation pressure relates to the disk internal structure. At the flow top (disk ``surface''), where the optical depth is zero, the radiation fields are assumed to be those above a static flat disk with uniform intensity. In order to match this boundary condition, one relation is imposed on the boundary quantities, and the mass-loss rate $J$ and terminal speed $v_{\\rm s}$ are both determined by the initial conditions $\\tau_0$ and $P_0$, as eigenvalues. In particular, in order for the flow to exist, the initial radiation pressure is bound in some range: at the upper limit, the loaded mass diverges and the flow terminal speed becomes zero, while, at the lower limit, the loaded mass becomes zero and the terminal speed approaches $(3/8)c$, which is the terminal speed above the luminous disk under the approximation of the order of $(v/c)^1$.  Among various types of astrophysical jets and outflows, some are believed to be accelerated by continuum radiation, and to which the present mechanism can be applied. For example, in supersoft X-ray sources, including accreting white dwarfs, high-velocity mass outflows were reported (e.g., Cowley et al. 1998). The wind velocity measured in these objects is about 3000--5000~km~s$^{-1}$, which is of the order of the escape velocity of white dwarfs. In addition, the line profile depends on the inclination angle, and the P~Cyg profiles are seen. Hence, in supersoft X-ray sources, the outflows should blow off from the innermost accretion disk. However, the gas temperature of the inner disk is so high that the line-driven mechanism cannot operate, and therefore, mass outflows in these objects would be driven by continuum radiation.  In SS~433, a prototype of astrophysical jets, the mass-accretion rate highly exceeds the Eddington one, and twin jets with $0.26~c$ would be accelerated by radiation pressure (see, e.g., Cherepashchuk et al. 2005 for a recent obsevation). The luminosity of microquasar GRS~1915$+$105 exceeds the Eddington one ($L_{\\rm Edd}$) during its outburst. In quiscent phase, its luminosity never drop below $\\sim 0.3 L_{\\rm Edd}$ (Done, Gierli\\'nski 2004).  In broad absorption-line quasars (BAL QSOs), strong broad absorption lines in the UV resonance are always blueshiftd relative to the emission-line rest frame, which indicates the presence of outflows from the nucleus, with velocities as large as $0.2~c$. These moderately high-velocity outflows are supposed to be accelerated by line or continuum radiation (e.g., Murray et al. 1995; Proga et al. 2000; Proga, Kallman 2004). In these objects, the outflows must consist of dense clouds, or be shielded from the central strong source, in order to operate a line-driven mechanism. Otherwise, the material is highly ionized by the central source, and, instead of line-acceleration, continuum acceleration may work (see, e.g., Chelouche, Netzer 2003).  Finally, narrow-line Seyfert 1 galaxies and luminous quasars are also believed to harbor supercritically accreting black holes, and high-velocity outflows are reported as warm absorbers (cf. Blustin et al. 2005). In these objects, radiative acceleration by continuum may play a dominant role.     The radiative flow investigated in the present paper must be quite {\\it fundamental problems} for accretion disk physics and astrophysical jet formation, although the present paper is only the first step, and there are many simplifications at the present stage.  For example, we have ignored the gravitational field produced by the central object. This means that the flow considered in the present paper would correspond to normal plasmas in the super-Eddington disk, pair plasmas in the sub-Eddington disk, or dusty plasmas in the luminous disk. In other cases, or even for normal plasmas in the super-Eddington disk, the gravitational field would affect the flow properties. In particular, the gravitational field in the vertical direction is somewhat complicated (e.g., Fukue 2002, 2004), and the influence of the gravitational field is important. The effect of gravity will be considered in a seperate paper (Fukue 2005b).   We also ignored the gas pressure. This means that the radiation field is sufficiently intense. In general cases, where the gas pressure is considered, there usually appear sonic points (e.g., Fukue 2002, 2004), and the flow is accelerated from subsonic to supersonic. The cross section of the flow also has a similar influence. In this paper we consider a purely vertical flow, and the cross section of the flow is constant. If the cross section of the flow increases along the flow, the flow properties, such as the transonic nature, would be influenced.  In the present paper, we have considered the internal heating processes in a simple form. Since the heating prcesses should couple with the viscous process or other heating processes, such as a nuclear process, we should treat internal heating more carefully.   In order to take account of the effect of radiation drag, we consider the subrelativistic regime, where the terms of the first order of $(v/c)$ are retained. As a result, the terminal speed $(3/8)c$ for the subrelativistic regime appears in the extreme case. This is quite consistent with the previous knowledge. As is well known, the terminal speed above the luminous flat disk in the full relativity is $[(4-\\sqrt{7})/3]c \\sim 0.45c$ (Icke 1989). To consider highly relativistic cases, all of the terms up to $(v/c)^2$ should be retained (cf. Kato et al. 1998). A related problem is the boundary condition at the flow top. In the present paper, we impose the boundary condition such that the radiation fields at the flow top are those above the luminous flat disk. Rigorously speaking, this boundary condition should be modified, since the disk gas, itself, is now moving. In the highly relativistic case, the deviation would be serious. The exact boundary conditions are derived and discussed in a separate paper for a fully special relativistic case (Fukue 2005c).  \\vspace*{1pc}  This work has been supported in part by a Grant-in-Aid for the Scientific Research (15540235 J.F.) of the Ministry of Education, Culture, Sports, Science and Technology."
    },
    "0601/astro-ph0601375_arXiv.txt": {
        "abstract": "In this short communication I compare recent findings suggesting a low binary star fraction for late type stars with knowledge concerning the forms of the stellar initial and present day mass functions for masses down to the hydrogen burning limit.  This comparison indicates that most stellar systems formed in the Galaxy are likely single and not binary as has been often asserted.  Indeed, in the current epoch two-thirds of all main sequence stellar systems in the Galactic disk are composed of single stars.  Some implications of this realization for understanding the star and planet formation process are briefly mentioned. ",
        "introduction": "\\label{sec:introduction}  Ever since Mitchell (1767) pointed out that the observed frequency of visual double stars was too high to be due to random chance, the study of binary stars has occupied an important place in astrophysics.  William Herschel (1802) discovered and cataloged hundreds of visual pairs and produced the first observations of a rudimentary binary orbit.  In doing so he established that the double stars were indeed physical pairs and that Newtonian physics operated nicely in the distant sidereal universe.  By the beginning of the twentieth century tens of thousands of binary stars were known and cataloged (e.g., Burnham 1906).  By the middle to late twentieth century the first systematic attempts to establish the binary frequency of main sequence F and G stars suggested that a very high fraction (70 - 80\\%) of all such stellar systems consist of binary or multiple stars (Heintz 1969; Abt \\& Levy 1976; Abt 1983).  The most comprehensive and complete study of the multiplicity of G stars was performed by Duquennoy \\& Mayor (1991) who argued that two-thirds of all such stellar systems are multiple.  It has often been assumed but never clearly demonstrated that similar statistics applied to stars of all spectral types.  This assumption has led to the commonly held opinion that most all stars form in binary or multiple systems with the Sun (and its system of planets) being atypical as a single star. But how robust is the assumption that the binary statistics for G stars is representative of all stars?  Over the last decade two important developments have occurred in stellar research which directly bear on this question.  First, the functional form of the stellar initial mass function (IMF) has been better constrained by observations of both field stars (e.g., Kroupa, 2002) and young embedded clusters (e.g., Muench et al.  2002).  The IMF has been found to peak broadly between 0.1 - 0.5 \\msun, indicating that most stars formed in the Galactic disk are M stars.  Second, surveys for binary stars have suggested that the binary star frequency may be a function of spectral type (e.g., Fischer \\& Marcy 1992).  In particular, there have been a number of attempts to ascertain the binary frequency of M type stars and even for L and T dwarfs, objectss near and below the hydrogen burning limit.  These studies suggest that the binary frequency declines from the G star value, being only around 30\\% for M stars (e.g., Leinert et al.  1997; Reid \\& Gizis 1997; Delfosse et al.  2004; Siegler et al.  2005) and as much as a factor of 2 lower for L and T dwarfs (e.g., Gizis et al.  2003).  I argue in this communication that these two facts together suggest that most stellar systems in the Galaxy consist of single rather than binary or multiple stars. ",
        "conclusions": "\\label{sec:discussion}    The primary result of this paper is the recognition that most stellar systems in the Galaxy consist of single rather than binary stars.  This fact has important consequences for star and planet formation theory.  For example, contrary to the current accepted paradigm that most, if not all, stars form in binary or multiple systems (e.g., Larson 1972, 2001; Mathieu 1994), this result could indicate that the theoretical frameworks developed to explain the formation of single, sunlike stars (e.g., Shu, Adams \\& Lizano 1987) have wide applicability. Indeed, when appropriately modified for a cluster-forming environment (e.g., Myers 1998; Shu, Li \\& Allen 2004), they may even describe most star forming events in the Galaxy.  On the other hand, most stars could still initially form in binary or multiple systems provided that most such systems promptly disintegrate via dynamical interactions or decay in an early, perhaps even protostellar, stage of evolution (e.g., Kroupa 1995; Sterzik \\& Durisen 1998, Reipurth 2000).  The current paradigm that most, if not all stars, form in binaries was strengthened by early multiplicity surveys of pre-main sequence (PMS) stars.  In particular, surveys of the PMS population of the Taurus cloud indicated a binary fraction that was twice that of field G stars (Ghez et al.  1993; Leinert et al. 1993; Reipurth \\& Zinnecker 1993).  However, most field stars are now known to have formed in embedded clusters, environments quite different than represented by the Taurus PMS population (e.g., Lada \\& Lada 2003).  Binary surveys of both young embedded and Galactic clusters have revealed binary fractions indistinguishable from that of the field (e.g., Petr et al.  1998; Duch\\^ene, Bouvier \\& Simon 1999; Patience \\& Duch\\^ene 2001). The most simple and straightforward hypothesis to explain these two facts and the finding of a high SSF in this paper is that the most common outcome of the star formation process is a single rather than multiple star.   Observations of dust emission and extinction of molecular cloud cores have found that the shape of the primordial or dense core mass function is very similar to that of the stellar IMF except that the core mass function is offset to higher mass by a factor of 2-3 (e.g., Stanke et al.  2005, Alves, Lombardi \\& Lada 2005).  These observations indicate that a 1-to-1 mapping of core mass to stellar mass, modified by a more or less constant star formation efficiency of 30-50\\%, is possible, if not likely.  This idea is consistent with single star systems being most often produced once the cores undergo collapse.  The fact that stellar multiplicity is a function of stellar mass, however, may provide important clues to the nature of the physical process of star formation.  For example, Durisen, Sterzik \\& Pickett (2001) have shown that if individual protostellar cores can further fragment and produce small N clusters, the dynamical decay of these clusters into binary and single stars can in certain circumstances produce a binary star fraction that declines with decreasing primary mass, similar to what is observed.  However, to be consistent with the SSF derived here and to simultaneously produce reasonable binary component separations, such models would require N $\\geq$ 5, within a region $\\sim$ 300 AU in size (Sterzik \\& Durisen 1998).  This would correspond to a stellar surface density ($\\sim$ 7.5 $\\times$ $10^5$ stars pc$^{-2}$) about two orders of magnitude higher than the peak density (7.2 $\\times$ 10$^3$ stars pc$^{-2}$) measured for the rich Trapezium cluster (Lada et al.  2004).  Such ultra-dense protostellar groups have not yet been identified, but could be revealed with high resolution infrared imaging surveys of deeply embedded candidates.  A related possibility, proposed by Kroupa (1995) and collaborators, posits that all stars are formed in binaries in modestly dense embedded clusters.  Dynamical interactions between these systems can disrupt some binaries and modify the separations of others.  These models can produce the observed dependance of binary frequency with mass, but at the expense of a SSF (50\\%) that is too low to be consistent with that derived here. These models could be made consistent with the high Galactic  SSF by assuming more compact configurations for the birth clusters, however it is unclear whether the required higher cluster densities would remain consistent with observed values.  Another possibility is that binary star formation is related to the initial angular momentum content of the primordial cores.  In this case the initial angular momentum of a protostellar core would be expected to be a function of core mass, with low mass cores being endowed with considerably less angular momentum than high mass cores.  A systematic molecular-line survey of cores of varying mass within a molecular cloud could test this idea.  A related possibility is that turbulence may play a role in the propensity for a core to fragment.  For example, Shu, Li \\& Allen (2004) posit that the break in the stellar IMF at 0.5 \\msun\\ is a result of the transition from turbulent to thermal support of the envelopes of dense pre-collapse cloud cores.  The more massive the core, the more turbulence is required to insure its support. Ammonia observations of dense cores in fact do suggest that massive cores are more turbulent than low mass cores (Jijina, Myers \\& Adams 1999).  Perhaps increased cloud turbulence in the more massive dense cores can also promote, in some fashion, more efficient core fragmentation and a higher incidence of binary star formation.  In this context it would be interesting to know if the trend of increasing stellar multiplicity with stellar mass continues to the more massive A, B and O stars, as has been suggested in some studies (e.g., Preibisch, Weigelt, \\& Zinnecker 2001, Shatsky \\& Tokovinin 2002).  Finally I note that the large fraction of single star systems in the field is consistent with the idea that most stars could harbor planetary systems unperturbed by binary companions and thus extra-solar planetary systems that are characterized by architectures and stabilities similar to that of the solar system could be quite common around M stars, provided planetary systems can form around M stars in the first place.  \\vskip -0.1in"
    },
    "0601/astro-ph0601696_arXiv.txt": {
        "abstract": "Structure and spectral energy distributions (SEDs) of externally irradiated circumstellar disks are often computed on the basis of the two-temperature model of Chiang \\& Goldreich. We refine these calculations by using a more realistic temperature profile which is continuous at all optical depths and thus goes beyond the two-temperature model. It is based on the approximate solution of the radiation transfer in the disk obtained from the frequency-integrated moment equations in the Eddington approximation. We come up with a simple procedure (``constant $g_z$ approximation'') for treating the vertical structure of the disk in regions where its optical depth to stellar radiation is high. This allows us to obtain expressions for the vertical profiles of  density and pressure at every point in the disk and to determine the shape of its surface. Armed with these analytical results we calculate the full radial structure of the disk and demonstrate that it favorably agrees with the results of direct numerical calculations. We also describe a simple and efficient way of the SED calculation based on our adopted temperature profile. Resulting spectra provide very good match (especially at short wavelengths) to the results of more detailed (but also more time-consuming) SED calculations solving the full frequency- and angle-dependent radiation transfer within the disk. ",
        "introduction": "\\label{sect:intro}   Protoplanetary disks around T Tauri stars (stellar mass $M_\\star\\lesssim 2~M_\\odot$) as well as the disks around more massive Herbig Ae/Be stars ($M_\\star\\gtrsim 2~M_\\odot$) belong to the class of circumstellar disks that are strongly affected by the radiation of their central stars. Irradiation modifies both the vertical structure of the disk and the radial dependence of disk properties. As a result, observational signatures of centrally irradiated discs are quite different from those of the conventional viscously heated $\\alpha$-disks (Shakura \\& Sunyaev 1972).  Kenyon \\& Hartmann (1987) have pointed out that the observed infrared spectra of protoplanetary disks can be understood as resulting from passive reprocessing of the stellar radiation assuming that disk surfaces have flared geometry. Later Chiang \\& Goldreich (1997; hereafter CG97) have analytically derived hydrostatic, radiative equilibrium models of disks around T Tauri stars.  Their calculation was based on a {\\it two-temperature approximation} for the disk thermal structure in which the outer ``superheated'' layer has been assumed isothermal at the temperature characteristic for dust grains directly illuminated by unattenuated stellar radiation. The isothermal midplane region was found to have lower temperature set by the balance between the oblique stellar illumination of the disk surface and the isotropic infrared emission of the disk interior. The success of this simple model has led to its wide acceptance for interpreting the observations of circumstellar disks.  The goal of this paper is to improve our understanding of the circumstellar disk structure by going beyond the simple two-temperature approximation of CG97. Calvet \\etal (1991; hereafter C91) have studied the vertical radiative transfer within irradiated dusty circumstellar disks and obtained approximate analytical expression for the vertical temperature profile which is superior to the two-temperature approximation of CG97. In this work we employ this temperature distribution, which is reviewed in detail in \\S \\ref{subsect:rad_transfer}, to calculate the disk structure and spectrum. Under certain conditions it is possible to derive approximate analytical solutions for the vertical disk structure as demonstrated in \\S \\ref{sect:hydro_eq}. This allows us to efficiently compute the radial dependence of various disk properties in \\S \\ref{sect:pure_radial} and to compare them with purely numerical solutions. Finally, in \\S \\ref{sect:spectra} we demonstrate how the use of the realistic temperature profile of C91 improves the calculation of the spectral energy distribution (SED) of circumstellar disks.  Throughout this study we assume the major heating source of the disk to be the radiation of the central star of mass $M_\\star$, having radius $R_\\star$ and effective temperature $T_\\star$. Viscous heating is considered subdominant. Circumstellar disk extending from $R_{in}$ to $R_{out}$ is assumed to be geometrically thin thus allowing us to neglect radial transport of radiation. This simplifies radiative transfer within the disk and reduces it to a one-dimensional problem. Surface density of the disk $\\Sigma_0$ varies as \\ba \\Sigma_0(a)=\\Sigma_\\star\\left(\\frac{R_\\star}{a}\\right)^{\\delta}, \\label{eq:Sigma} \\ea where $a$ is the distance from the central star and $\\Sigma_\\star\\equiv \\Sigma(R_\\star)$. We use $\\delta=1$ in our calculations. Our results will be illustrated using three fiducial models of circumstellar disks around Herbig Ae/Be star, T Tauri star, and a brown dwarf, analogous to those adopted by Dullemond \\& Natta (2003; hereafter DN03). Properties of these models are summarized in Table \\ref{table} ($\\Sigma_1\\equiv\\Sigma_\\star(R_\\star/1~\\mbox{AU})^{\\delta}$ is the value of $\\Sigma_0$ at $1$ AU). ",
        "conclusions": "\\label{sect:discussion}   We have explored structure of the irradiated circumstellar disk based on a realistic temperature profile (\\ref{eq:tem_calvet}), which is more accurate than the two-temperature approximation of CG97. Our results on the disk SEDs presented in \\S  \\ref{sect:spectra} persuasively demonstrate that this approach is superior to that of CG97 or DDN as it naturally reproduces the spectral behavior at short wavelengths without making any artificial assumptions.  This improvement is significant because of the recently emerged interest to the structure of the SED around $\\lambda\\sim 2$ $\\mu$m. Many circumstellar disks exhibit excess emission in this band (Hillenbrand \\etal 1992) which has been generally interpreted as the evidence for the existence of the dust sublimation region in the inner parts of these disks (DDN). Examination of Figures \\ref{fig:AeBe_1} \\& \\ref{fig:TT} reveals that the use of the two-temperature approximation can significantly overestimate the contribution of the disk to the total flux in this band. As a result, one would {\\it underestimate} the amount of emission produced by the dust sublimation region and this can seriously affect the interpretation of the data. On the contrary, our approach reproduces realistic disk SEDs (obtained by DN03 using rather time-consuming procedure for solving full radiation transfer) very accurately at short wavelengths and, thus, can be relied upon when determining the emissivity of the dust sublimation region from the data. Note that our method of SED calculation based on equation (\\ref{eq:tem_calvet}) is no more computationally intensive than the two-temperature approximation of CG97 which is routinely used for analyzing  protoplanetary disk spectra.  Another goal of this study was to introduce and test the constant $g_z$ approximation (\\S \\ref{subsect:g_z}) for treatment of the irradiated circumstellar disk structure. We have demonstrated by comparison with numerical results (\\ref{sect:pure_radial}) that within its region of applicability [see equation (\\ref{eq:validity})] this approximation works very well. Its use has allowed us to obtain analytical solutions for the vertical disk structure and to determine the position of the disk surface as a function of distance from the star rather accurately through the disk and stellar parameters. Improved description of the density and temperature structure of gas in the disk photosphere at $\\tau_\\star\\sim \\tau_c$ obtained in \\S \\ref{subsect:g_z} allows one to improve the predictions of molecular line intensities compared to the two-temperature approximation [see Dullemond \\etal (2002) for comparison of the line intensities obtained using full radiation transfer and its simplified representation given by equation (\\ref{eq:tem_calvet})].  Finally, solutions presented in \\S \\ref{subsect:g_z} allow fast and efficient exploration of the large phase space of circumstellar disk parameters when fitting disk SEDs. Their analytical nature will also be useful for  treating more complicated problems, e.g. determining the structure of the dust sublimation region, ice lines, and so on.\\\\"
    },
    "0601/astro-ph0601028.txt": {
        "abstract": "Current plans call for the first Terrestrial Planet Finder mission, TPF-C, to be a monolithic space telescope with a coronagraph for achieving high contrast.  The coronagraph removes the diffracted starlight allowing the nearby planet to be detected.  In this paper, we present a model of the planet measurement and noise statistics.  We utilize this model to develop two planet detection algorithms, one based on matched filtering of the PSF and one using Bayesian techniques.  These models are used to formulate integration time estimates for a planet detection with desired small probabilities of false alarms and missed detections. ",
        "introduction": "\\label{sec:intro}  In 2004, NASA  announced that it plans to launch the Terrestrial Planet Finder-C (TPF-C), a space based visible light telescope equipped with a coronagraph to detect earthlike planets around other stars, by 2016.  Work is proceeding on developing a baseline design that would meet the scientific requirements.  This currently consists of a $8 \\times 3.5$ meter off-axis Cassegrain optical telescope combined with a starlight suppression system (SSS)  to achieve the needed starlight rejection.  The SSS is equipped to utilize either conventional Lyot coronagraphs or pupil masks (shaped or apodized).  At Princeton, we have been studying  what we call shaped pupil coronagraphs as an alternative to the traditional Lyot design. \\citep{ref:KVSL,VSK02,VSK03,Kasdin04,Van04}  These are coronagraphs with apodized entrance pupils that rely solely on one/zero binary openings, resulting in a more manufacturable and robust design.  Critical to TPF-C's success is the ability to survey as many stars as possible and maximize the chances of detecting potential planets.  To do this, a thorough understanding of the data analysis is required in order to obtain reasonable estimates of the integration time for each observation.  \\cite{Brown05} discusses at length the question of survey completeness and its relation to integration time.  In this paper, we focus on a more careful and complete model of the measurement and noise statistics and develop a probabilistic detection criteria.  We also analyze a Bayesian detection approach to see if detection time improvements can be made.  These expressions can be used for further completeness studies and estimates of the efficacy  of the TPF-C search program. ",
        "conclusions": ""
    },
    "0601/hep-th0601162_arXiv.txt": {
        "abstract": "Using the recently introduced method to calculate bubble abundances in an eternally inflating spacetime, we investigate the volume distribution for the cosmological constant $\\Lambda$ in the context of the Bousso-Polchinski landscape model.  We find that the resulting distribution has a staggered appearance which is in sharp contrast to the heuristically expected flat distribution.  Previous successful predictions for the observed value of $\\Lambda$ have hinged on the assumption of a flat volume distribution.  To reconcile our staggered distribution with observations for $\\Lambda$, the BP model would have to produce a huge number of vacua in the anthropic range $\\Delta\\Lambda_A$ of $\\Lambda$, so that the distribution could conceivably become smooth after averaging over some suitable scale $\\delta\\Lambda\\ll\\Delta\\Lambda_A$. ",
        "introduction": "The cosmological constant problem is one of the most intriguing mysteries that we now face in theoretical physics. The observed value of the cosmological constant $\\Lambda$ is many orders of magnitude smaller than theoretical expectations and is surprisingly close to the present matter density of the universe,\\footnote{Here and below we use reduced Planck units, $M_p^2/8\\pi =1$, where $M_p$ is the Planck mass.} \\beq \\Lambda_0\\sim\\rho_{m0}\\sim 10^{-120}. \\label{Lambdarho} \\eeq As of now, the only plausible explanation for these enigmatic facts has been given in terms of the multiverse picture, which assumes that $\\Lambda$ is a variable parameter taking different values in different parts of the universe \\cite{Weinberg87,Linde87,AV95,Efstathiou,MSW,GLV,Bludman,AV05}. The probability for a randomly picked observer to measure a given value of $\\Lambda$ can then be expressed as \\cite{AV95} \\be P_{obs}(\\Lambda)\\propto P(\\Lambda)n_{obs}(\\Lambda), \\label{Pobs} \\ee where $P(\\Lambda)$ has the meaning of the volume fraction of the regions with a given value of $\\Lambda$ and $n_{obs}(\\Lambda)$ is the number of observers per unit volume.\\footnote{$P(\\Lambda)$ is often called the prior probability. Here we avoid this terminology, since it is usually used to characterize one's ignorance or prejudice, while the volume factor $P(\\Lambda)$ should be calculable, at least in principle.}  Disregarding the possible variation of other ``constants'' and assuming that the density of observers is roughly proportional to the fraction of matter clustered in large galaxies, \\be n_{obs}(\\Lambda) \\propto f_G(\\Lambda), \\label{nobs} \\ee one finds that the function $n_{obs}(\\Lambda)$ is narrowly peaked around $\\Lambda =0$, with a width \\beq \\Delta\\Lambda_A\\sim 100\\Lambda_0 \\sim 10^{-118}. \\label{DeltaLambdaA} \\eeq  In general, the volume factor $P(\\Lambda)$ depends on the unknown details of the fundamental theory and on the dynamics of eternal inflation.  However, it has been argued \\cite{AV96,Weinberg96} that it should be well approximated by a flat distribution, \\be P(\\Lambda)\\approx {\\rm const}. \\label{flat} \\ee The reason is that the anthropic range (\\ref{DeltaLambdaA}), where the function $n_{obs}(\\Lambda)$ is substantially different from zero, is much narrower than the full range of variation of $\\Lambda$, which is expected to be set by the Planck scale.  A smooth function varying on this large characteristic scale will be nearly constant within the tiny anthropic interval.  Combination of Eqs.~(\\ref{Pobs})-(\\ref{flat}) yields the distribution \\be P_{obs}(\\Lambda)\\propto f_G(\\Lambda), \\label{PnG} \\ee which can be readily calculated using the Press-Schechter approximation for $f_G$. The observed value of $\\Lambda$ is within the $2\\sigma$ range of this distribution -- an impressive success of the multiverse paradigm.  One should keep in mind, however, that the successful prediction for $\\Lambda$ hinges on the assumption of a flat volume distribution (\\ref{flat}).  We emphasize that the form of the volume distribution is important. If, for example, one uses $P(\\Lambda) \\propto\\Lambda$ instead of (\\ref{flat}), the $2\\sigma$ prediction would be $10 <\\Lambda/\\Lambda_0 <500$ and the observed value of $\\Lambda$ would be ruled out at a 99.9\\% confidence level \\cite{Pogosian}.  The heuristic argument for a flat distribution (\\ref{flat}) sounds plausible, but it needs to be verified in specific models.  The simplest model with a variable effective cosmological constant is that of a scalar field $\\phi$ with a very slowly varying potential $V(\\phi)$ \\cite{Banks84,Linde87}. In such models, $\\Lambda$ takes a continuum range of values. It has been verified that the resulting volume distribution for $\\Lambda$ is indeed flat for a wide range of potentials \\cite{Weinberg00,GV00,GV03}. The main challenge one has to face in this type of model is to justify the exceedingly flat potential which is required to keep the field $\\phi$ from rolling downhill on the present Hubble time scale.  A model with a discrete spectrum of $\\Lambda$ was first suggested by Abbott \\cite{Abbott}. He considered a scalar field with a ``washboard'' potential, having many local minima separated by barriers. Transitions between the minima could occur through bubble nucleation.  An alternative discrete model, first introduced by Brown and Teitelboim \\cite{BT}, assumes that the cosmological constant can be expressed as \\be \\Lambda = \\Lambda_{bare} +F^2/2. \\label{BT} \\ee Here, $\\Lambda_{bare}$ is the bare cosmological constant, which is assumed to be large and negative, and $F$ is the magnitude of a four-form field, which can change its value through the nucleation of branes. The change of the field strength across the brane is \\be \\Delta F =\\pm q, \\label{DeltaF} \\ee where the ``charge'' $q$ is a constant fixed by the model.  In order to explain observations, the discrete spectrum of $\\Lambda$ has to be very dense, with separation between adjacent values \\beq \\Delta\\Lambda\\lesssim \\Lambda_0, \\label{dense} \\eeq which in turn requires that the charge $q$ has to be very small. If this is satisfied, analysis shows that the flat volume distribution (\\ref{flat}) is quite generic \\cite{GV01}. But once again, the exceedingly small charge $q$ required by the model appears to be unnatural.\\footnote{Some ideas on how such small parameters could arise in particle physics have been suggested in \\cite{Weinberg00,Donoghue00,FMSW,BDM,DV01}.}  In an effort to remedy this problem, Bousso and Polchinski (hereafter BP) extended the Brown-Teitelboim approach to include multiple four-form fluxes \\cite{BP}. They considered a model in which $J$ different fluxes give rise to a J-dimensional grid of vacua, each labeled by a set of integers $n_a$. Each point in the grid corresponds to a vacuum with the flux values $F_a = n_a q_a$ and a cosmological constant \\be \\label{totalLambda} \\Lambda=\\Lambda_{bare}+\\frac{1}{2}\\sum_{a=1}^{J}F_a^2=\\Lambda_{bare} +\\frac{1}{2}\\sum_{a=1}^{J}n_a^2q_a^2. \\ee This model is particularly interesting because multiple fluxes generally arise in string theory compactifications. The model can thus be regarded as a toy model of the string theory landscape. BP showed that with $J\\sim 100$, the spectrum of allowed values of $\\Lambda$ can be sufficiently dense, even in the absence of very small parameters, e.g., with $|\\Lambda_{bare}|\\sim 1$, $q_a\\sim 0.1$.  In the cosmological context, high-energy vacua of the BP grid will drive exponential inflationary expansion. The flux configuration in the inflating region can change from one point on the grid to the next through bubble nucleation. Bubbles thus nucleate within bubbles, and each time this happens the cosmological constant either increases or decreases by a discrete amount. This mechanism allows the universe to start off with an arbitrary large cosmological constant, and then to diffuse through the BP grid of possible vacua, to generate regions with each and every possible cosmological constant, including that which we inhabit.  Our goal in this paper is to study the volume distribution for $\\Lambda$ in the BP model. In particular, we would like to check whether or not this distribution is approximately flat, as suggested by the heuristic argument of \\cite{AV96,Weinberg96}. Until recently, such an analysis would have been rather problematic, since the calculation of the volume fractions in an eternally inflating universe is notoriously ambiguous. The volume of each type of vacuum diverges in the limit $t\\to\\infty$, so in order to calculate probabilities, one has to impose some kind of a cutoff. The answer, however, turns out to be rather sensitive to the choice of the cutoff procedure. If, for example, the cutoff is imposed on a constant time surface, one gets very different distributions depending on one's choice of the time variable $t$ \\cite{LLM}. (For more recent discussions, see \\cite{VVW,Guth00,Tegmark}.)  Fortunately, a fully gauge-invariant prescription for calculating probabilities has been recently introduced in \\cite{GSPVW}. It has been tried on some simple models and seems to give reasonable results. Here we shall apply it to the BP model.  As BP themselves recognized, their model does not give an accurate quantitative description of the string theory landscape. In particular, it does not explain how the sizes of compact dimensions get stabilized. This issue was later addressed by Kachru, Kallosh, Linde and Trivedi (KKLT) \\cite{KKLT}, who provided the first example of a metastable string theory vacuum with a positive cosmological constant. Apart from the flux contributions in (\\ref{totalLambda}), the vacuum energy in KKLT-type vacua gets contributions from non-perturbative moduli potentials and from branes.  The $4D$ Newton's constant in these vacua depends on the volume of extra dimensions and changes from one vacuum to another.  Douglas and collaborators \\cite{Douglas,AshokDouglas,DenefDouglas} studied the statistics of KKLT-type vacua. Their aim was to find the number of vacua with given properties (e.g., with a given value of $\\Lambda$) in the landscape. Our goal here is more ambitions: we want to find the probability for a given $\\Lambda$ to be observed.  In our analysis, we shall disregard all the complications of the KKLT vacua, with the hope that the simple BP model captures some of the essential features of the landscape. At the end of the paper we shall discuss which aspects of our results can be expected to survive in more realistic models.  We begin in the next section by summarizing the prescription of Ref.~\\cite{GSPVW} for calculating probabilities.  We shall see that the problem reduces to finding the smallest eigenvalue and the corresponding eigenvector of a large matrix, whose matrix elements are proportional to the transition rates between different vacua. The calculation of the transition rates for the BP model is reviewed in Section III.  Some of these rates are extremely small, since the upward transitions with an increase of $\\Lambda$ are exponentially suppressed relative to the downward transitions. In Section IV we develop a perturbative method for solving our eigenvalue problem, using the upward transition rates as small parameters. This method is applied to the BP model in Section V. We find that the resulting probability distribution has a very irregular, staggered appearance and is very different from the flat distribution (\\ref{flat}). The implications of our results for the string theory landscape are discussed in Section VI. ",
        "conclusions": "In this paper we have used the prescription of Refs. \\cite{GSPVW,ELM} to determine the bubble abundances $p_j$ in the BP model. We found that the resulting distribution is very irregular, with values of $p_j$ soaring and plummeting wildly as $\\Lambda_j$ changes from one value to the next. This distribution is drastically different from the flat distribution (\\ref{flat}) which was used as a basis for the successful anthropic prediction for $\\Lambda$.  Apart from the bubble abundance factor $p_j$, the volume distribution (\\ref{PpZ}) includes the slow-roll expansion factor $Z_j$. In any realistic model, bubble nucleation should be followed by a period of slow-roll inflation, at least in some bubble types. The expansion factor $Z_j$ is, of course, model-dependent, but there is no reason to expect that it will somehow compensate for the wild swings in the values of $p_j$ as we go from one value of $\\Lambda_j$ to the next.  Another point to keep in mind is that, in a realistic setting, vacua with different values of the fluxes $F_a$ may have different low-energy physics, so the density of observers $n_{obs}(\\Lambda)$ would also be very different. We should therefore focus on the subset of vacua in the BP grid which differ only by the value of $\\Lambda$ and have essentially identical low-energy constants. Once again, there seems to be no reason to expect any correlation between these constants and the up and down swings in the bubble abundances. We conclude that the staggered character of the distribution $P_j \\equiv P(\\Lambda_j)$ is expected to persist, even in more realistic versions of the model.  This conclusion is not limited to the BP model. It is likely to arise in any landscape scenario, where a dense spectrum of low-energy constants is generated from a wide distribution of states in the parameter space of the fundamental theory. Vacua with nearly identical values of $\\Lambda$ may then come from widely separated parts of the landscape and may have very different bubble abundances and volume fractions.  Given the staggered character of the volume distribution, what kind of prediction can we expect for the observed value of $\\Lambda$? The answer depends on the number ${\\cal N}_A$ of possible vacua with $\\Lambda_j$ within the anthropic range (\\ref{DeltaLambdaA}), $\\Delta\\Lambda_A \\sim 10^{-118}$. (We count only vacua in which all low-energy constants other than $\\Lambda$ have nearly the same values as in our vacuum.)  Suppose the volume factors in the distribution $P_j$ span $K$ orders of magnitude. ($K\\sim 300$ in our numerical example in Section V.) We can divide all vacua into, say, $10K$ bins, such that the values of $P_j$ in each bin differ by no more than $10\\%$. Suppose now that there are ${\\cal N}_A \\ll 10K$ vacua in the anthropic range $\\Delta\\Lambda_A$. We can then expect that most of these vacua will be characterized by vastly different volume factors $P_j$, so that the entire range will be dominated by one or few values of $\\Lambda_j$ having much higher volume fractions than the rest.  Moreover, there is a high likelihood of finding still greater volume fractions if we go somewhat beyond the anthropic range - simply because we would then search in a wider interval of $\\Lambda$.  We could, for example, find that a vacuum with $\\Lambda_1 \\sim 10^{-114}\\sim 10^6\\Lambda_0$ has a volume fraction 200 orders of magnitude greater than all other vacua in the range $0<\\Lambda\\lesssim \\Lambda_1$. Galaxy formation is strongly suppressed in this vacuum: the fraction of matter that ends up in galaxies is only $f_G(\\Lambda_1)\\sim 10^{-110}$.  However, this suppression is more than compensated for by the enhancement in the volume fraction.  If this were the typical situation, most observers would find themselves in rare, isolated galaxies, surrounded by nearly empty space, all the way to the horizon.  This is clearly not what we observe. The dominant value could by chance be very close to $\\Lambda=0$, but if such an ``accident'' is required to explain the data, the anthropic model loses much of its appeal.  In the opposite limit, ${\\cal N}_A\\gg 10K$, the number of vacua in the anthropic interval $\\Delta\\Lambda_A$ is so large that they may scan the entire range of $P_j$ many times. Then, it is conceivable that the distribution will become smooth after averaging over some suitable scale $\\delta\\Lambda$. If $\\delta\\Lambda$ can be chosen much smaller than $\\Delta\\Lambda_A$, then it is possible that the effective, averaged distribution will be flat, as suggested by the heuristic argument in the Introduction.  The successful prediction for $\\Lambda$ would then be unaffected.\\footnote{Joe Polchinski has informed us that a similar argument, indicating that the anthropic explanation for the observed $\\Lambda$ requires a large number of vacua in the anthropic range, was suggested to him by Paul Steinhardt.}  The above argument is somewhat simplistic, as it assumes that the vacua in the BP grid are more or less randomly distributed between the $10K$ bins, with roughly the same number of vacua in each bin. Such an ``equipartition'' is not likely to apply to the most abundant vacua, but it may hold for the vacua in the mid-range of $P_j$. Finding the conditions under which equipartition applies would require a statistical analysis that goes beyond the scope of the present paper.  In summary, it appears that the staggered volume distribution resulting from the BP model is in conflict with observations, unless it yields a huge number of vacua in the anthropic range of $\\Lambda$. Counting only vacua which have nearly the same low-energy physics as ours, we should have much more than $10K$; hence, the total number of vacua should be many orders of magnitude greater. The large number of vacua in the anthropic range is only a necessary condition for the distribution $P_j$ to average out to the flat distribution (\\ref{flat}). Further analysis will be needed to find whether or not this actually happens, and if so, then under what conditions.  It would also be interesting to analyze other simple models of the landscape, such as the ``predictive landscape'' of Arkani-Hamed, Dimopoulos and Kachru \\cite{AHDK}, and see what similarities and differences they have compared to the BP model.  Throughout this paper we assumed that the brane charges $q_a$ are not particularly small. This assumption may be violated in certain parts of the landscape, e.g., in the vicinity of conifold points, resulting in a much denser spectrum of vacua \\cite{FMSW,AshokDouglas,DenefDouglas,Giryavets}. Infinite accumulations of vacua may occur near certain attractor points \\cite{DVattractor,Gia}.  Implications of these effects for the probability distribution on the landscape remain to be explored."
    },
    "0601/astro-ph0601139_arXiv.txt": {
        "abstract": "We have detected absorption lines from the High Velocity Cloud \\wb\\ in the spectrum of the star \\tee.  This detection sets an upper distance limit to the cloud of $8.8^{+2.3}_{-1.3}\\kpc$.  Non-detection (at $>4\\sigma$ confidence) in the star \\hee\\ at $7.7\\pm 0.2\\kpc$ sets a probable lower limit. The equivalent width of the \\CaK\\ line due to the HVC (\\teeWCaK) corresponds to a column density of \\teeNCa.  Using an \\HI\\ spectrum from the Leiden/Argentine/Bonn survey, we calculated \\teeNCaHI. These distance limits imply an \\HI\\ mass limit of $3.8 \\times 10^5\\Msun < \\mbox{M}_{\\HI} < 4.9 \\times 10^5\\Msun$.  The upper distance limit imposed by these observations shows that this HVC complex has a probable Galactic or circum-Galactic origin. Future metallicity measurements will be able to confirm or refute this interpretation. ",
        "introduction": "High Velocity Clouds (HVCs) are clouds of neutral hydrogen (\\HI) gas with velocities that are inconsistent with a simple model of differential Galactic rotation. Since a kinematic distance cannot be derived, the exact location and nature of the HVCs has been the topic of much speculation. Determination of the distances to HVCs is vitally important, as such measurements can constrain their likely origins, and establish their physical parameters, many of which scale with distance.  The HVCs have been ascribed to many different origins, from local Galactic processes, to distant proto-galaxies. \\citet{shapiro-field-76} were the first to suggest the existence of a Galactic Fountain, in which supernovae drive hot gas into the halo, which then condenses and returns to the disk \\citep[e.g.][]{bregman-80}. Extra-Galactic origins have been discussed almost since the discovery of HVCs \\citep{oort-66-HVC-origins}. \\citet{blitz-etal-99} have inspired vigorous debate by reviving and further refining this scenario, noting that it naturally explains the kinematics of clouds towards the bary- and anti-barycenter of the Local Group. They suggested that these distant HVCs could be dark-matter dominated objects, which some have claimed might be the solution to the ``satellite problem'' of Cold Dark Matter simulations \\citep[namely, that many more dark matter haloes are predicted to exist than satellite galaxies observed, e.g. ][]{klypin-etal-99-missing-satellites}.  Despite many years of effort, few HVCs have been detected in absorption against stellar probes.  Of the traditional complexes, only M and A have been detected \\citep{danly-etal-93-ComplexM,vanwoerden-etal-99-nature}. Several other smaller clouds have also been detected in the spectra of background stars \\citep[e.g.][]{bates-etal-90-HVC-dist,sembach-savage-massa-91-HVC-dist}. \\citet{wakker-01-distance-metallicity} provides a much more comprehensive summary.  \\citet{wannier-wrixon-wilson-72} were the first to report on the positive velocity clouds in the third and fourth Galactic quadrants.  Fainter than the traditional negative velocity complexes such as A and C, they were divided into four separate complexes WA -- WD \\citep[][hereafter WvW91]{wvw91}. These complexes span a wide range of position and velocity space, comprising a total of 64 clouds in the WvW91 catalog, with mostly moderate deviation velocities\\footnote{Deviation velocity is the amount by which a velocity deviates from allowed Galactic rotation velocities \\citep{wakker-91-II-distribution}.}. \\wb\\ ($l = 225 - 265\\degr$; $b = 0 - 60\\degr$) is composed of 29 separate clouds, with a total \\HI\\ mass $\\mbox{M(\\HI)} = 0.236\\,d^{2}(\\kpc)\\,S({\\rm Jy}\\kms) \\sim 6.5\\times 10^{3}\\,d^{2}\\Msun$, where $S$ is the integrated flux and $d$ is the distance. \\citet{robertson-etal-91-complexWB} have observed the Wannier Clouds in absorption towards the Seyfert galaxy PKS\\,0837--12. Along this line of sight they derive a N(\\Ca)/N(\\HI) ratio of $160 \\times 10^{-9}$. Using a higher resolution \\HI\\ spectrum, \\citet{wakker-01-distance-metallicity} revised this value upwards to $280 \\times 10^{-9}$. Both sets of authors note the anomalously high calcium abundance (one of the highest measured for any HVC), while \\citeauthor{robertson-etal-91-complexWB} concluded that a distance is necessary to discern between the various possible explanations they offered.  The Wannier clouds, along with the extreme positive velocity clouds (population EP) lie in the direction of the Local Group anti-barycenter. They are explicitly included in the \\citet{blitz-etal-99} scenario as being very distant and undergoing infall towards the Local Group barycenter. Thus a confirmed upper limit on the distance would stand in opposition to this suggestion and may have broader consequences for the HVCs as a whole. ",
        "conclusions": "\\label{sec: discussion}   The upper distance limit we obtain places WB firmly inside the halo of the Milky Way. The non-detection towards \\hee\\ sets a probable lower limit. Although a scenario in which the two clouds are physically disjoint is not excluded, it is unlikely given the location and kinematics of the gas. A distance of $7.7\\kpc < \\mbox{d} < 8.8\\kpc$ implies a total \\HI\\ mass of $3.8 \\times 10^5\\Msun < \\mbox{M}_{\\HI} < 4.9 \\times 10^5\\Msun$ for the entire complex. We measured a \\Ca\\ abundance $N(\\Ca)/N(\\HI) = 81\\pm16\\times10^{-9}\\cm$ ($\\sim 0.04$ solar) which is about a factor of four less than that seen towards PKS\\,0837--120 \\citep[$280 \\times 10^{-9}$; ][]{wakker-01-distance-metallicity}, $\\sim 35\\degr$ away from \\tee.  Since the \\HI\\ data have poor spatial resolution (in the case of the LAB) and poor velocity resolution (HIPASS), we cannot resolve individual cloud structures; the \\HI\\ column density estimates are likely to be accurate only to a factor of two.  While the broad characteristics of the \\HI\\ and optical data match, we cannot resolve any condensations within the cloud, which are suggested by the multiple absorption components seen in the optical data. Higher resolution interferometric or large single-dish data (e.g. ALFA at Arecibo), with good velocity resolution, are thus crucial.  In the Galactic rest frame, the measured velocity of \\wb\\ (96\\kms) implies that the complex is in-falling with $v_{\\mbox{\\scriptsize GSR}} = -30\\kms$. A maximum distance of 9\\kpc\\ places \\wb\\ $7\\kpc$ above the disk of the Milky Way, some $12\\kpc$ from the Galactic center. Both fountain and infalling gas interpretations are consistent with the data.  Under an infall scenario, the anomalously high N(\\Ca)/N(\\HI) measurement might be explained in terms of low depletion due to low dust content or hydrogen ionization (which has not been taken into account here).  If the gas is closer to the disk, then \\wb\\ may be fountain gas returning to the disk. The \\Ca\\ abundance is not implausible since \\Ca\\ does not accurately trace the total gas metallicity, although a more accurate metallicity estimate would be very helpful for this discussion.  The distance bracket towards \\wb\\ is an important step towards the final goal of determining the distances to a large sample of High Velocity Clouds. It is one of only a few upper distance limits to HVCs, and the first direct limit towards a cloud that was predicted to be very distant by the \\citet{blitz-etal-99} model.  Our results add to the growing body of evidence that the High Velocity Clouds are much more likely to belong to a Galactic or circum-Galactic population, rather than a distant, extra-Galactic one."
    },
    "0601/astro-ph0601413_arXiv.txt": {
        "abstract": "We present measurements obtained with the {\\it Spitzer Space Telescope\\/} in five bands from 3.6--24\\micron\\ of the northern inner radio lobe of Centaurus~A, the nearest powerful radio galaxy. We show that this emission is synchrotron in origin.  Comparison with ultraviolet observations from {\\it GALEX\\/} shows that diffuse ultraviolet emission exists in a smaller region than the infrared but also coincides with the radio jet. We discuss the possibility, that synchrotron emission is responsible for the ultraviolet emission and conclude that further data are required to confirm this. ",
        "introduction": "The radio source Centaurus~A and its host galaxy, NGC\\,5128, provide a rare opportunity to observe the detailed behaviour of a recently merged system supporting a powerful active nucleus with jets and extended radio emission.  At a distance of 3.4\\,Mpc \\cite{israel98}, such that  1\\arcs = 16.5\\,pc, Centaurus~A is the nearest powerful radio galaxy, and its activity, merger remnants, and star-formation may be resolved and studied in detail .  The host galaxy is believed to be a giant elliptical galaxy that recently merged ($\\sim$200\\,Myrs ago) with a small spiral galaxy, producing the prominent dust lanes seen in optical images  \\citep{bm54,QGF93}.  The active nucleus gives rise to a jet and counter-jet in the inner arcminute ($1\\arcm \\approx 1$\\,kpc) about the nucleus \\citep[][]{feigelson81,hardcastle03}.  On larger scales, giant radio lobes extend over $\\sim$6\\deg\\ on the sky, with inner radio lobes extending about 6\\arcm\\ to the north-east and south-west.  The focus of this {\\em Letter} is the Northern radio `jet', \\squig~3--4\\,kpc from the nucleus (see Figure 1).  For a comprehensive review of Centaurus~A see, for example, Israel~(1998), Morganti et al.~(1999) and Junkes et al.~(1993).  While jet/lobe emission is broadly understood in terms of synchrotron emission from electrons accelerated in the nuclear jet or in associated shocks, there remain uncertainties when interpreting the details of observed emission.  Observing the synchrotron spectrum, particularly at sufficiently high frequencies that the cut-off due to spectral aging is measured, allows the study of the energy distribution of the underlying electron population.  Since the radio emission is produced by long-lived electrons, it alone does not provide detailed information about the underlying physical structure of the emitting plasma.  At sites of acceleration, highly energetic electrons may emit at frequencies as high as X-rays \\citep[][]{smith83,feigelson81}, so a multi-frequency approach is required to study these regions.  Infrared (IR) observations provide important constraints at intermediate frequencies in modelling targets.  The {\\it Spitzer Space Telescope\\/} \\citep[][]{werner04} offers a powerful new capability to study IR emission from jets in \\hbox{AGN}.  The Infrared Array Camera \\citep[IRAC;][]{fazio04} and the Multiband Imaging Photometer \\citep[MIPS;][]{rieke04} provide imaging at 3.6 to 160\\microns, with sensitivity orders of magnitude better than any other telescope. Whilst  near-IR detections of FRII sources have been made recently (e.g. Floyd et al.~2006), no jet to date has been observed extensively in the IR.  The {\\it Spitzer Space Telescope\\/}'s advantage, in addition to offering several IR bands, is its sensitivity. In this {\\it Letter} we present IRAC and MIPS photometry of the northern radio lobe, based upon  effective integrations of 72\\,s and 160\\,s respectively, and use them to describe and interpret the jet SED of Centaurus \\hbox{A}.  The IRAC observations and general processing are described in Quillen et al.\\ (2006).  The MIPS data are presented here for the first time.  We use radio data at 843\\,MHz, 1.4\\,GHz, and 4.9\\,GHz from the literature:  the 843\\,MHz data are taken from the Sydney University Molonglo Sky Survey \\citep[SUMSS;][]{sumss};  the 1.4\\,GHz and 4.9\\,GHz data are from Condon et al.\\ (1996) and Burns et al.\\ (1983), respectively.   We show that the diffuse IR emission is also synchrotron in origin, and compare it to ultraviolet (UV) emission detected by the {\\it Galaxy Evolution Explorer\\/} \\citep[{\\it GALEX\\/};][]{martin05}.   \\S\\,\\ref{observations} presents the MIPS imaging and describes the IRAC, UV, and radio data sets.  \\S\\,\\ref{measure} describes the method for measuring the surface brightness of the jet and presents the derived photometry.  \\S\\,\\ref{discussion} interprets the results in terms of an underlying synchrotron spectrum and discusses the physical implications. ",
        "conclusions": "\\label{discussion}  Figure \\ref{bgfit_jetsed} (right) shows the surface brightness of the designated region of the jet of Centaurus~A as a function of frequency.  The solid line is a power law, $S\\propto \\nu^{-\\alpha}$, fit to the IR points for which $\\alpha$ = 0.68. This line comes remarkably close to the radio points suggesting that the IR emission is also synchrotron in origin, though the radio points are slightly flatter than the IR points. This suggests the possibility of a break in a synchrotron spectrum, rather than a simple, single power-law spectrum. However this is not a strong conclusion. We conclude that the IR emission is due to synchrotron emission and, when a single power-law is fitted to both the radio and IR points, $\\alpha$ = 0.72 (the dashed line in Figure \\ref{bgfit_jetsed}).  Synchrotron emitting electrons radiate at a rate proportional to their energy.  The most energetic electrons lose their energy the fastest, leading to a slope that is one-half power steeper in the high frequency spectrum,  at frequencies much above the break frequency, $\\nu_B = eB/2\\pi m_{\\rm e}c$ (e.g. Beams and Jets in Astrophysics, ed. Hughes, 1991). Our observations indicate that this break must occur at frequencies higher than $10^{13}$\\,Hz.   Though detailed modelling is required for an exact evaluation of the break frequency, an approximate limit may be gained.  Following Miley (1980), we calculate the magnitude of the magnetic field on the assumption that the system is at its minimum energy density.   Using Equation~3 of Miley (1980) at 1.4\\,GHz, we find $B\\approx3$\\,\\hbox{nT}.  Using the highest IR frequency observed (83\\,THz), an upper limit to the lifetime of the emitting electrons is found to be 30,000\\,yrs.  It is important to know how far the synchrotron spectrum extends without a break, since very high frequency emission is produced by short-lived, high energy particles and may imply in situ particle acceleration. In both the near UV and far UV {\\it GALEX\\/} images there are two kinds of emission: extended sources, likely associated with star-formation sites; and diffuse emission which coincides with the jet. Figure \\ref{overlayuv_plot} shows the 1.4\\,GHz radio contours overlaid on the  8\\micron\\ IRAC image in greyscale (left) and  the near UV {\\it GALEX\\/} image (right; smoothed with a 7\\farcs6 circular Gaussian for clarity). The UV emission is seen only at the inner portion of the jet, whereas the IR emission follows the jet to the north, consistent with the idea that emission at all three wavelengths is synchrotron emission.  The synchrotron lifetime of the UV emitting electrons is shorter than that of the IR emitting electrons, and so they do not live/emit long enough to trace the bulk motion downstream. northward at this point, indicative of interactions with the environment in the region.  \\begin{figure*} \\plotone{overlayuv_fig.ps} \\figcaption{ Left: 8\\microns\\ IRAC image of the jet with contours describing the 1.4\\,GHz radio emission. Right: Near UV {\\it GALEX\\/} image (smoothed by a 7\\farcs6 circular Gaussian) with 1.4\\,GHz radio contours.  The two sources to the North and East are discrete and are likely to be associated with star-formation sites. Both images show the region of the aperture used to measure the surface brightness (Figure \\ref{mipsirac}). \\vspace*{1.0cm} \\label{overlayuv_plot} } \\end{figure*}  The surface brightness of the diffuse UV emission is compared to that of the radio and IR emission in Figure \\ref{bgfit_jetsed} (right).  The UV points are significantly below the extrapolation of the synchrotron emission based upon the radio and IR data.  Several explanations must be considered. Firstly, the diffuse UV emission might not be synchrotron at all. We consider this unlikely because the jet emission extends  to X-ray frequencies \\citep[][]{kraft03}, though it is possible that other mechanisms, such as diffuse star-formation, also contribute to the UV. Secondly, we could be seeing the effects of extinction due to dust. Suppose that the best-fit power law to the radio and infrared measurements extends without break to ultraviolet frequencies, but that there is intervening dust.  Assuming a Calzetti extinction law, only $E(B-V) = 0.56$ is required to produce  the observed ultraviolet emission. Given the violent dynamic history of Centaurus~A there is clearly a large presence of gas, and therefore dust, associated with the merger, which can be found in distinct regions far outside the galactic disk. Hence it is plausible that  sufficient extinction, at the shortest wavelengths, exists along specific lines of sight. Note that  if this were the case, the appropriate synchrotron lifetime for the UV emitting particles (\\squig~~6000 years) would be short enough to require {\\it in situ} particle acceleration, given  a projected distance of $\\sim$3\\,kpc from the nucleus, and bulk motion $\\sim$0.5\\,c in the kpc scale jet \\citep{hardcastle03}.  However, inclusion of X-ray measurements argues for a third possibility: a spectral break between infrared and ultraviolet frequencies.  Hardcastle et al.\\ (2006; henceforth H06) fit a broken power-law spectrum to all measurements from radio to X-ray frequencies, and determine that a break occurs at $\\sim$0.3\\,THz. Figure 5 of H06 shows that this broken power law accounts well for the radio, infrared, and X-ray measurements.  However, it significantly underpredicts the ultraviolet measurements, even assuming no extinction from dust intrinsic to Centaurus A (Galactic extinction was included).  If intrinsic extinction is important this discrepancy becomes worse.  Neither dust alone, nor a broken power-law accounts for the data in an entirely satisfactory way. Submillimeter and 70\\micron\\ {\\it Spitzer\\/} measurements with higher SNR would fill in the gap in the observed spectrum above and below 1\\,THz, providing significant new constraints on the overall spectrum, and would help in resolving the issue."
    },
    "0601/astro-ph0601625_arXiv.txt": {
        "abstract": "Magnetic fields play an important role in producing and modifying the photospheric chemical peculiarities of intermediate-mass main sequence stars. This article discusses the basic theory and methods of measurement used to detect and characterise stellar magnetic fields, and reviews our current knowledge of selected characteristics of magnetic fields in intermediate-mass stars. ",
        "introduction": "Magnetic fields play a fundamental role in the physics of the atmospheres of a significant fraction of stars on the H-R diagram. The magnetic fields of early-type stars (mainly the intermediate-mass main sequence Ap and Bp stars) have quite different characteristics (and probably a different origin) than those of late-type stars like the sun (e.g. Mestel 2003). In early-type stars, the surface magnetic field is static on a timescale of at least many decades, and appears to be \"frozen'' into a rigidly rotating atmosphere. The magnetic field is organized on a large scale, permeating the entire stellar surface, with a relatively high field strength (typically of a few hundreds up to a few tens of thousands of gauss). The presence of a magnetic field strongly influences energy and mass transport within the atmosphere, and results in the presence of strong chemical abundance nonuniformities in photospheric layers (one of our reasons for having this meeting!). Due its global nature, the magnetic field of Ap stars can be much more easily detected and studied than that of late-type stars. Therefore, Ap stars represent a primary target for spectropolarimetric observations, for the development and testing of the instruments and modelling techniques, and for evaluation of physical models.ssssssssss ",
        "conclusions": "Using the same basic physical principles demonstrated by Zeeman (1897), modern spectropolarimetric methods are being combined with sophisticated data analysis techniques to clarify the magnetic properties of intermediate mass stars. Some of the most important results of recent investigations are: \\begin{enumerate} \\setlength{\\itemsep}{-4pt} \\item When viewed at high resolution using Magnetic Doppler Imaging, the surface magnetic field topologies of A and B stars show significant diversity (see Kochukhov, this volume). \\item A ``magnetic dichotomy'' is confirmed to exist for main sequence A and B star (see Shorlin et al. 2002, Auri\\`ere et al., this volume). \\item Magnetic Ap/Bp stars of low mass (2-3~$M_\\odot$) have been firmly established at very early phases of their main sequence evolution (see Pohnl et al. 2003, 2005 Bagnulo et al. 2005a, and in preparation). \\item The first magnetic pre-main sequence progenitors of magnetic Ap/Bp stars appear to have been identified (see Drouin 2005, Wade et al. 2005). \\end{enumerate}"
    },
    "0601/astro-ph0601555_arXiv.txt": {
        "abstract": "We describe 2D gasdynamical models of jets that carry mass as well as energy to the hot gas in galaxy clusters. These flows have many attractive attributes for solving the galaxy cluster cooling flow problem: Why the hot gas temperature and density profiles resemble cooling flows but show no spectral evidence of cooling to low temperatures. Using an approximate model for the cluster A1795, we show that mass-carrying jets can reduce the overall cooling rate to or below the low values implied by X-ray spectra. Biconical subrelativistic jets, described by several {\\it ad hoc} parameters, are assumed to be activated when gas flows toward or cools near a central supermassive black hole. As the jets proceed out from the center they entrain more and more ambient gas. The jets lose internal pressure by expansion and are compressed by the ambient cluster gas, becoming rather difficult to observe. For a wide variety of initial jet parameters and several feedback scenarios the global cooling can be suppressed for many Gyrs while maintaining cluster temperature profiles similar to those observed. The intermittancy of the feedback generates multiple generations of X-ray cavities similar to those observed in the Perseus Cluster and elsewhere. ",
        "introduction": "Many computationally elaborate calculations of jet-heated cooling flows have been proposed to explain why the hot gas in galaxy clusters cools much slower than the rate expected from the cluster X-ray luminosity, ${\\dot M}_{cf} \\approx  (2 \\mu m_p / 5 k T_{vir}) L_x$ (e.g. Cavaliere et al. 2001; Reynolds, Heinz, \\& Begelman 2001,2002; Bruggen \\& Kaiser 2002; McCarthy et al. 2003; Hoeft et al. 2003; Basson \\& Alexander 2003; Omma et al. 2004; dalla Vecchia et al. 2004; Zanni et al. 2005). In principle the energy created by gas accreting at rates very much less than ${\\dot M}_{cf}$ onto supermassive black holes in cluster-centered galaxies can easily balance the radiative losses $L_x$. Consequently, AGN jets have been regarded as a plausible means to distribute this energy throughout the cluster gas. Unfortunately, few if any jet-heated flow simulations closely resemble in projection the X-ray emission and radial temperature profiles observed in galaxy groups and clusters. Surprisingly often these calculations do not include radiative losses and therefore lack the essential physics to prove that cooling can be stopped by the jets. Nor do these calculations typically continue for many Gyrs. Multi-Gyr calculations are essential to determine if the proposed type of heating can keep the group/cluster gas from cooling while also preserving time-averaged temperature and entropy profiles commonly observed, most of which are very similar. In general, the long term effectiveness of powerful jet heating in stopping the cooling has not been convincingly demonstrated. As a first step toward a more successful solution, adopted here, jet-heating scenarios are sought that are consistent with the global observations of the hot cluster gas. Once this problem is solved, the next step is to determine if real black holes can indeed supply energy in the form and amount required.  We consider here jet-dominated flows in which gas near the cores of cluster-centered elliptical galaxies is accelerated in bipolar outflows when triggered by a supermassive black hole feedback mechanism. Both mass and energy are transported out from the center. The feedback appears largely as an outward flow of mass rather than energy, although the energetics of successful flows must be understood {\\it a postiori}. We find that these more massive jets have many desirable long term attributes. We have discussed elsewhere how heated, buoyant gas can rise upstream in cooling flows, transporting mass and energy to distant cluster gas in a manner that sharply reduces the overall cooling rate and preserves the thermal profiles observed (Mathews, et al. 2003; Mathews, Brighenti \\& Boute 2004). We now turn our attention to the possibility that mass-loaded subrelativistic jets can accomplish the same desirable results.  Another inspiration for the calculations described here are recent HI (Morganti et al. 2004; 2005), UV (e.g. Crenshaw et al. 1999; Kriss 2003) and X-ray (George et al. 1998; Risaliti et al. 2005) observations of AGNs showing blue-shifted absorption or emission lines along the line of sight. Neither the optical depth nor the covering factor of the outflowing gas can be accurately determined from these observations, but the outflowing mass flux can be comparable to the Eddington limit, implying that relatively little gas is being directly accreted by the central black hole. The outflow velocity is typically several 100 km s$^{-1}$ but increases to several 1000 km s$^{-1}$ in a few more luminous objects (Kriss 2004). At least half of all AGNs exhibit outflows so it is plausible that they exist in all objects and with substantial covering factors. Winds from accretion disks are one possible explanation for the outflows (Narayan \\& Yi 1994; Konigl \\& Kartje 1994; Blandford \\& Begelman 1999; Proga 2000; Soker \\& Pizzolato 2005), particularly those with higher velocities.  The highly uncertain outflowing mass flux observed in low luminosity galaxies (${\\dot M} \\sim 1$ $M_{\\odot}$ yr$^{-1}$), is far less than the mass flux that arrives at the central galaxies in group and cluster cooling flows, ${\\dot M} \\sim 10-300$ $M_{\\odot}$ yr$^{-1}$. However, outflowing gas in these more massive flows is likely to be too hot, too rarefied and too highly ionized to produce blueshifted UV or X-ray lines and would be difficult to observe. Indeed, currently available UV and X-ray absorption observations may naturally select the densest, coldest and most slowly moving gas in each outflow. Many authors (e.g. Churazov et al. 2002; 2005; Peterson \\& Fabian 2005) argue that the mechanical outflow luminosity $L_{mech}$ from massive black holes can greatly exceed their bolometric luminosity, and this may apply to all radiatively inefficient black holes (Hopkins et al. 2005).  As in the Blandford-Begelman accretion model, we assume that the majority of gas accreted at the outer edges of the accretion disks ultimately flows away from the disk surface as a fast but nonrelativistic wind. Such disk flows are unlike the relativistic e$^{\\pm}$ radio-emitting jets driven from much smaller regions near the central hole (Blandford \\& Znajek 1977), but the two types of outflow may coexist. In contrast, winds driven by radiation pressure from thin disks, such as those described by Proga (2000) are confined by strong radial, polodial fields (e.g. Blandford \\& Payne 1982), and may be directed mostly along the equatorial plane. However, more axis-oriented outflows may be possible from thicker disks or with different, more vertical field geometries (e.g. Everett, K\\\"onigl \\& Kartje 2001). Nevertheless, fast wind-driven bipolar flows, such as we consider, may at their origin fill cones of significant solid angles. For example, Proga (2003) finds that most of the mass flux in disk winds at high velocities $v \\gta 1000$ km/s, which we require in our flows, lies within 20 or 30 degrees from the polar axes. Such broad outflows are expected to entrain additional ambient gas as they proceed out from the central region. Therefore, for our exploratory calculations we adopt a bipolar wind geometry similar to the prevailing geometry in models of jet-heated cooling flows cited above, but allow the jets to have larger angular sizes at their source. Most of the mass outflow in our jets arises not from the origin but by entrainment of ambient gas at larger radii -- such a model is supported by the recent observations of Sun, Jerius \\& Jones (2005) discussed below. Omma et al. (2004) considered the initial transient flow resulting from a bipolar outflow of this sort, but did not address the important question whether or not such jets can shut down the cooling flow for many Gyrs while preserving the observed thermal gradient in the hot gas. This is our objective here.  In the 2D gasdynamical models described here nonrelativistic outflows are generated by assigning a fixed velocity to gas that flows into a biconical source region (radius of a few kpc with half angle $\\theta_j \\sim 5 - 20^{\\circ}$) at the cluster center. The acceleration of gas in the source bicone is activated by a feedback recipe triggered as gas flows into the innermost zones. The 2D biconical outflows, which may for example represent a disk wind, proceed along the axis of symmetry of the computational grid. Even when the initial outflow has rather substantial opening angles (i.e. $\\theta_j \\sim 20^{\\circ}$), the flow rapidly concentrates within $\\sim 30$ kpc into a much narrower jet. This compression occurs because the rapid pressure drop in the jet due to expansion causes the jet to be compressed and narrowed by the ambient gas pressure which decreases less rapidly with radius in the cluster gas. As the jet proceeds, it entrains additional ambient gas and its mass flux increases. These solutions have two excellent attributes: after many Gyrs the time-averaged gas temperature profile resembles those observed, $dT/dr > 0$ in $r \\lta 0.1 r_{vir}$, and very little gas cools below $\\sim T_{vir}$. The jet itself is difficult to observe. Finally, because of the intermittent nature of the feedback jet excitation, multiple generations of large X-ray cavities are created with 2 or 4 visible at any time, very similar to Perseus (e.g. Fabian et al. 2005). ",
        "conclusions": "Subrelativistic jet flows that entrain ambient gas may be the essential key for solving the cluster cooling flow problem: Why the hot gas temperature and density profiles resemble cooling flows but show no spectral evidence for cooling below $\\sim T_{vir}/3$ at rates expected from the luminosity $L_x$. Many scenarios have been previously considered in which the gas is heated by a variety of mechanisms including jets. In most previous jet calculations it has been assumed that the jets are primarily sources of energy that reheat the cluster gas with no significant outward mass transport. Overall, these simulations have not been successful in reducing the cooling rate while maintaining the observed temperature and density profiles for many Gyrs. Because of the large number of cooling core clusters observed, cooling must be sharply reduced or arrested for many Gyrs.  The mass-carrying jets considered here are a variant of the circulation flows we have discussed in which buoyant bubbles provide an outward mass transport while most of the X-ray emission comes from a normal cooling interbubble inflow (Mathews, et al. 2003; Mathews, Brighenti \\& Boute 2004).  Our computational results for mass-carrying jets are robust in that we find satisfactory multi-Gyr solutions for a significant range of parameters describing the initial central jet outflow. The important global features of our flows are also insensitive to computational resolution. The jet outflow is stimulated by cooling or inflowing gas near the central supermassive black hole, but the success of our longterm solutions is not strongly dependent on the specifics of this feedback mechanism. The natural intermittancy of the feedback generates multiple generations of X-ray cavities similar to those observed in the Perseus Cluster and elsewhere. Nevertheless, the physics at the source of outflow is poorly understood and observational support is limited. More detailed models of mass-carrying jets will be necessary before they can be fully accepted.  One possible objection to the jet driven mass circulation described here is that SNIa-enriched gas that enters the source cone (within the central E galaxy) is transported out only along the jet axis, unlike the observed iron abundance pattern which is spherically symmetric around the central galaxy (de Grandi et al. 2004). Either the radio jet precesses (Gower, et al. 1982) as in the cluster observed by Gitti et al. (2005) or the jet axis direction was altered by black hole mergers at early times when most of the SNIa iron was produced.  \\vskip.1in  Studies of the evolution of hot gas in elliptical galaxies at UC Santa Cruz are supported by NASA grants NAG 5-8409 \\& ATP02-0122-0079 and NSF grant AST-0098351 for which we are very grateful."
    },
    "0601/astro-ph0601249_arXiv.txt": {
        "abstract": " ",
        "introduction": "Despite compelling indirect astrophysical and cosmological evidence, the fundamental nature of non-baryonic dark matter remains elusive (see Refs.~\\cite{dmreviews1,dmreviews2} for recent reviews). Upgrades of direct detectors looking for scattering of dark-matter particles from nuclei, future neutrino telescopes, and space-based antimatter and gamma-ray detectors will dramatically enhance the possibility to discover dark-matter particles in the Galactic halo. Given our complete ignorance on the nature of dark matter, any anomalous experimental result that might point to dark matter should be carefully analyzed, and possibly cross-correlated both with other observational constraints and with existing theoretical models.  Among indirect-detection techniques that seek products of dark-matter annihilations, those searching for neutrinos and photons play a special role. Unlike electrically charged particles, like antiprotons or positrons, which diffuse in the Galactic magnetic field, the arrival direction of neutrinos and photons points to where the dark-matter annihilation took place. In particular, the forthcoming launch of the Gamma Ray Large Area Space Telescope (GLAST) \\cite{glast} and the rapidly developing field of ground based Atmospheric Cherenkov Telescopes (ACTs) \\cite{ACTreview} make the search for energetic photons from dark-matter annihilation a particularly promising and exciting endeavor.  Since the rate per unit volume for dark-matter annihilation depends on the square of the dark-matter density, there may be great advantage to seeking astrophysical locations where the dark-matter density is believed to be high.  The center of the Milky Way has been viewed as a promising target as it might host a dark-matter spike and/or because the Galactic halo could feature a steep cusp towards its center \\cite{pierobuck}. On the other hand, diffuse gamma-ray backgrounds produced by the spallation of cosmic rays on interstellar gas, and the uncertainties on the distribution of molecular hydrogen in the galactic ridge, may blur the dark-matter-induced signal (but for a different viewpoint see Ref.~\\cite{deboer}).  The presence of other gamma-ray emitters in the central region of our Galaxy, including those associated with the supermassive central black hole and the supernova remnant Sgr A$^*$, however, also make it difficult to distinguish a putative dark-matter signal and its astrophysical background \\cite{galcenter,profumo}.  It has therefore been suggested that one should instead seek dark-matter-induced gamma rays from other sources that feature large dark-matter densities but that remain devoid of background sources.  In particular, several extragalactic sources have been considered \\cite{tyler,baltz,pieri,extragal,evans}.  Among these, dwarf spheroidal (dSph) galaxies are among the most dark-matter dominated systems, featuring mass-to-light ratios as high as $M/L \\sim250\\, M_\\odot/L_\\odot$ \\cite{dsphrev}.  Moreover, at least four dSph galaxies with very large $M/L$ lie within 100 kpc from the Milky Way center.  The dark-matter-induced photon flux depends on the dark-matter distribution, and particularly on its inner structure, which is unfortunately poorly known in the case of local dSphs. Ref.~\\cite{baltz} modeled the dark-matter halos of the four nearby dSph, Carina, Draco, Ursa Minor, and Sextans, with King profiles. Even with optimistic assumptions on the background-rejection capabilities of ACTs, they found that the gamma-ray flux expected from supersymmetric models from local dSphs is typically at least three orders of magnitude below the background. They find that if the dark-matter distribution is clumpy, the signal can be boosted by at most a factor of 40, which would still be insufficient. A re-examination of the Draco dSph galaxy in Ref. \\cite{tyler} arrived at somewhat different conclusions. This work showed that, assuming a steeply cusped isothermal dark-matter halo, a large portion of the supersymmetric parameter space could produce a signal visible at both forthcoming ACTs and at GLAST.  This work also pointed out that, depending on the magnetic-field strength inside Draco, the radio-continuum limits for Draco obtained in Ref.~\\cite{dracoradio} could also rule out sizeable portions of the same parameter space. Ref.~\\cite{pieri} conducted an extensive analysis of the 44 nearest galaxies in the Local Group, pointing out other promising extragalactic candidates, such as M31. A systematic comparison between the case of the Galactic center and that of various nearby dSph galaxies (Sagittarius, Draco, and Canis Major) was carried out in Ref.~\\cite{evans}, where a wide array of halo profiles was also employed. Depending on the latter, Ref.~\\cite{evans} concludes that if dark matter is supersymmetric, dSph galaxies may indeed produce a detectable signal in ACTs and at GLAST.  Recently, a gamma-ray excess from the direction of Draco has been detected by the Solar 2 Heliostat Array CACTUS, located in Barstow, California \\cite{cactus}. Although the robustness of the signal has still to be tested, the possibility of ascribing the excess over the off-source background to dark-matter annihilation in Draco's halo is certainly intriguing, as Draco is a dark-matter dominated system (see Sec.~\\ref{sec:dracodm}) that is not expected to host any other significant gamma-ray source \\cite{tyler,nogammatyler}. This possibility was recently envisaged in Ref.~\\cite{hooper}. In the present analysis, we study the impact on the gamma-ray flux of various, astrophysically motivated, halo models \\cite{mash}. We point out that if the angular resolution of the experiment can be improved and understood, then measurement of the angular distribution of the gamma-ray excess can discriminate between cuspy (e.g., the Navarro-Frenk-White profile \\cite{nfw}) and cored profiles and determine the halo scale radius (Sec.~\\ref{sec:dracoang}). In the absence of a full analysis of the raw counts reported by CACTUS \\cite{cactus}, the estimate of the flux of photons to be attributed to dark-matter annihilation appears particularly critical. Here, we conservatively estimate the putative gamma-ray flux detected by CACTUS \\cite{cactus} bracketing it between an upper estimate where all the CACTUS excess counts over the background are attributed to the dark-matter signal (following the approach of Ref.~\\cite{hooper}), and a lower estimate where only the counts in the innermost angular region are supposed to originate from dark-matter annihilations. We then scrutinize the dark-matter interpretation of the CACTUS excess from Draco on model-independent grounds on the planes defined by the particle mass and annihilation cross section (Sec.~\\ref{sec:dracopdm}). We specialize to the particular case of supersymmetric dark matter in Sec.~\\ref{sec:susy}, where we also make predictions for the expected detection rates for models consistent with the Draco excess with other dark-matter searches. We consider in Sec.~\\ref{sec:line} the possibility that the bulk of the excess photons detected by CACTUS come from the monochromatic $\\gamma\\gamma$ line.  In Sec.~\\ref{sec:decay}, we show that the signal cannot come from the decay of a very long-lived dark species. We finally draw our conclusions in Sec.~\\ref{sec:conclusions}. ",
        "conclusions": "\\label{sec:conclusions}  We considered in this paper the possibility that the gamma-ray excess observed by CACTUS from the direction of the Draco dSph galaxy originates from WIMP annihilation. We summarize below the main results of the present analysis: \\begin{itemize}  \\item We showed that future measurements, with $\\sim0.1$ degree angular resolution, should allow us to distinguish between cored and cuspy halos, even though the current CACTUS angular resolution is not good enough. \\item We estimated the putative gamma-ray flux from dark-matter annihilation considering the two extreme possibilities that the dark-matter signal consists of (1) all the excess counts reported by CACTUS, and (2) the excess over background corresponding to Draco's innermost region. \\item We analyzed, in a model-independent approach, the regions of the particle-mass versus annihilation-cross-section parameter space compatible with the estimated CACTUS excess, imposing the constraints from the null results of gamma-ray searches reported by EGRET and by the VERITAS collaboration observation with the Whipple-10m ACT, and from the absence of a statistically relevant excess in the CACTUS data for photon energies above 150 GeV.  The annihilation cross section range lies well above that expected with standard thermal relic dark matter production. \\item The total excess counts over background reported by CACTUS can only be reproduced with a very cored dark-matter halo and a hard photon spectrum.  If this is the correct explanation for the excess, then higher-angular-resolution measurements should still show a $\\sim1$ degree spread around Draco's center.  Almost all CACTUS-compatible models will be within reach of GLAST. \\item In the case of supersymmetric dark matter, the annihilation cross sections needed to reproduce the CACTUS signal are close to the maximal theoretically allowed values and would imply a negligible thermal relic abundance. CACTUS-compatible supersymmetric models give typically very large detection rates at direct-detection experiments, and a sizeable neutrino flux from neutralino annihilation in the Sun, which in principle could allow a cross-check of the gamma-ray signal that GLAST should see within this scenario. \\item The possibility of a dominantly monochromatic origin of the CACTUS excess is not viable within supersymmetry, and requires, for consistency with the EGRET null result from Draco, dark-matter models that annihilate predominantly to two photons and thus produce a very suppressed continuum photon spectrum. \\item A decaying-dark-matter scenario is ruled out by the resulting inferred gamma-ray flux from the center of the Milky Way and in the diffuse gamma-ray background. \\end{itemize}   \\vspace*{1cm} \\noindent{ {\\bf Acknowledgments} } \\\\ \\noindent We would like to thank Max Chertok, Francesc Ferrer, Jeter Hall, Reshmi Mukherjee, Mani Tripathi, and Piero Ullio for helpful discussions.  This work was supported in part by DoE DE-FG03-92-ER40701 and FG02-05ER41361 and NASA NNG05GF69G."
    },
    "0601/astro-ph0601280_arXiv.txt": {
        "abstract": "We present a study of a faint fuzzy star cluster system in the nearby SB0 galaxy NGC 5195 interacting with the famous spiral galaxy NGC 5194 (M51), based on $HST$ ACS $BVI$ mosaic images taken by the Hubble Heritage Team. We have found about 50 faint fuzzy star clusters around NGC 5195 which are larger than typical globular clusters with effective radii $r_{eff}> 7$ pc and red with $(V-I)>1.0$. They are mostly fainter than $M_V\\approx -8.3$ mag. From the comparison of $BVI$ photometry of these clusters with the simple stellar population models, we find that they are as massive as $\\approx 10^5 M_{\\odot}$ and older than 1 Gyr. Strikingly, most of these clusters are found to be scattered in an elongated region almost perpendicular to the northern spiral arm of NGC 5194, and the center of the region is slightly north of the NGC 5195 center, while normal compact red clusters of NGC 5195 are located around the bright optical body of the host galaxy. This is in contrast against the cases of NGC 1023 and NGC 3384 where spatial distribution of faint fuzzy clusters shows a ring structure around the host galaxy. We suggest that at least some faint fuzzy clusters are experiencing tidal interactions with the companion galaxy NGC 5194 and must be associated with the tidal debris in the western halo of NGC 5195. ",
        "introduction": "Faint fuzzy star clusters are peculiar star cluster populations which are large ($r_{eff}> 7$ pc), faint ($M_V>-8.0$), and red ($(V-I)>1.0$). They are much larger than typical Galactic globular clusters, but are as red as old globular clusters. These star clusters were introduced as a new family of star clusters by \\citet{larb00} in their study of the nearby lenticular galaxy NGC 1023. Although extensive searches over four early-type galaxies (three S0's and one E) were made by \\citet{lar01}, these clusters were found only in two SB0 galaxies NGC 1023 (SB0) and NGC 3384 (SB0), and were not found in NGC 3115 (S0) and NGC 3379 (E1).  The main characteristics of these clusters derived from the imaging and spectroscopic studies \\citep{larb00, larb02, bro02} are summarized as follows: (1) They are larger than normal old globular clusters. Their effective radii range from 7 to 15 pc, while the effective radii of typical globular clusters in our Galaxy are 2--3 pc; (2) They are faint. Most of them are fainter than $M_V>-8$ mag; (3) They are associated with disks of their host galaxies, showing a systematic rotation similar to the rotation of the disk; (4) Those in NGC 1023 are moderately metal-rich with [Fe/H]$=-0.58\\pm0.24$ dex, and are $\\alpha$ element-enriched with [$\\alpha$/Fe]$=0.3\\sim 0.6$ dex; and (5) those in NGC 1023 are older than 7--8 Gyr; (6) their origin is regarded to be associated with the interaction of host galaxies with their neighbors.  During this study, \\citet{peng05} reported that 12 early-type galaxies in Virgo cluster harbor a significant number of diffuse star clusters, and that nine of them are S0s. These clusters have low luminosities ($M_V>-8$), a broad spectrum of size ($3<r_h<30$ pc), and a broad range of color ($0.3<(g-z)_0<1.9$), and they are predominantly red with $(g-z)_0>1.0$. These clusters are mainly characterized by the $g$-band low surface brightness $\\mu_g>20$ mag arcsec$^{-2}$. The red population of these diffuse star clusters shares common features with faint fuzzy clusters, and faint fuzzy clusters are considered to be a subset of the diffuse star clusters.   Recently, \\cite{lee05} reported the discovery of a dozen of faint fuzzy clusters in NGC 5195, which made this galaxy the third known SB0 galaxy inhabited by this rare family of star clusters. However, due to the limited spatial coverage and filters ($VI$) of the $HST$ WFPC2 data they used, details of these clusters remain to be known: How many such clusters exist in NGC 5195? How wide are they distributed? How old are they?  In the course of our study of star clusters in the M51 system, We have used $HST$ ACS mosaic images of NGC 5195 to get the answers to the above questions. In this letter, we report a spatial distribution of these faint fuzzy star clusters in NGC 5195. ",
        "conclusions": "\\citet{bro02} noted that faint fuzzy clusters were found in only two late-type SB0 galaxies, NGC 1023 and NGC 3384, among four early type galaxies investigated. Various observations that found several features of interaction in NGC 1023 \\citep{tul80, cap86} and NGC 3384 \\citep{bus96, sil03} suggest that the formation of faint fuzzy clusters in these galaxies are maybe related to the interaction of their host galaxies with neighbors \\citep{bro02}. This is supported by theoretical studies \\citep{fel02, fel05, bur05}. \\citet{peng05} noted that only two of the galaxies hosting diffuse star clusters have nearby neighbor galaxies. We found from an inspection of the POSS maps and the data in NED\\footnote{The NASA/IPAC Extragalactic Database, http://nedwww.ipac.caltech.edu.} of their sample that 40\\% of 10 bright galaxies with $M_B<-18$ show hints of interaction (e.g., existence of neighbor galaxies, asymmetric outer halo) and 60\\% of them are SB0 or SAB0. Diffuse star clusters in bright galaxies ($M_B<-18$) are found either in SB0s or in interacting galaxies. The only exception is VCC 1720 (SA), which happens to be the faintest among the bright sample ($M_B=-18.9$). Therefore, major interaction of the host galaxy with its companions or the existence of a bar appears to be a key to the understanding of the origin of faint fuzzy clusters and diffuse star clusters in these galaxies.  NGC 5195 is also a peculiar SB0 galaxy interacting with another barred spiral galaxy NGC 5194 which is two or three times more massive than NGC 5195 \\citep{sch77}. However, the estimated ages of the faint fuzzy clusters suggest that it is not likely that these clusters were formed during the encounter with NGC 5194. This is in contradiction with the theoretical merging models proposed for the formation of faint fuzzy clusters in NGC 1023 \\citep{fel02}. This result indicates that the faint fuzzy clusters might have formed during the interaction of galaxies much earlier than the age range investigated in the previous dynamical models or they might have formed via some other mechanisms.  To date, NGC 1023 is the only galaxy which has been confirmed to have a rotating disk system composed of faint fuzzy clusters \\cite[]{bro02, bur05}. \\citet{peng05} found that diffuse star cluster populations in nine S0 galaxies in Virgo cluster are both aligned with the host galaxy light and associated with galactic disks. Nonetheless, the faint fuzzy clusters in NGC 5195 do not show any definite evidence in their spatial distribution supporting their association with the optical disk of NGC 5195. The elongated distribution of the faint fuzzy clusters looks different from the bright optical body of NGC 5195 that is rather circular with ellipticity $e=0.21$ and inclined about $37^{\\rm o}$ (north side to our direction) \\cite[]{dev91}. However, faint stellar light is seen in the region where the faint fuzzy clusters are found. In addition, the center of the elongated distribution of the faint fuzzy clusters is off about $1'$ northwest from the center of NGC 5195. Hence, it is not evident whether the faint fuzzy clusters belong to the disk of NGC 5195.  What if the elongated distribution should be the revelation of any physical structure unknown? From HI observation, NGC 5195 was found to be nearly deficient in HI gas except for the gas associated with the spiral arm of NGC 5194 touching its east side \\cite[]{rot90}. Only in the very center of NGC 5195, a high content of gases is detected from CO and HCN observation by \\cite{koh02}. Therefore we conclude that the elongated structure should be very old and devoid of gas.  It is also possible that the faint fuzzy clusters showing the elongated distribution as a whole may belong to the debris of tidal interactions between NGC 5194 and NGC 5195. \\citet{cha04} found 20 large globular clusters with $r_{eff}>7$ pc in the four $HST$ WFPC2 fields of NGC 5194, much more than in other spiral galaxies in their sample. Nine of these clusters are similar to faint fuzzy clusters found in NGC 5195. This indicates that faint fuzzy clusters might be dispersed not only over NGC 5195 but also over NGC 5194 through the interactions. It is noted, however, that about a dozen faint fuzzy clusters were found in one $HST$ WFPC2 field of NGC 5195 by \\citet{lee05}, while only nine clusters were found in the four HST WFPC2 fields of NGC 5194.   Recently, \\citet{bur05} suggested a scenario that the passage of a small companion galaxy through or near the center of a disk galaxy could result in the annular distribution of faint fuzzy clusters in NGC 1023. In M51 system, the mass of NGC 5195 is 1/2 to 1/3 of that of NGC 5194 and the faint fuzzy clusters are not in a well-defined annular but in an elongated distribution, which is quite a different case from NGC 1023. Therefore \\citet{bur05}'s scenario may not be applied to the case of NGC 5195.  In summary, we present a study of faint fuzzy clusters in NGC 5195 based on the analysis of the HST ACS $BVI$ mosaic images. We found about 50 such clusters in the $6.8' \\times 4.0'$ field around NGC 5195. Most of the faint fuzzy clusters appear to be distributed over an elongated region whose center is about one arcmin off from the center of NGC 5195. while compact red clusters are located around NGC 5195. Some faint fuzzy clusters are at least regarded as parts of tidal debris produced by the interactions between NGC 5194 and NGC 5195. Spectroscopic studies are needed to investigate the age, metallicity, and kinematics of these faint fuzzy clusters and to understand their origin."
    },
    "0601/astro-ph0601024_arXiv.txt": {
        "abstract": "{During the past years photometric surveys, later complemented by follow-up radial-velocity measurements, have revealed the presence of several new extra-solar transiting planets, in very short period orbits. Many of the host stars are extremely faint (V$\\sim$16), making high-precision spectroscopic measurements challenging.} {We have used the UVES spectrograph (VLT-UT2 telescope) to obtain high resolution spectra of 5 stars hosting transiting planets, namely for OGLE-TR-10, 56, 111, 113 and TrES-1. The immediate objective is to derive accurate stellar parameters and chemical abundances.} {The stellar parameters were derived from an LTE analysis of a set of \\ion{Fe}{i} and \\ion{Fe}{ii} lines. } {Complementing the spectroscopic information with photometric transit curves and radial-velocity data from the literature, we have then refined the stellar and planetary radii and masses. The obtained data were also used to study and discuss the relation between the stellar metallicity and orbital period of the planets.} {} ",
        "introduction": "Although the majority of known extrasolar planets was discovered using the radial-velocity technique\\footnote{For a continuously updated list see table at http://obswww.unige.ch/Exoplanets}, some major photometric transit searches are now delivering important results, giving a new breath to the study of exoplanets. Many candidates have been announced by surveys like OGLE or TrES \\citep[][]{Udalski-2002,Alonso-2004}, and in a few cases, the planetary nature was confirmed by follow-up radial-velocity measurements \\citep[][]{Konacki-2003,Bouchy-2004,Pont-2004, Alonso-2004,Bouchy-2005a,Konacki-2005}. Together with the few transiting planets found in the context of radial-velocity surveys \\citep[][]{Charbonneau-2000,Henry-2000,Sato-2005,Bouchy-2005b}, these discoveries are now providing information about the physical properties of the giant planets themselves (e.g. radius, mean density), and opening the possibility to confront the observed properties with those predicted by the models.  The new detections have also raised a lot of scientific problems. Part of the discovered transiting planets \\citep[][]{Konacki-2003,Bouchy-2004} have orbital periods of the order of 1-2 earth days. This is in clear contrast with the known pile-up of planets with periods above 3-days observed for (hot-jupiter) giant planets discovered by the radial-velocity technique \\citep[see e.g.][]{Santos-2005b}. Whether this is simply the result of a detection bias, or if this can be explained by some other physical process, is now a matter of debate \\citep[e.g.][]{Gaudi-2005}. On the other hand, among the confirmed transiting planets, HD\\,209458 \\citep[][]{Charbonneau-2000,Henry-2000} has the lowest, and most anomalous mean density \\citep[][]{Baraffe-2005,Guillot-2002}. Curiously also, a relation between the planet mass and the orbital period also seems to exist \\citep[][]{Mazeh-2005}. The understanding of these issues may give us some important clues about the processes of giant planet formation and evolution.  In order to access these problems we have to be able to derive accurate parameters for the planets. But the derivation of the planetary properties (mass, radius, and as a consequence their mean density) may depend considerably on the deduced parameters for the stellar host. For some cases, namely for the planets discovered in the context of the OGLE survey, the host stars are extremely faint (V$\\sim$16), making difficult the task of obtaining precise spectroscopic parameters.  In this paper we present accurate stellar parameters and chemical abundances for 5 stars known to be orbited by transiting planets, namely for OGLE-TR-10, OGLE-TR-56, OGLE-TR-111 and OGLE-TR-113, as well as for the brighter planet-host TrES-1. We use the obtained values to derive revised stellar and planetary radii and masses, as well as to study the relation between the presence of planets and the stellar metallicity \\citep[e.g.][]{Gonzalez-1998,Santos-2001,Santos-2004b} for short orbital period planets. In a separate paper we derive and discuss the ages of stars with short period transiting planets \\citep[][]{Melo-2006}. ",
        "conclusions": "We have used high resolution and S/N spectra taken with the UVES spectrograph (VLT-UT2) to derive accurate stellar parameters and metallicities for 5 stars known to be transited by giant planets. The resulting parameters were used to study both the relation between the stellar metallicity and the existence of giant planets in short period orbits, and to derive revised masses and radii for the transiting planets.  The results show that the stars with transiting planets studied in the current paper follow the general trend of high metallicities found for the majority of the planet hosts \\citep[][]{Gonzalez-1998,Santos-2001}. A weak, but not significant correlation is found for the average metallicity as a function of the orbital period among the short period planets. The addition of more data is needed to confirm/deny this result."
    },
    "0601/astro-ph0601354_arXiv.txt": {
        "abstract": " ",
        "introduction": "The so-called data explosion in astronomy promises exciting new scientific developments, but brings with it many technical challenges, in collecting, storing, transporting, processing and visualising data. Virtual Observatories (VO) have developed to meet some of these technical challenges. Falling under the broad area of e-science (which incorporates other scientific domains facing similar challenges, such as genetics and particle physics) the aim of Virtual Observatory research is to provide the tools necessary for dealing with this data.  The Australian Virtual Observatory (Aus-VO)\\footnote{{\\tt http://www.aus-vo.org}} was started in 2003 with the aim of both contributing to the international VO effort, and developing tools of use to Australian astronomers. Australia has many areas of expertise (for example radio astronomy) and it makes sense to focus our efforts on providing modern tools for working in these areas. In this context the ATNF decided to develop a range of tools for storing, processing and visualising data from the Australia Telescope Compact Array. The aim was that these tools would be useful to astronomers now, and at the same time let us explore the technology that would be necessary for developing software for future telescopes such as the Square Kilometre Array (SKA).  This paper is based on a talk given at the ASA Annual Meeting in Sydney in July 2005. After giving some background about the ATCA and the Virtual Observatory, we discuss three tools developed at the ATNF over the last few years. Firstly the Australia Telescope Online Archive which contains all of the data collected so far by the ATCA; secondly a prototype data reduction pipeline for ATCA data; and finally the Remote Visualisation System for viewing large datasets. ",
        "conclusions": "The ATOA has been public since December 2004. We hope that it will encourage the reuse of ATCA data for projects other than those it was originally intended for. The framework used for the ATOA could easily be extended to include data from other telescopes and can be updated as additional metadata is required.  The most significant improvement of the ATOA over existing online archives (such as NRAO\\footnote{{\\tt http://archive.nrao.edu/archive/e2earchive.jsp}} and MAST\\footnote{{\\tt http://archive.stsci.edu/}}) is the data delivery mechanism. Most existing archives do not support on demand delivery of data over the web, instead requiring the user to submit a form requesting files that then have to be transfered to a publicly accessible ftp site or to other media (such as CD) for physical delivery. In the ATOA, the batch downloading of multiple files is handled by a streaming TAR or ZIP archiving algorithm that performs dynamic archiving as files are streamed over the web, requiring no additional disk space on the server for these operations.  In developing the ATCA data model and considering the type of metadata required for automatic processing we identified several new metadata types that would be useful to store in the RPFITS files. As a result the following fields have been added to the RPFITS files and will be available in all future ATCA data: \\begin{itemize} \\item four calibrator codes \\begin{description} \\item[C] (standard phase calibrator) \\item[F] (primary flux calibrator) \\item[B] (bandpass calibrator) \\item[P] (pointing calibrator) \\end{description} \\item Pointing offsets \\item Weather data: added rain gauge and phase rms and difference \\item Attenuator settings at start of scan \\item Subreflector position \\item Correlator configuration \\item Scan type \\item Coordinate type \\item Line mode \\item CACAL counter \\end{itemize} These will help both automatic processing systems and astronomers assess the data quality in the observations they are interested in. A full e-logbook system will be used in the future as currently the logs are all stored on paper at the telescope and hence are not easily accessible to ATOA users.  We have developed a prototype pipeline for processing of raw data for single-pointing continuum images. This is attached to the ATOA to provide an improved service for users of the ATOA. At this stage the image quality is suitable for previewing the data in archive to see if it is of interest. Further manual processing would then be required to obtain images of scientific quality.  A significant challenge in developing the ATOA and the prototype pipeline were integrating pre-existing software with modern software tools. For example, the Glish scripting language has no web service libraries and so an extra layer had to be developed between the data processing level and the web services. If re-implementing from scratch, a language such as Python would be a better alternative for developing the pipeline.  In developing these tools we have started to explore the techniques necessary for astronomical software development in the VO era. This is essential for future telescopes and surveys that Australia will produce. Making access to existing Australian data as easy as possible will maximise its use in the international community."
    },
    "0601/astro-ph0601162_arXiv.txt": {
        "abstract": "We present results from the analysis of the optical spectra of 47 radio-selected narrow-line Seyfert 1 galaxies (NLS1s). These objects are a subset of the First Bright Quasar Survey (FBQS) and were initially detected at 20 cm (flux density limit $\\sim1 $ \\rm mJy) in the VLA FIRST Survey. We run Spearman rank correlation tests on several sets of parameters and conclude that, except for their radio properties, radio-selected NLS1 galaxies do not exhibit significant differences from traditional NLS1 galaxies. Our results are also in agreement with previous studies suggesting that NLS1 galaxies have small black hole masses that are accreting very close to the Eddington rate. We have found 16 new radio-loud NLS1 galaxies, which increases the number of known radio-loud NLS1 galaxies by a factor of $\\sim 5$. ",
        "introduction": "Narrow line Seyfert 1 galaxies (hereafter NLS1s) are a class of Active Galactic Nuclei (AGN) which, like Seyfert 1 galaxies, exhibit both permitted and forbidden optical emission lines.  Unlike standard broad-line Seyfert 1 galaxies, NLS1 galaxies additionally exhibit the following observational properties \\citep{goo89}: (1) the FWHM of the H$\\beta$ line is less than 2000 km s$^{-1}$, (2) permitted lines are only slightly broader than the forbidden lines, (3) the forbidden-line emission is relatively weak, i.e., [OIII]$\\lambda5007$ / H$\\beta <$ 3, and (4) FeII and other high ionization emission line complexes are unusually strong, i.e., stronger than ``normal'' Seyferts, especially in the 4500$-$4680 \\AA\\/ and 5105$-$5395 \\AA\\/ wavelength ranges \\citep{ost85}.   It is not well understood at present what mechanism drives these differences between normal broad-lined Seyfert 1 and NLS1 galaxies. Two opposing models exist, each with their own successes and shortcomings. The first model suggests NLS1 galaxies are viewed with the accretion disk in a face-on orientation. This case is similar to the Unified Model explanation for Seyfert 1 galaxies, with the crucial difference being that the emission lines appear narrower than typical Seyfert 1 galaxies because the optical line emitting clouds are primarily confined to a disk perpendicular to the symmetry axis, reducing the velocity dispersion along the line of sight \\citep{ost85,goo89, Puchnarewicz92,bol96}. The appearance of strong FeII emission could also be explained by a pole-on geometry if the FeII emission originates in the accretion disk as has been suggested by \\cite{fab89}.   It has been alternatively proposed that the unique properties exhibited by NLS1 galaxies are a result of having accretion rates much closer to the Eddington limit than ``normal'' broad-line Seyfert galaxies \\citep{lao97}.  Assuming that the efficiency of conversion of accretion energy to radiation is the same in both near- and sub-Eddington sources, those sources accreting near the Eddington rate (NLS1s) should contain smaller mass black holes compared to sub-Eddington sources of the same luminosity. Smaller mass black holes in NLS1 galaxies would lead to narrower optical emission lines if the motion of the BLR region clouds are dominated by Keplerian velocities. This model also provides an explanation for many other peculiar properties of the NLS1 galaxies. For example, when the accretion rate approaches Eddington, the disk is likely ``puffed up'' by radiation pressure. This would result in a larger X-ray-heated volume that can (a) generate the observed FeII emission and (b) shield the NLR from UV ionizing radiation, resulting in weaker [OIII] emission.  Additional fuel which favors the accretion rate model, is provided by the work of \\cite{bor92}. They measured many continuum and line properties for a variety of AGN and examined various correlations and inverse correlations between myriad parameters. They then reduced this large set of data to a few principal components, eigenvectors of the correlation matrix. The most important eigenvector is principal component 1 (PC1) which links the strength of FeII emission, [OIII] emission, and H$\\beta$ line asymmetry, but is dominated by the inverse correlation between the strengths of FeII and [OIII]. The most popular interpretation for what PC1 physically represents is a sequence in L/L$_{\\rm Edd}$, the Eddington ratio. The idea is that the vertical structure of the accretion disk, governed by the Eddington ratio, drives line strengths and continuum components through its illumination of broad-line clouds and an extended narrow-line region. NLS1 galaxies lie at one extreme of the PC1 correlation (strong FeII / weak [OIII]), lending support to the idea that near Eddington accretion rates (around small mass black holes) ultimately determine whether an object manifests itself as a NLS1 or Seyfert 1 galaxy.   But more than optical emission line properties distinguish NLS1 galaxies from other classes of AGN; their X-ray properties also show marked differences compared to other types of Seyfert galaxies. The best known X-ray property of NLS1 galaxies is the presence of a soft X-ray excess \\citep{bra94}). NLS1 galaxies also exhibit unusually strong X-ray variability \\citep{bol97,lei99}.  Finally, NLS1 galaxies are known to generally exhibit steeper hard (2-10 keV) X-ray spectra than broad-line Seyfert 1 galaxies \\citep{BMA97}.   Most of these X-ray properties can be explained within the framework of either the orientation or accretion-rate models, although each requires some special assumptions or modifications.  Within the general confines of the orientation model, \\cite{madau88} has examined the properties of geometrically thick accretion disks and predicts a strong soft X-ray emission when these disks are viewed face-on. In the accretion-rate models, the observed soft excesses in NLS1 galaxies is attributed to thermal emission from a viscously heated accretion disk that results at an accretion rate close to the Eddington limit \\citep{pou95}. Likewise, enhanced variability in NLS1 galaxies can be explained via either the orientation or accretion rate models. If NLS1 galaxies do accrete near Eddington and have smaller mass black holes, their size scale is smaller and the variability can be more rapid. In the orientation model, the variability may be enhanced via relativistic boosting of the emission.  The trend toward harder spectra in normal broad-lined Seyfert galaxies is the one observational characteristic which is difficult to explain in the orientation model. In this model, the hard photon index is expected to become softer as the inclination angle increases \\citep{haardt93}, the opposite of which must be true for this model to explain differences seen between NLS1 and normal Seyfert 1 galaxies. However, when this model is generalized, and the hot phase is assumed to be in localized regions on the disk, the dependence on inclination decreases, and the index depends rather on the height of the hot phase blobs above the disk. In this case, there is a need to invoke a mechanism which links the height of sites which produce the bulk of the observed emission with orientation in order to fit within the orientation based scheme. This is a difficult proposition without invoking some ad hoc physical structure that reflects or absorbs radiation asymmetrically. In the high accretion rate model, however, the steeper hard X-ray spectrum of NLS1 galaxies can by explained if the hot phase luminosity is smaller in NLS1 galaxies, presumably due to the smaller central black hole mass.   Despite years of work, both the orientation and accretion rate models are still essentially equally viable; both provide a framework which unifies NLS1 galaxies and normal Seyfert 1 galaxies, yet neither is entirely satisfactory.  Additional observational constraints would go a long way toward resolving the question. Radio observations may well provide those constraints; they have undoubtedly been crucial in the formulation of unification scenarios for other classes of AGN. But while Seyferts of all types have been studied extensively at optical and X-ray wavebands, little is known about the radio properties of NLS1 galaxies in part because they are typically faint radio sources, if they are known to be radio sources at all. One of the few systematic studies of the radio properties of NLS1 galaxies is by \\cite{ulv95}. They observed a group of 15 NLS1 galaxies (of which only 9 were detected in radio) and found relatively modest radio powers for their sample (10$^{27}$ - 10$^{30}$ erg s$^{-1}$ Hz$^{-1}$). These 15 NLS1 galaxies were all radio quiet, i.e., the ratio of radio luminosity to optical luminosity is relatively small. \\cite{mor03} show in their study that NLS1 galaxies can be radio intermediate and even radio luminous, and, as a whole, are more radio luminous than classical Seyferts.    The work of \\cite{ulv95} is a radio study of optically selected NLS1 galaxies while the \\cite{mor03} work is a radio study of X-ray and IR selected NLS1 galaxies.  There are no studies of a radio-selected sample of NLS1 galaxies. This is partly because only recently has it become apparent that NLS1 galaxies can exhibit significant radio luminosities \\citep{siebert99,gliozzi01,osh01}. Additionally, it is only recently that new radio surveys have emerged which cover a wide area of sky to depths of only a few mJy. It is from one such survey, the FIRST (Faint Images of the Sky at Twenty Centimeters; \\cite{bec95}) that we have compiled the first known sample of radio-selected NLS1 galaxies, and in the process increased by a factor of five the number of known radio-loud NLS1 galaxies.   In this paper we present 47 NLS1 galaxies along with 15 BLS1 galaxies with very narrow lines selected from FIRST and analyze their optical and radio properties. We will use an H$_{0} = 65$ km s$^{-1}$, $\\Omega_{matter} = 0.3$, $\\Omega_{\\Lambda} = 0.7$ cosmology throughout.   The structure of this paper is as follows: In sections 2 and 3 we describe the sample selection criteria. In section 4 and 5 we describe the data and the data analysis and in section 6 the method we use for estimating central black hole mass in explained. We present our results and correlations in Section 7 and discuss the radio properties in section 8. The conclusion follows in section 9. ",
        "conclusions": "In this paper we presented the optical and radio properties of a sample of radio selected Narrow-line Seyfert 1 galaxies. We find 16 new radio loud NLS1 galaxies to add to the three that were previously discovered to increase the number of known radio loud NLS1 galaxies significantly. We have run Spearman rank correlation tests between several sets of parameters and find that, for the most part, our results agree with previous studies of NLS1 galaxies. When compared with X-ray selected and optically selected samples of NLS1 galaxies, the properties of radio-selected NLS1 galaxies do not exhibit many obvious differences from NLS1 galaxy samples selected in other manners. We also use Student's t-test to determine whether radio loud NLS1 galaxies are statistically different from their radio quiet counterparts and conclude that there are few differences.  We estimate black hole masses with a relationship between FWHM(H$\\beta)$, optical luminosity, and M$_{BH}$. We also estimate bulge velocity dispersion from FWHM([OIII]) and find that NLS1 galaxies do not obey the tight correlation between bulge mass and black hole mass, confirming the result of \\cite{gru04}.  Our sample seems to agree with the canonical model of NLS1 galaxies, in that they have small black hole masses and are accreting near the Eddington rate.  We have found a few discrepancies between our sample and previous studies of NLS1 galaxies and other AGN. We have found an anti-correlation between black hole mass and radio loudness showing that as $M_{BH}$ increases, $\\log \\rm R^{*}$ decreases. In previous studies of AGN, it has been found that radio loudness should increase as black hole mass increases. We test several correlations from \\cite{gru99} and \\cite{bor02} and find that, while many of the properties in our sample agree with theirs, there are a few exceptions. We do not find an anti-correlation between the strength of FeII and [OIII], nor do we find any connection between FWHM(H$\\beta$) and EW([OIII]).  The most significant result found in this study is that we may now have to consider that NLS1 galaxies can be strong radio emitters. Although the galaxies in this study do not show the prodigious radio luminosities of radio loud quasars, the luminosities we find are significantly greater than the typical Seyfert galaxy.  Our sample of radio selected NLS1 galaxies adds significantly to the total of published NLS1 galaxies. Further investigation into this growing class of radio loud NLS1 galaxies could help define what is different physically between NLS1 and traditional Seyfert 1 galaxies. The authors are involved in a study of the X-ray properties of the most radio luminous objects in this sample, which we hope will reveal what is physically different about these NLS1 galaxies."
    },
    "0601/astro-ph0601481_arXiv.txt": {
        "abstract": "We show that in hybrid inflation it is possible to generate large second-order perturbations in the cosmic  microwave background due to the instability of the tachyonic  field during preheating. We carefully calculate this effect from the tachyon contribution  to the gauge-invariant curvature perturbation, clarifying some confusion in the literature concerning nonlocal terms in the tachyon curvature perturbation; we show explicitly that such terms are absent.  We quantitatively compute the nongaussianity generated by the tachyon field during the preheating phase and translate the experimental constraints on the  nonlinearity parameter $f_{NL}$ into constraints on the parameters of the model. We also show that nonscale-invariant second-order perturbations from the tachyon field with spectral index $n=4$ can become larger than the inflaton-generated first-order perturbations, leading to stronger constraints than those coming from nongaussianity.  The width of the excluded region in terms of the logarithm of the dimensionless coupling $g$, grows linearly with the log of the ratio of the Planck mass to the tachyon VEV, $\\log(M_p/v)$; hence very large regions are ruled out if the inflationary scale $v$ is small.  We apply these results to string-theoretic brane-antibrane inflation, and find a stringent upper bound on the string coupling, $g_s < 10^{-4.5}$. ",
        "introduction": "Inflation \\cite{Guth} has become the dominant paradigm for the generation of the anisotropies of the Cosmic  Microwave Background (CMB) (see \\cite{RiottoReview,BrandenbergerReview} for a review). During inflation quantum fluctuations are redshifted out of the horizon where they become ``frozen'' and later re-enter the horizon after Big Bang Nucleosynthesis.  Observation of the temperature anisotropies of the CMB directly constrain the curvature perturbation somewhat before cosmological scales re-enter the horizon.  Current observations show that the curvature perturbation is gaussian with a scale-invariant spectrum to within  experimental limits, which is consistent with the predictions of inflation.  Nevertheless, some nongaussianity is expected to be generated during inflation \\cite{inflationaryNG}-\\cite{otherNG}  (see \\cite{NGreview} for a review).  It is interesting to study the level of nongaussianity, usually parameterized by a dimensionless nonlinearity parameter $f_{NL}$, in various scenarios for inflation, since it can be an important way to discriminate between different models.  For example, the curvaton mechanism \\cite{curvaton} for the generation of cosmological perturbations may be  ruled out by future non-observation of nongaussianity \\cite{curvatonNG} (see, however, \\cite{curvatonNGexception}).  One way to generate significant levels of nongaussianity at the end of  inflation is by preheating \\cite{preheatNG0}-\\cite{preheatNG6}. Preheating refers to the  nonperturbative production of particles which arises either due to coherent oscillations \\cite{preheating} or tachyonic (spinodal) instability \\cite{tachyonic1}-\\cite{tachyonic3} (see also \\cite{ReheatTheory,Bassett}). It has also been suggested that preheating might affect the first order variation of the  curvature perturbation, perhaps dominating the contribution of the inflaton and thus modifying the standard predictions of inflation \\cite{firstorderreheating}-\\cite{FB2}.  In this paper we carefully reconsider the generation of nongaussian fluctuations from tachyonic preheating in hybrid inflation \\cite{hybrid} using second order cosmological perturbation theory similarly to \\cite{Acquaviva,EV}.  See \\cite{doubleinflation} for constraints on the parameter space of hybrid inflation coming from the WMAP data.  During hybrid inflation the inflaton field $\\varphi$ is displaced from its preferred value, and inflation is driven by the false vacuum energy of the ``waterfall'' or tachyon field $\\sigma$ which is trapped in its false  vacuum by interactions with the inflaton.  When the inflaton reaches some critical value the tachyon  effective mass squared  becomes negative and its fluctuations grow exponentially due to the spinodal instability.  The exponential growth continues until the fluctuations of the tachyon start to oscillate about the true vaccuum.  At this stage the tachyon fluctuations become nonperturbative and their back-reaction brings inflation to an end.  Here we study the evolution of the second-order curvature perturbation due to the tachyonic instability and find that it can lead to a large level of nongaussianity, if the tachyon was  lighter than the Hubble scale $m_\\sigma < 3H/2$, during a sufficient part of the inflationary epoch. (If $m_\\sigma \\gg H$, the fluctuations of $\\sigma$ are exponentially suppressed during inflation.)  In the hybrid inflation model, it can be natural to have $m_\\sigma < 3H/2$ during inflation, since $m_\\sigma^2$ changes sign at the end of inflation.  If the inflaton is rolling slowly enough, $m_\\sigma^2$ will naturally already be close to zero even before the end of inflation, at the time of horizon crossing.  The scenario is summarized as follows.  During inflation the squared mass of the tachyon field varies linearly with the number of e-foldings of inflation, in the vicinity of the point where it passes through zero: $m^2_\\sigma = c H^2 N$, taking $N=0$ to be the value where $m^2_\\sigma=0$, and $c$ is a function of the parameters of hybrid inflation.    Inflation starts at some $N\\equiv N_i < 0$, and ends at $N\\equiv N_* > 0$, for a total of  $N_e = -N_i + N_*$ e-foldings.  We are going to limit our inquiry to regions of parameter space where the linear behavior of  $m^2_\\sigma$ is valid during the full duration of inflation, since this technical assumption makes the calculations tractable; however we will show that this assumption must often be satisfied anyway, as a consequence of the experimental limit on the inflationary spectral index.  At $N=0$, fluctuations of the tachyon field start to grow exponentially, and can make a large contribution to the curvature perturbation at second order in cosmological perturbation theory, resulting in a scale-noninvariant component both in the CMB  temperature power spectrum, and its nongaussian correlations.   Our goal is to derive constraints on the parameter space of hybrid inflation from these effects.  We start in section \\ref{II} by defining the perturbations up to second order in the metric and inflationary fields.  In section \\ref{III} the hybrid inflation model is reviewed, starting with the dynamics of the fields at zeroth order, and then their first order perturbations, with emphasis on the dynamics of the tachyon field fluctuations toward the end of inflation. In section \\ref{IV} we solve for the second order curvature perturbation which is induced by the tachyonic fluctuations. This is used in section \\ref{V} to compute the bispectrum (three-point function) and the nonlinearity parameter $f_{NL}$, as well as contributions to the spectrum itself. In section \\ref{VI} we incorporate experimental constraints on  nongaussianity as well as on the spectrum, to derive excluded regions in the parameter space of the hybrid inflation model. In section \\ref{VII} we consider the possibility of generating significant nongaussianity in brane inflation. We conclude in section \\ref{VIII}, comparing our results with previous treatments.  We give technical details concerning mode functions in de Sitter space, the second order perturbed Einstein equations, the inflaton curvature perturbation, Fourier transforms of convolutions, the construction of the tachyon curvature perturbation and the approximation of an adiabatically varying tachyon mass in appendices A-F. ",
        "conclusions": "\\label{VIII}  We have carefully calculated the contribution from tachyonic preheating to the spectrum and bispectrum of the curvature perturbation in a model of  hybrid inflation, and derived constraints on the model parameters  using experimental limits on the power spectrum and nongaussianity in the CMB.  The tachyonic contribution to the spectrum, which we parametrized by $f_L$, has spectral index $n=4$, and the experimental limit on $f_L$ turns out to always be a stronger constraint than that coming from nongaussianity.  Thus we do not expect to see nongaussian signals in future data sets, if hybrid inflation is the underlying model.  Rather, one should look for an excess of power at small scales to see the effects we have investigated.  We have studied a simplified model of hybrid inflation, in which the tachyon is a real scalar field.  This model is not viable as it stands, because domain walls will be copiously produced at the end of inflation, and they will quickly dominate the universe.  We need to assume that there are some small operators in the Lagrangian which violate the symmetry $\\sigma\\to -\\sigma$, and make the domain walls unstable. Although our results are strictly valid only for hybrid inflation with real fields, we have made a preliminary attempt to apply them to the warped brane-antibrane inflation model, by matching the coefficients of the lowest dimension operators in the corresponding Lagrangians.  To the extent that this is valid, we find an interesting limit on the string coupling, $g_s < 10^{-4.5}$.  One achievement of this work is our demonstration that nonlocal contributions to the second order curvature perturbation, as defined herein, vanish when care is taken to consistently include all terms of a given order in the slow-roll expansion.  The nonlocal terms which appear at intermediate steps when using the generalized longitudinal gauge have been the bane of some previous attempts to do similar calculations.  They should not appear in any physical quantities since they violate causality.  Any estimates made using the nonlocal terms are likely to grossly overestimate perturbations on large scales, since these terms are singular at small wave numbers.  Some readers may feel uneasy about an event at the end of inflation being able to affect fluctuations whose wavelength corresponds to the horizon at the beginning of inflation.  However, as has been emphasized in \\cite{FB1,FB2}, there is no violation of causality here.  Consider a collection of comoving observers who are within the horizon at the beginning of inflation, but are about to lose causal contact with each other.  If they agree to set up perturbations with small wavelengths in their Hubble patches at the end of inflation, this will induce large-scale correlations between different Hubble patches.  Since the tachyon field is synchronized across these different Hubble patches it is able to amplify the large-scale perturbations in a coherent way.   One avenue for future study is the more realistic class of models where $\\sigma$ is complex and the defects which form are cosmic strings instead of domain walls. However, this is a much more technically difficult model to study than the case of the real scalar.  In the latter case, the VEV of $\\sigma$ remains zero even during the tachyonic preheating phase, because the field rolls in opposite directions on either side of each domain wall, and the VEV averages to zero.  In the complex case, the modulus of $\\sigma$ grows uniformly during preheating, and the randomly varying phase is only a gauge degree of freedom.  We can no longer decouple the equations for the fluctuations of the metric and the inflationary fields when $\\sigma$ has a VEV.  It thus remains an interesting question whether the realistic hybrid inflation model also gets significant nongaussianity from  tachyonic preheating \\cite{inprogress}.   \\bigskip{\\bf Acknowledgments.}  We would like to thank A.\\ Mazumdar and A. Jokinen for stimulating discussions which motivated us to start on this work.  We also thank R.\\ Brandenberger, C.\\ Burgess, D.\\ Lyth, K.\\ Malik,  A.\\ Notari and D.\\ Seery for helpful comments. We thank the referee for thoughtful and insightful comments which helped us to strengthen our arguments.  This work was supported by NSERC of  Canada and FQRNT of Qu\\'ebec.    \\renewcommand{\\theequation}{A-\\arabic{equation}} \\setcounter{equation}{0}"
    },
    "0601/astro-ph0601029.txt": {
        "abstract": "We have studied $\\sim$ 2100 early-type galaxies in the SDSS DR3 which have been detected by the GALEX Medium Imaging Survey (MIS), in the redshift range $0<z<0.11$. Combining GALEX $UV$ photometry with corollary optical data from the SDSS, we find that, at a 95 percent confidence level, \\emph{at least} $\\sim$ 30 percent of galaxies in this sample have $UV$ to optical colours consistent with \\emph{some} recent star formation within the last Gyr. In particular, galaxies with a $NUV-r$ colour less than 5.5 are \\emph{very} likely to have experienced such recent star formation, taking into account the possibility of a contribution to $NUV$ flux from the UV upturn phenomenon. We find quantitative agreement between the observations and the predictions of a semi-analytical $\\Lambda$CDM hierarchical merger model and deduce that early-type galaxies in the redshift range $0<z<0.11$ have $\\sim $ 1 to 3 percent of their stellar mass in stars less than 1 Gyr old. The average age of this recently formed population is $\\sim$ 300 to 500 Myrs. We also find that `monolithically' evolving galaxies, where recent star formation can be driven \\emph{solely} by recycled gas from stellar mass loss, \\emph{cannot} exhibit the blue colours ($NUV-r<5.5$) seen in a significant fraction ($\\sim$ 30 percent) of our observed sample. ",
        "introduction": "One of the most important unresolved debates in contemporary astrophysics concerns the formation mechanism of early-type galaxies. In the classical `monolithic' model \\citep[e.g.][]{Larson74,Chiosi2002}, early-type stellar populations form in a single, short, highly efficient burst of star formation at high redshift $(z>>1)$, which is followed by passive ageing to present day. Such a scenario can explain many of the observed properties of early-type galaxies, both at local and high redshift \\citep{Chiosi2002}, without invoking hierarchical merger-driven processes \\citep[e.g.][]{Toomre1972,Toomre77,KWG93,SP99,Benson2000,Cole2000,Hatton2003,Khochfar2003}. Observational evidence that might \\emph{conclusively} rule out either the monolithic or the merger-based scenario has remained elusive (see \\citet{Peebles2002} for a critical review of the two paradigms).  The properties of optical colour-magnitude relations (CMRs) of early-type galaxies have often been taken as evidence in favour of the monolithic scenario. The lack of redshift evolution of the slope and scatter in optical CMRs \\cite[e.g.][]{BLE92,Bower98,Ellis97,Stanford98,Gladders98,VD2000} is consistent with a high formation redshift ($z>2$), as is the evolution in their zero-point. However, the predicted star formation histories (SFHs) of early-type galaxies in the merger paradigm make such a conclusion far less clear cut. For example, the predicted SFHs of (cluster) early-types in the merger scenario are \\emph{quasi-monolithic}, with an \\emph{overwhelming} majority of the stellar mass forming before a redshift of 1 \\citep{Kaviraj2005a}. A careful treatment in the merger framework suggests that, the observed optical CMR can \\emph{also} be reconciled comfortably with merger models. From a photometric point of view, optical colours are, at best, \\emph{degenerate} with respect to the competing theories of early-type galaxy formation and it is very difficult to discriminate between them using optical colours alone \\citep{Kaviraj2005a}.  Evidence does exist for morphological evolution in galaxies, suggesting that formation mechanisms of early-types are \\emph{at least not uniquely monolithic}. Although approximately 80 percent of galaxies in the cores of present day clusters have early-type morphologies \\citep{Dressler80}, a higher fraction of spiral galaxies have been reported in clusters at $0.3<z<0.8$ \\citep[e.g.][]{BO84,Dressler97,Couch98,VD2000}, along with increased rates of merger and interaction events \\citep[e.g.][]{Couch98,VD99}. This is supported by recent results which suggest that the mass density on the red sequence (which is dominated by early-type systems) has doubled since $z=1$ \\citep{Bell2004}. The strengths of age-sensitive spectral indices observed in some early-type systems require luminosity-weighted ages which are consistent with the presence of at least some recent star formation (RSF) in these systems \\citep[e.g.][]{CRC2003,Trager2000a,Trager2000b,Proctor2002}. Furthermore, some nearby early-type galaxies, such as NGC 5128 \\citep[e.g.][]{Rejkuba2001,Peng2003b,Yi2004,Kaviraj2005b,Rejkuba2004}, NGC 205 \\citep{Hodge73,Burstein1988}, NGC 5102 \\citep{Pritchet79,Burstein1988,Deharveng1997} and certain early-type systems recently mapped by the SDSS \\citep{Fukugita2004} exhibit unambiguous signs of significant ongoing or recent star formation. Others, such as NGC 2865 \\citep{Bica1987}, NGC 5128 \\citep{Schiminovich94}, NGC 3921 \\citep{Schweizer1996} provide direct residual evidence of the recent interactions that created them (see also papers by the SAURON collaboration e.g. Falcon-Barroso et al. 2005 and references therein). It is worth noting that, contrary to the traditional notion of early-type galaxies being (cold) gas-poor systems, observational evidence over the last thirty years (e.g. from $IRAS$) have shown that this is \\emph{not} the case. Cold gas, which supplies the fuel for star formation, has been detected in significant numbers of nearby early-type systems \\citep[e.g.][also see the review by \\citet{Knapp1999}]{Knapp1989,Knapp1996,Young2005}. Thus it is perhaps correct to \\emph{expect} RSF in early-type galaxies, rather than consider them inert, passively evolving systems!  Given the apparently-evolving early-type fraction and accumulating spectro-photometric evidence for RSF in nearby early-type galaxies, it is vital to establish whether RSF may only be a sporadic feature of a few early-types or a more widespread phenomenon in the early-type population. The confirmation of low-level star formation \\emph{in a significantly large sample of low-redshift early-type galaxies} would put powerful constraints on models for their formation.  Spectroscopic indicators of RSF are already available, such as the commonly used $H_{\\beta}$ index, higher order Balmer lines such as $H_{\\gamma}$ and $H_{\\delta}$ and the D4000 break. From a photometric point of view, rest-frame ultraviolet flux is highly sensitive to RSF. Given the apparent degeneracy in optical colours, ultra-violet photometry offers the best route to finding tell-tale signatures of RSF in early-type systems at low redshift - if indeed such star formation is present in these systems!  In a recent work, \\citet[][hereafter FS2000]{Ferreras2000} studied a sample of early-type galaxies in the cluster Abell 851 at $z=0.41$. They used $F300W$ (rest-frame near-ultraviolet ($NUV$), 2000 angstroms) and optical photometry to perform one of the first studies of the $NUV$-optical CMR in a sample of early-type galaxies. They found that the slope and scatter of the $NUV$-optical CMR was consistent with some early-types having $\\sim$ 10 percent of their stellar mass in stars younger than $\\sim$ 500 Myrs. Detailed modelling of this data \\citep{Ferreras2002} indicated that secondary bursts of star formation at low redshifts lead to a natural explanation of the large scatter in the $NUV$-optical CMR.  More recently, \\citet{Deharveng2002} studied the rest-frame flux around $2000$ angstroms from a sample of 82 nearby early-types using data from the \\emph{FOCA}, \\emph{SCAP} and \\emph{FAUST} experiments. The focus of this study was primarily to investigate the far-ultraviolet ($FUV$) emission of early-type galaxies, thought to originate from old extreme horizontal branch (EHB) stars and their progeny \\citep{Yi97}. However, their analysis suggested that the $2000-V$ colours in some early-type galaxies are significantly bluer than would be expected from old populations alone and cannot be explained without invoking some RSF.  Low redshift photometry in the far-ultraviolet ($FUV$; 1530 angstroms) and near-ultraviolet ($NUV$; 2310 angstroms) passbands from the GALEX mission (Martin et al. 2005), unprecedented both in terms of its quality and quantity, provides a unique opportunity to study the RSF-sensitive $UV$ emission from a variety of nearby early-type galaxies across a range of luminosities and environments. In this paper we study $\\sim$ 2100 early-type galaxies detected by GALEX, in the redshift range $0<z<0.11$. We focus mainly on the $NUV$ passband, because the $FUV$ is sensitive not only to RSF but also to the $UV$ upturn flux from EHB stars which could be expected in old early-type populations. The $NUV$ is less sensitive to $UV$ upturn and therefore a better RSF indicator.  A preliminary study of the $UV$ emission of early-types, based on GALEX detections of SDSS early-type galaxies listed in the catalog of Bernardi et al. (2003d, B2003 hereafter) was presented in \\citet{Yi2005}. Comparing the photometry of early-types in this sample to the spectral energy distribution (SED) of a strong nearby $UV$-upturn galaxy (NGC 4552), they did not find more than 2 galaxies (out of 162) which fit the typical UV-optical shape of a system with a significant amount of $UV$ upturn flux. They concluded that the large scatter in the $NUV-r$ colours cannot be generated from a scatter in late-stage stellar evolution alone. We also direct readers to \\citet{Boselli2005} who were the first to perform a detailed study of the $UV$ CMR in the Virgo cluster. However, we note that our sample covers a wide variety of environments and we sample almost two magnitudes deeper (in $r$-band) in dense environments \\cite[see][]{Schawinski2005} compared to the Virgo sample of \\citet{Boselli2005}.  We begin this study by describing the construction of an early-type catalog, similar to that used in \\citet{Yi2005}, but in which the morphological classification of objects and removal of potentially $UV$-contaminating AGN are performed in a more robust manner. Using a simple parametrisation of the SFH, we identify galaxies which appear to have had \\emph{some} star formation within the last Gyr. We use a semi-analytical $\\Lambda$CDM model, calibrated to accurately reproduce the (cluster) optical CMR in the redshift range $0<z<1.27$ \\citep[see][]{Kaviraj2005a}, to reproduce the observed $NUV$ CMR. Finally, we discuss the comparative roles of quiescent and merger-driven processes in driving the residual star formation found in some of our early-type sample and investigate whether an \\emph{extreme} monolithic scenario is capable of reproducing the $NUV$ colours of large ($>L_*)$, blue ($NUV-r<5.5$) early-type galaxies in our sample.  %................................................................................................. ",
        "conclusions": "We have studied $\\sim$ 2100 early-type galaxies in the SDSS DR3 which have been detected by the GALEX medium depth (MIS) survey, in the redshift range $0<z<0.11$. The early-type sample has been selected through careful morphological inspection, with potentially $UV$-contaminating AGN removed through the use of both optical spectral data (from the SDSS) and radio data (from the VLA FIRST survey). At a 95 percent confidence level, \\emph{at least} $\\sim$ 30 percent of early-type galaxies in this sample have optical and $UV$ photometry consistent with \\emph{some} recent star formation within the last Gyr.  Our analysis indicates that, while optical CMRs cannot distinguish between early-type galaxies which have had star formation within the last Gyr and those which havent, the $NUV$ CMR is an excellent diagnostic of RSF - any galaxy with $NUV-r<5.5$ has a high likelihood of containing recent star formation (RSF), even after taking into account the possibility of a contribution to the $NUV$ spectrum from $UV$ upturn flux.  %This theoretical estimate, derived from the parametric analysis %presented in this study, is supported by the $NUV$ colour of one %of the strongest local UV upturn systems, NGC 4552 ($NUV-r \\sim$ %5.4) which can be taken as a limiting case to the contamination of %$NUV$ flux by $UV$ upturn.  Comparison of the observations to predictions of a semi-analytical $\\Lambda$CDM hierarchical merger model yields good quantitative agreement (across the $UV$ and optical spectrum), if we assume that (a) very young stars ($<30$ Myrs old) are contained in birth clouds which increase dust extinction by a factor of $\\sim 3$ compared to the ISM alone and (b) that star formation is driven by random (pseudo-instantaneous) starbursts which are Poisson distributed in time, so that there is a small timelag between the last starburst and the observation of the galaxy - the timelags are distributed exponentially in time, with maximum and minimum values of 200 and 20 Myrs respectively.  Combining our parametric analysis of the $UV$ + optical photometry of the observed early-types, with the properties of the predicted population in the semi-analytical model, we conclude that early-type galaxies in the redshift range $0<z<0.11$ are likely to have $\\sim$ 1 to 3 percent of their stellar mass in stars less than 1 Gyr old. The ($V$-band luminosity weighted age) of this recent star formation is $\\sim$ 300 to 500 Myrs.  Finally, we find that monolithically evolving galaxies, where RSF can be produced solely from recycled gas due to stellar mass loss, do not exhibit the blue colours ($NUV-r<5.5$) seen in some large early-type galaxies in our observed sample. While the degeneracy between the monolithic and merger paradigms cannot be broken for red early-types even with $UV$ photometry, a monolithic scenario is a very unlikely channel for the evolution of large blue early-type systems. Such blue galaxies require \\emph{additional} fuel for star formation which must necessarily have an external origin.  %................................................................................................."
    },
    "0601/astro-ph0601604_arXiv.txt": {
        "abstract": "MERLIN measurements of 1.6-GHz OH masers associated with Orion-BN/KL are presented, and the data are compared with data on other masers, molecular lines, compact radio continuum sources and infrared sources in the region. OH masers are detected over an area 30~arcsec in diameter, with the majority lying along an approximately E-W structure that extends for $\\sim$18~arcsec, encompassing the infrared sources IRc2, IRc6 and IRc7. Radial velocities range from --13 to +42~km~s$^{-1}$. The system of OH masers shows a velocity gradient together with non-circular motions.  The kinematics are modelled in terms of an expanding and rotating disc or torus. The rotation axis is found to be in the same direction as the molecular outflow.  There is an inner cavity of radius $\\sim$1300~au with no OH masers.  The inner cavity, like the H$_{2}$O `shell' masers and SiO masers, is centred on radio source I. Some of the OH masers occur in velocity-coherent strings or arcs that are longer than 5~arcsec (2250~au).  One such feature, Stream A, is a linear structure at position angle $\\sim$45\\degr, lying between IRc2 and BN.  We suggest that these masers trace shock fronts, and have appeared, like a vapour trail, 200~yr after the passage of the runaway star BN.  The radio proper motions of BN, source I and source n project back to a region near the base of Stream A that is largely devoid of OH masers. The 1612-MHz masers are kinematically distinct from the other OH masers. They are also more widely distributed and appear to be associated with the outflow as traced by H$_{2}$O masers and by the 2.12-$\\mu$m emission from shocked H$_{2}$. The magnetic field traced by the OH masers ranges from 1.8 to 16.3~mG, with a possible reversal. No OH masers were found associated with even the most prominent proplyds within 10~arcsec of $\\theta^{1}$ Ori C. ",
        "introduction": "The Orion A molecular cloud (L1641) is the nearest giant molecular cloud (GMC) containing the nearest regions of massive star-formation.  Because of its proximity it suffers little galactic extinction, and has been studied in far greater detail than other comparable regions in the Galaxy. The most recent  formation of high--mass stars in the Orion GMC is within the Becklin--Neugebauer--Kleinmann--Low (BN/KL) region, about 1~arcmin Northwest from  the Trapezium star $\\theta^{1}$ Ori C.  The explosive nature of this event is demonstrated dramatically by the infrared H$_{2}$ and [FeII] images of Allen \\& Burton (1993) and more recently Kaifu et al. (2000). `Bullets' of matter are being ejected at hundreds of km s$^{-1}$ to form a bipolar cone of molecular `fingers' whose axis is perpendicular to the hot core traced by NH$_{3}$ emission (Wilson et al. 2000). These  `fingers' are tipped by Herbig--Haro (HH) objects whose shocked gas is visible at optical wavelengths (see Graham, Meaburn \\& Redman 2003 for a recent association of these phenomena).  At radio and millimetre wavelengths there are at least two molecular outflows, one associated with the optical and near-infrared features just described, and a second low-velocity outflow first detected through proper motions of H$_{2}$O masers (Genzel et al. 1981).  The powerful infrared source IRc2 was originally thought to be the source of the outflow, but more recent observations, in particular of SiO masers (e.g. Greenhill et al. 1998), suggest that radio continuum source I of Menten \\& Reid (1995) is the more likely outflow source.   Source I is significantly offset (by $\\sim$0\\farcs5) to the South of IRc2.  Source I and the radio counterpart to BN have proper motions away from each other, corresponding to transverse speeds of 12 and 27~km~s$^{-1}$ respectively (Plambeck et al. 1995;  Rodr\\'{i}guez et al. 2005).  Here and elsewhere we assume a distance of 450~pc.  The radio  proper motions indicate that BN and source I were within 225~au of each other 525$\\pm$30~years ago. Tan (2004) has suggested that the large proper motion of BN is due to its ejection from the $\\theta^{1}$~Ori~C system $\\sim$4000~years ago.  Bally \\& Zinnecker (2005) on the other hand propose that the motions of BN and I are the result of a stellar merger $\\sim$500~years ago, which ejected BN and produced the energy powering the molecular outflow.  Further analysis of the radio proper motion data by G\\'{o}mez et al. (in press) shows that a third source n also has proper motions that project back to the position on the sky where BN and source I were close, $\\sim$500~years ago.  This raises the possibility that n is the third star dynamically neccessary to eject BN from the multiple system.  Distinguishing between these different possibilities is crucial to understanding the history of recent star-formation in the region.  We have observed hydroxyl (OH) masers in the Orion-BN/KL region as part of an ongoing study of masers associated with molecular outflows from massive young stars (Hutawarkorn \\& Cohen 2005, and references therein).  OH masers have the potential to trace physical conditions, including magnetic fields which are detected and measured through Zeeman splitting and OH maser polarization. Previous MERLIN OH observations by Norris (1984) detected 80 1665-MHz masers distributed over a region of 10~arcsec, which were interpreted in terms of a rotating torus surrounding IRc2. Johnston, Migenes \\& Norris (1989) detected 175 masers spread over a region of 20~arcsec, with the majority in an E-W structure 14~arcsec in extent. The masers were found in clusters or  `clumps'  that correlated with NH$_{3}$ emission. The present MERLIN observations more than double the number of OH masers known and show that the distribution is even more extensive in position and velocity (Section~\\ref{distribution}).  The association of the masers with radio continuum sources is described in Sections~\\ref{distribution} and~\\ref{arcs}, while the association with near- and mid-infrared sources is described in Section~\\ref{comparison}.  Secondary targets in the same MERLIN observations were the `proplyds' (O'Dell, Wen \\& Hu 1993) in the vicinity of $\\theta^{1}$ Ori C. These are the solar system--sized circumstellar envelopes discovered by Laques \\& Vidal (1979) which surround low--mass YSOs (Churchwell et al 1987; Meaburn 1988; McCaughrean \\& Stauffer 1994; O'Dell \\& Wong 1996; Bally et al. 1998a). These stellar cocoons have dusty molecular disks with dense ($\\geq$ 10$^{6}$ cm$^{-3}$) surfaces ionized by the intense flux of Lyman radiation from $\\theta^{1}$ Ori C (see Graham et al. 2002). In these circumstances maser emission could be anticipated from the molecular disks which, if detectable, could be an important tool for investigating proto-planetary environments.  Upper limits for such emission are reported in Section~\\ref{proplyds1}. ",
        "conclusions": "The distribution of OH masers in the Orion-BN/KL region is far more extensive than previously realized, covering a region of 30~arcsec extent and a radial velocity range from --13 to +42~km~s$^{-1}$.  The bulk of the emission can be modelled in terms of a rotating and expanding torus, centred on IRc2 or radio source I, with an inner cavity of $\\sim$1300~au radius.  The rotation axis has the same position angle and inclination to the line-of-sight as the molecular outflow and the large scale magnetic field inferred from mm- and submm-polarization (Section~\\ref{model}).  The dynamical timescale is similar to that of the explosive event that produced the widespread shocked H$_{2}$ emission.  It is likely that the OH masers trace the interaction between the low-velocity molecular outflow and the molecular hot core, and that they acquired their expansional motions in the same event that produced the outflow.  Of particular interest is a string of masers, Stream A, at $\\sim$21~km~s$^{-1}$, that extends at position angle $\\sim$45\\degr between IRc2 and BN, in the direction of the radio proper motions of these two dominant sources.  We suggest that Stream A may have appeared, like a vapour trail, in the wake of the runaway star BN.  The proper motions of BN and sources I and n project back to the base of Stream A (Fig.~\\ref{ra-dec}), $\\sim$4~arcsec Northwest of the centre of the OH maser torus, a position that is largely devoid of masers.   The 1612-MHz masers have a more widespread distribution than the other OH masers, with kinematics that are more like those of the H$_{2}$O masers associated with the outflow.  Many of these 1612-MHz masers are spatially associated with fingers of shocked H$_{2}$ emission (Fig.~\\ref{subcore}).  These OH masers are thought to require relatively low gas and dust temperatures for their inversion (Section~\\ref{1612}).  Apart from the OH 1612-MHz masers, the other OH masers have complementary distributions and kinematics to the  H$_{2}$O masers (Figs.~\\ref{ra-dec}--\\ref{dec-vel}), with essentially no overlap.  OH is also spatially and kinematically distinct from the class I methanol 25-GHz masers.  A possible correspondence between OH 1612-MHz and class II methanol 6.7-GHz masers reported by Voronkov et al. (2005) requires further study at higher angular resolution.  The magnetic field strength in the OH maser regions ranges from 1.8 to 16.3~mG, with a possible field reversal (Section~\\ref{zeeman}).   A search for OH masers associated with the proplyds around  $\\theta^{1}$ Ori C yielded only upper limits, which may indicate a low OH abundance in this more evolved region.   \\vspace{1.2cm} \\noindent \\large {\\bf ACKNOWLEDGMENTS}\\normalsize \\vspace{0.5cm}  We thank the anonymous referee for helpful comments and in particular for drawing our attention to the paper by G\\'{o}mez et al. We thank Busaba Hutawarakorn for assistance with the data reduction, Ralph Shuping for providing a high resolution version of the Keck image in Fig.~5, and Masa Hayashi for providing high resolution Subaru data used in Figs. 6 and 7. NG gratefully acknowledges the Thai Astronomy Co-operative Research Network  for their support. MERLIN is a national facility operated by the University of Manchester on behalf of PPARC."
    },
    "0601/astro-ph0601432_arXiv.txt": {
        "abstract": "A search for RR Lyrae stars has been conducted in the publicly available data of the Northern Sky Variability Survey ({\\it NSVS}). Candidates have been selected by the statistical properties of their variation; the standard deviation, skewness and kurtosis with appropriate limits determined from a sample 314 known RRab and RRc stars listed in the GCVS.  From the period analysis and light curve shape of over 3000 candidates 785 RR Lyrae have been identified of which 188 are previously unknown. The light curves were examined for the Blazhko effect and several new stars showing this were found. Six double-mode RR Lyrae stars were also found of which two are new discoveries. Some previously known variables have been reclassified as RR Lyrae stars and similarly some RR Lyrae stars have been found to be other types of variable, or not variable at all. ",
        "introduction": "Statistical studies of the relative numbers of the different classes of RR Lyrae variable stars and especially the incidence rate of multi-periodicity may give indications of the metallicity of different stellar systems and of their evolution \\citep[see e.g. ][]{mos}. Exhaustive studies have been done in the Magellanic Clouds by the {\\it MACHO} collaboration \\citep[][ and 2003]{macho} and the {\\it OGLE} survey \\citep{oglelmc}. Other studies have searched for RR Lyrae stars in the Galaxy, such as {\\it QUEST} \\citep{quest} and also {\\it OGLE} \\citep{collinge}. Most of the stars found in these studies are faint, and limited to a small region of sky (the galactic bulge and the equator for {\\it OGLE} and {\\it QUEST} respectively). The Robotic Optical Transient Search Experiment \\citep[{\\it ROTSE-1}, ][]{rotse} found a fairly large number of previously unknown bright (magnitude $< 15$) RR stars in a part of the sky.  This paper sets out to extend this search for field galactic RR Lyrae stars to the whole northern sky in the {\\it ROTSE-1} data, made publicly available via the Internet \\citep[Northern Sky Variability Survey - {\\it NSVS},][]{skydot}.   \\section[]{Methodology}  With only photometric data, and no spectral information, the type of a short period variable can only be determined once the (phased) light curve is known, and hence only once the period is known.  However, determining the period from sparse data is a computationally demanding process.  Therefore it was decided to limit the number of objects by statistical parameters involving much less computation: standard deviation, skewness, kurtosis and the mean square of successive differences of the magnitude data.  By looking at the values for these statistics for the RR Lyrae stars listed in the General Catalogue of Variable Stars \\citep[GCVS, ][ and its online edition]{gcvs}, limiting conditions were then derived.  The aim was to set these limits as strict as possible, so that not too many objects needed to be checked, but still to include the majority of RR Lyrae stars. Those GCVS stars that did not follow the criteria, can then provide an estimate for the completeness of the survey.   \\subsection[]{Data}  The {\\it ROTSE-1} was an unfiltered CCD survey of the sky from the north pole to declination $\\sim -38$, reaching magnitude $\\sim 15$ with varying levels of completeness. The survey lasted nominally for one year but depending on the circumstances coverage of individual objects may be significantly less than this. Objects typically have 100 to 500 measurements with a median photometric accuracy of 0.02 mag for 10th magnitude stars and a positional accuracy of 2\". The spatial resolution of 14\" compromises the photometry in crowded fields, typically with $|b| < 20$, but also at higher galactic latitudes for stars with companions within $\\sim 45\"$.  The data are publicly available from the Sky Database for Objects in Time-Domain (SkyDOT) web site \\citep{skydot} and it is possible to select data with respect to 8 extraction flags and 7 photometric correction flags. The default selection for good data sets all but one, {\\sc PATCH}, of the photometric correction flags and only one, {\\sc SATURATED}, of the extraction flags. However, experience of working with the data has shown that observations with the extraction flag {\\sc APINCOMPL} set are often completely out of range and should be rejected. On the other hand data with the photometric correction flag {\\sc RADECFLIP} set are often indistinguishable from the other data. So the data have been selected with the {\\sc SATURATED} and {\\sc APINCOMPL} flags set and the {\\sc PATCH} and {\\sc RADECFLIP} flags unset.  It was also decided that only stars with 100 good {\\it NSVS} observations or more were to be considered, in order to get good statistics and reliable period determinations, and to possibly detect multiperiodicity (double-mode pulsation or a Blazhko effect). With fewer data points, statistics may be influenced substantially by erroneous observations, such as those introduced by e.g. a close companion. Also, it is then not always possible to derive the correct period: in view of the sampling frequency, and the rather short total time span of the available data (less than a year) alias frequencies will be more important.  RR Lyrae stars are fairly blue stars (spectral types A and F). The {\\it NSVS} survey however observed only in one colour (unfiltered CCD), so colour information has to be retrieved from another source, such as {\\it 2MASS} \\citep{2mass}. The {\\it NSVS} positions are not very accurate however, and matching them to {\\it 2MASS} coordinates may be troublesome in crowded fields. In view of this and of reddening aspects, it was decided not to use colour information as a filter.   \\subsection[]{Control group: GCVS stars}  The GCVS stars that were to be considered for the control group, had to be well-known and have an accurate position. Therefore only the GCVS types RRab or RRc were taken (no RR, RR:, RRab: or RRc: stars, i.e. the GCVS classification should be precise enough to give the exact subtype). Because of the limited number of RRd stars in the GCVS (the RR(B) class), these were not taken into account either. For practical purposes, as far as their statistical parameters are concerned, these double-mode stars can be considered to be RRc stars. The known RR Lyrae stars in the GCVS were further limited to the constellations And to Ori for which precise positions had been determined by the GCVS team at the time this study started. Many stars in other constellations did not have accurate enough coordinates, which could lead to misidentifications with the {\\it NSVS} stars. It is however still possible that a faint RR Lyrae star unobservable by the {\\it ROTSE} camera, lies close to a brighter (constant) companion, leading to a false identification. Because of this, the success rates for finding an RR star, may be underestimated. On the other hand, especially at fainter magnitudes, some stars which should have been detectable, will not have been registered, thus overestimating the success rate.  With the above restrictions imposed, 582 {\\it NSVS} objects were identifed as GCVS RR Lyrae stars by their {\\sc HTM} identification \\citep[see][]{skydot}. {\\it NSVS} synonyms, the same object observed in overlapping {\\it NSVS} fields, have been counted separately here. This will be done also for the remainder of this section, as it will not change the statistics very much.  The further restriction that there needed to be at least 100 good data points limited the sample to 314 objects (273 RRab and 41 RRc), or 54\\% of the total number of stars identified. Compare this to the overall 42\\% of objects with at least 100 good points (8393519 out of a total of 19995106 {\\it NSVS} objects). 60\\% of the GCVS stars that are on average brighter than magnitude 14, have more than 100 observations, and 68\\% if only objects North of the equator are counted.   \\subsection[]{Skewness}  RRab stars have a typical light curve, spending more time near minimum than near maximum, while RRc stars have more symmetric light curves. It distinguishes them from eclipsing binaries, by far the most common type of variable star found in the {\\it NSVS} database. As a result, the distribution of magnitudes of an RRab star shows a negative skewness, while those of an RRc star shows a skewness near zero (note that a perfect sine curve has a skewness equal to 0). The distribution of the skewness values for the selected GCVS stars is given in Fig.~\\ref{skew}.  \\begin{figure} \\includegraphics[width=84mm]{RRNSVSfig1.eps} \\caption{Distribution of the skewness values of the {\\it NSVS} data for RRab and RRc stars from the GCVS.} \\label{skew} \\end{figure}  The maximum value for the skewness statistic was set to 0.5, which selects 303 out of 314 stars (96\\%). Most of the GCVS objects with higher skewness values did not satisfy the standard deviation criterion either. Therefore the chosen restriction seems to be an appropriate one. The objects not selected include HM~Aql (probably a constant star, see below) and UU~Cam (often referred to as an EW type variable) and a few objects near the magnitude limit resulting in a disproportionally high number of bright observations, and hence a larger skewness. One star had two very faint data points, outside its normal range, most probably data errors, also influencing skewness to rise abnormally.  Some of the objects selected with positive skewness are eclipsing variables of the W~UMa type (EW). Especially for skewness values $> 0.5$, most of the variables are EW stars or other eclipsing binaries. Some of these have been reported by \\cite{ecl}. In a number of cases it is impossible to distinguish between EW and RRc on the basis of the light curve alone. These stars have not been withheld.   \\subsection[]{Mean square of successive differences}  The mean square of successive differences ({\\sc MSSD}) gives an indication for the timescale of the variation of a star as compared to the sampling timescale. The statistic considered further in this paper is in fact the rescaled value \\begin{equation} \\theta = 1-{\\sc MSSD}/2\\sigma^2, \\end{equation} with $\\sigma$ the standard deviation. For a purely random distribution of the data $\\theta = 0$ \\citep{cuypers}, and for a star that only brightens or fades during the total time interval, $\\theta$ will approach 1. RR Lyrae stars are rapidly changing stars (their periods range between about 0.2 and 1 day), compared to the frequency of observation by {\\it ROTSE} (up to 5 points per night). Therefore a raw light curve will show almost random variation. The criterion $\\theta < 0.65$ was chosen.  307 (98\\%) of the selected GCVS stars satisfy this criterion.  Because of the {\\it ROTSE} observing regime, in some cases it is possible that a star with a period close to an integer fraction of a day, will show a raw light curve that resembles one of a much longer period star. This extreme aliasing is the case for RU~Boo (period 0.4927d), V1949~Cyg (0.4989d), AW~Lyr (0.4975d), BT~Leo (0.4997d) and FY~Aqr.  For the latter, a period has not been given in the GCVS, but from {\\it ASAS3} \\citep{asasI} and {\\it NSVS} data a period of 1.0229 days may be derived; see also \\cite{dom04} for another period determination of this star.   \\subsection[]{Kurtosis}  Kurtosis is a measure of the \"peakedness\" of a distribution and it was found to be useful here as a discriminator against stars with extreme values. A maximum value of 4.5 was set for the kurtosis value of the magnitudes (a perfect sine has kurtosis = 1.5). It excludes stars which have some unusually faint or bright data points. This criterion effectively removes stars with bad data points making the standard deviation erroneously large. It may however hide true variation in some particular cases, as some stars with highly deviating points which are really variable are excluded as well. This did not affect the majority of the RR Lyrae stars as the criterion selects 301 (96\\%) of the GCVS stars, All of the objects that failed this test also violated other criteria as well. Overall with all criteria so far applied (without limits on the standard deviation) 287 (255 RRab and 32 RRc) out of 314 (91\\%) are left.     \\subsection[]{Standard deviation}  Being the most important parameter for the recognition of a variable star, a cut is difficult to define for the standard deviation, as it strongly depends on the average magnitude of the stars, especially for fainter objects. From a Fourier fit to the observations, \"theoretical\" standard deviations (i.e. without observational errors) for the selected RR Lyrae stars were found to be always larger than 0.1 mag, unless the star was in a close pair which could not be resolved by the instrument. Most RRab stars have \"theoretical\" standard deviations between 0.15 and 0.30 mag, RRc variables between 0.10 and 0.20 mag (note that in the case of a perfect sine curve, the standard deviation $\\sigma = \\Delta m/2\\sqrt2$, with $\\Delta m$ the total amplitude from minimum to maximum).  A survey for Cepheids in Milky Way fields of the {\\it NSVS} \\citep{cep} has shown that the standard deviation needed to be at least about twice that of the average standard deviation at the star's magnitude, to be certain the star is a genuine variable star, and not one with unusually high observational scatter.  Therefore this restriction was chosen. Without imposing it, the number of stars to be checked grows exponentially with decreasing standard deviation. True variables exist which have a lower standard deviation, as shown by some of the GCVS stars, but these cases are rather rare and/or hard to confirm, as true variation gets masked by observational scatter. In addition, inaccuracies on the other statistics calculated, increase as well.  The cutoff for the standard deviation $\\sigma$ is graphically illustrated in Fig.~\\ref{stdev}. It plots $\\sigma$ for the 314 GCVS RR Lyrae stars against their average magnitude. The thin full lines represent the average $\\sigma$ (lower line) and the chosen lower limit $2\\sigma$ (upper line) at the given magnitude for all stars with more than 100 data points in the {\\it NSVS} database. 0.1\\% of the stars have a value of $\\sigma$ above the upper dashed line, and 1\\% above the lower dashed line.  Almost all the RR Lyrae stars belong to the latter group.  It is probably worth looking at the brightest stars with low standard deviation. HM~Aql \\citep{harwood} looks constant in {\\it NSVS} data as well as in {\\it ASAS3} data \\citep{asasV}. Also V1510~Cyg and HU~Cam appear to be constant in {\\it NSVS} data. There might be an identification problem for these stars or a bright close companion.  \\begin{figure} \\includegraphics[width=84mm]{RRNSVSfig2.eps} \\caption{{\\it NSVS} standard deviation compared to magnitude for the RRab (open circles) and RRc stars (filled circles) selected from the GCVS. The lower and upper dashed lines represent respectively the 99 and 99.9 percentiles of the {\\it NSVS} standard deviations $\\sigma$ as a function of magnitude. The bottom thin full line gives the average standard deviation for a given magnitude ($1 \\sigma$ level). The upper thin full line is the $2 \\sigma$ level and it represents the chosen cutoff limit. The total number of stars with at least 100 data points, is given per 0.1 mag bin by the thick line (right axis).} \\label{stdev} \\end{figure}  \\subsection[]{Detection probability}  Only 205 RRab and 16 RRc stars out of the total of 314 GCVS objects are left after applying the criterion on the standard deviation. So the overall detection possibility for RRab stars is 75\\%, for RRc stars it is only 40\\%. However, for stars on average brighter than mag 13, the success rate is 91 and 83\\% respectively. For stars brighter than magnitude 14, these rates are reduced to 83 and 57\\%. If these numbers are then applied to the requirement that the stars needed to be observed at least 100 times, it follows that 56\\% of the RRab stars on average brighter than magnitude 14 in the northern hemisphere are detectable with the assumed criteria and nearly 40\\% of the RRc stars. These detection probabilities will be compared with the results from the {\\it ASAS3} survey. Because of the different detection possibilities, it is hard to accurately determine the ratio between the number of RRab and RRc stars. ",
        "conclusions": "This paper has enlarged the known number of galactic field RR Lyrae stars in the northern sky. The number of known stars showing multiperiodicity (Blazhko effect and double-mode pulsation) has been enlarged as well.  The number of RR Lyrae stars brighter than magnitude 14 has been estimated, showing that there is an over-abundance of RRab stars in the southern hemisphere."
    },
    "0601/astro-ph0601118_arXiv.txt": {
        "abstract": "The evolutionary history of  galaxies is thought to be strongly conditioned by the environment. In order to quantify and set limits on the role of nurture one must identify and study an isolated sample of galaxies. But it is not enough to identify a small number of the \"most isolated\" galaxies. We begin with 950 galaxies from the Catalogue of Isolated Galaxies (Karachentseva 1973) and reevaluate isolation using an automated star-galaxy classification procedure on large digitised POSS-I fields. We define, compare and discuss various criteria to quantify the degree of isolation for these galaxies: Karachentseva's revised criterion, local surface density computations and an estimation of the external tidal force affecting each isolated galaxy. Comparison of multi-wavelength ISM properties, in particular the H$\\alpha$ emission line, will allow us to separate the influence of the environment from the one due to the initial conditions at formation. ",
        "introduction": "The role of the environment on galaxy evolution is still not fully understood. In order to quantify and set limits on the role of nurture one must identify and study a sample of isolated galaxies. The AMIGA project, \"Analysing the Interstellar Medium of Isolated  Galaxies\", is doing a multi-wavelength study of a large sample of isolated galaxies in order to examine their star formation activity and interstellar medium (Verdes-Montenegro et al. 2005, http://www.iaa.csic.es/AMIGA.html). ",
        "conclusions": ""
    },
    "0601/astro-ph0601268_arXiv.txt": {
        "abstract": "We report the definite spectroscopic identification of $\\simeq40$ OB supergiants, giants and main sequence stars in the central parsec of the Galaxy. Detection of their absorption lines have become possible with the high spatial and spectral resolution and sensitivity of the adaptive optics integral field spectrometer SPIFFI/SINFONI on the ESO VLT.  Several of these OB stars appear to be helium and nitrogen rich. Almost all of the $\\simeq80$ massive stars now known in the central parsec (central arcsecond excluded) reside in one of two somewhat thick ($\\langle|h|/R\\rangle\\simeq0.14$) rotating disks.  These stellar disks have fairly sharp inner edges ($R\\simeq1\\arcsec$) and surface density profiles that scale as $R^{-2}$. We do not detect any OB stars outside the central $0.5$~pc. The majority of the stars in the clockwise system appear to be on almost circular orbits, whereas most of those in the `counter-clockwise' disk appear to be on eccentric orbits.  Based on its stellar surface density distribution and dynamics we propose that IRS~13E is an extremely dense cluster ($\\rho_\\mathrm{core}\\gtrsim3\\times10^8\\sunmass$~pc$^{-3}$), which has formed in the counter-clockwise disk.  The stellar contents of both systems are remarkably similar, indicating a common age of $\\simeq6\\pm2$~Myr.  The K-band luminosity function of the massive stars suggests a top-heavy mass function and limits the total stellar mass contained in both disks to $\\simeq1.5\\times10^4\\sunmass$. Our data strongly favor in situ star formation from dense gas accretion disks for the two stellar disks.  This conclusion is very clear for the clockwise disk and highly plausible for the counter-clockwise system. ",
        "introduction": "The Galactic Center (GC) is a unique laboratory for studying galactic nuclei. Given its proximity, processes in the GC can be investigated at resolutions and detail that are not accessible in any other galactic nucleus (unless otherwise specified we adopt a distance of 8~kpc for simplicity of comparison to earlier work; we specifically use the most recent value $R_0=7.62\\pm0.32$~kpc by Eisenhauer et al. 2005 when the $\\simeq5\\%$ error could lead to a significant bias).  The GC has many features that are thought to occur in other nuclei (for reviews, see Genzel \\& Townes 1987; Morris \\& Serabyn 1996; Mezger, Duschl \\& Zylka 1996; Alexander 2005).  It contains the densest star cluster in the Milky Way intermixed with a bright \\HII\\ region (Sgr~A West or the `mini-spiral') and hot gas radiating at X-rays. These central components are surrounded by a $\\simeq1.5$~pc ring/torus of dense molecular gas (the `circum-nuclear disk', CND). At the very center lies a very compact radio source, Sgr~A*. The short orbital period of stars (in particular the B star S2) in the central arcsecond around Sgr~A* show that the radio source is a $3$--$4\\times10^6\\sunmass$ black hole (BH) beyond any reasonable doubt (Sch\\\"odel et al. 2002; Ghez et al.  2003).  The larger Galactic Center region contains three remarkably rich clusters of young, high mass stars: the Quintuplet, the Arches, as well as the parsec-scale cluster around Sgr~A* itself (Figer 2003).  In seeing-limited near-infrared images of the central region of the Galactic Center, several bright sources dominate the $\\simeq20\\arcsec\\times20\\arcsec$ field centered on Sgr~A*. Among these, the IRS~16 cluster\\footnote{The sources named ``GCIRS'' for Galactic Center Infrared Source are often referred to simply as ``IRS'' sources in the GC-centric literature.} (Becklin \\& Neugebauer 1975) is a bright source of broad \\HeI~2.058~\\micron\\ line emission (Hall, Kleinmann \\& Scoville 1982).  IRS~16 has been since then resolved into a cluster of about a half a dozen stars (Forrest et al. 1987, Allen, Hyland \\& Hillier 1990, Krabbe et al. 1991, 1995, Tamblyn et al. 1996, Paumard et al. 2001). These appear to be post main-sequence OB stars in a transitional phase of high mass loss (Morris et al.  1996), between extreme O supergiants and Wolf-Rayet (WR) stars (Allen et al. 1990; Najarro et al. 1994, 1997; Trippe et al.  2005). They have been classified as Ofpe/WN9 stars (Allen, Hyland, \\& Hillier 1990) and have been suggested as Nitrogen-rich OB stars (OBN) stars by Hanson et al. (1996) and as Luminous Blue Variables (LBV) candidates by Paumard et al.  (2001).  There is no a-priori incompability between these tentative classifications that are based on different properties.  Several dozens of even more evolved WR stars have been observed in the same region (e.g. Krabbe et al.  1995; Blum, Sellgren \\& DePoy 1995; Paumard et al. 2001). The lack of OB stars in these earlier studies is puzzling. The question is whether this lack is due to a true depletion or is merely a selection effect due to veiling of the weak absorption lines in near main sequence stars by bright nebular emission.  Adaptive optics (AO) spectroscopy of the center-most arcsecond around Sgr~A* (mostly devoid of nebular emission) has already revealed a dozen massive stars. These stars appear to be main sequence late O and B stars and orbit the central black mass at distances as short as a few light-days (Sch\\\"odel et al.  2003; Ghez et al. 2003, 2005; Eisenhauer et al. 2005).  These observations show that massive star formation has occurred at or near the Galactic Center within the last few million years. This is surprising. All obvious routes to creating or bringing massive young stars in(to) the central region face major obstacles. In situ star formation, transport of stars from far out, scattering of stars on highly elliptical orbits and rejuvenation of old stars due to stellar collisions and tidal stripping have all been proposed and considered (for a recent review of the rapidly growing body of literature see Alexander 2005). No explanation at this point is the obvious winner (or loser). Perhaps the two most prominent and promising scenarios for explaining the young massive stars outside the central cusp, at radii of 3--$10\\arcsec$ from Sgr~A*, are \\begin{enumerate} \\item the `\\emph{in situ, accretion disk}' scenario (Levin \\& Beloborodov 2003; Genzel et al.  2003; Goodman 2003; Milosavljevic \\& Loeb 2004; Nayakshin \\& Cuadra 2005). Here the proposal is that stars have formed near where they are found today, very close to the central black hole. However, in situ star formation is impeded by the tidal shear from the central black hole and surrounding dense star cluster. To overcome this shear, gas clouds have to be much denser ($\\simeq10^{12}R_{1\\arcsec}^{-3}$~cm$^{-3}$) than currently observed (Morris 1993).  The tidal shear can be overcome if the mass accretion was large enough at some point in the past -- perhaps as the consequence of the in-fall and cooling of a large interstellar cloud -- such that a gravitationally unstable (outside of a critical radius) disk was formed. The stars were formed directly out of the fragmenting disk; \\item the `\\emph{in-spiraling star cluster}' scenario (Gerhard 2001; McMillan \\& Portegies Zwart 2003; Portegies\\linebreak Zwart et al. 2003; Kim \\& Morris 2003; Kim, Figer, \\& Morris 2004; G\\\"urkan \\& Rasio 2005). Here the idea is that young stars were originally formed outside the hostile central parsec and only transported there later on. Individual transport of stars by two body relaxation and mass segregation from further out takes too long a time ($\\simeq10^{7.5}$--$10^9$ years: Alexander 2005).  Stars in a bound, massive cluster can sink in much more rapidly owing to dynamical friction (Gerhard 2001). To sink from an initial radius of a few parsec or more to a final radius of $\\ll1$~pc within an O star lifetime (a few Myr) requires a cluster mass $>10^5\\sunmass$.  To prevent the final tidal disruption of such a cluster at too large a radius {--} resulting in the deposition of its stars there {--} the core of the original star cluster also has to be much denser ($>10^7\\sunmass$~pc$^{-3}$) and more compact ($\\ll1$~pc) than any known cluster.  However, as a helpful by-product, dynamical processes in such a hypothetical super-dense star cluster may then lead to the formation of a central, intermediate mass black hole (IMBH; Portegies-Zwart \\& McMillian 2002; G\\\"urkan et al, 2004). Such a black hole may help to stabilize the cluster core against tidal disruption and lessen the high density requirement somewhat (Hansen \\& Milosavljevic 2003). \\end{enumerate}  The `paradox of youth' in the central S-star cluster, with $>15$ apparently normal main sequence B stars residing in tightly bound orbits in the central light month around the central black hole, probably requires yet another explanation (recently Ghez et al.  2003, 2005; Genzel et al. 2003; Hansen \\& Milosavljevic 2003; Gould \\& Quillen 2003; Alexander \\& Livio 2004; Eisenhauer et al. 2005; Alexander 2005; Davies \\& King 2005). Perhaps the most promising route to get the B stars into the central arcsecond is a scattering process from the reservoir of massive, young stars at $\\simeq3$--$10\\arcsec$ (e.g. Alexander \\& Livio 2004).  To test these proposals, the detailed properties and dynamics of the massive stars in the central parsec must be studied. These properties include exact stellar type, spatial distribution and 3D space velocities. For this purpose, high-resolution imaging and spectro-imaging are required. The new adaptive optics assisted, near-infrared integral field spectrometer on the ESO-VLT, SPIFFI/SINFONI (Eisenhauer et al.  2003b; Bonnet et al. 2004) represents a key new capability for addressing the issues discussed above. We report in this paper SPIFFI/SINFONI observations in 2003, 2004 and 2005 that give important new information on the location, dynamics and evolution of the massive, early type stars in the central parsec.  We begin by discussing the SPIFFI/SINFONI observations and data analysis in Sect.~\\ref{sect:obs}. This is followed by presentation of our results in Sect.~\\ref{sect:res}.  In Sect.~\\ref{sect:discuss} we discuss the implications of our findings. Further technical details are presented in the Appendices. ",
        "conclusions": ""
    },
    "0601/astro-ph0601097_arXiv.txt": {
        "abstract": "Cold molecular gas has recently been found is several cooling flow clusters cores with single dish telescopes. High spatial resolution imaging of some of these clusters then revealed the peculiar morphology and dynamics of the CO emission lines, pointing out a perturbed very cold component in the cluster centers. We report here the observations of NGC 1275, in the Perseus cluster of galaxies. This object is the strongest cooling flow emitter in the millimeter. The 9 dual polarization pixels of the HERA focal plane array, installed on the 30m telescope, enabled to image the large scale emission of the cold molecular gas which is found to follow the very peculiar Halpha filamentary structure around the central galaxy. We discuss here this association and the non-rotating dynamics of the cold gas that argue for a cooling flow origin of the molecular component. ",
        "introduction": "These observations were made from 1st to 3rd January 2005 at the IRAM-30m telescope. We used the HERA (HEterodyne Receiver Array), a focal array of 18 SIS receivers, 9 for each polarisation, tuned at the CO(2-1) line, for NGC 1275 (226.56 GHz). The sampling was 6 arsec (full sampling). We built four such maps, covering the central 138x138$''$, and also a 5th one covering 66x66$''$ towards the north. Figure \\ref{figure1} shows in contours the CO(2-1) emission, overlaid on a Halpha image of the filamentary structure pointed out by Conselice et al. (2001). ",
        "conclusions": ""
    },
    "0601/astro-ph0601083_arXiv.txt": {
        "abstract": "We present near-infrared $J$, $H$, and $K$ images of four embedded stellar clusters in the Galaxy.  We find a significant fraction of pre-main-sequence stars present in at least one of the clusters.  For the clusters dominated by main-sequence stars, we determine the initial mass function (IMF) both by using the $K$ luminosity function and a global extinction correction and by deriving individual extinction corrections for each star based on their placement in the $K$ vs. $H-K$ color-magnitude diagram.  Based on our IMFs we find a significant discrepancy between the mean IMF derived via the different methods, suggesting that taking individual extinctions into account is necessary to correctly derive the IMF for an embedded cluster. ",
        "introduction": "\\label{sec:intro}  Embedded clusters are increasingly recognized as vital sites of star formation for both low- and high-mass stars.  Recent studies indicate that clusters may account for 70-90\\% of star formation and that embedded clusters (those still partially or fully enshrouded in their natal molecular cloud) may exceed the number of older, non-embedded open clusters by a factor of $\\sim$20 \\citep{elmegreen00,lada03}.  The stellar content of embedded clusters within well-known star formation regions can now be probed, where high extinction ($A_{V} \\gtrsim 10$) prohibits studies at optical wavelengths.  The IMF of such clusters has generally been found to be consistent with a Salpeter value with a slope of $\\Gamma = -1.35$ \\citep[e.g][]{okumura00,blum01,figueredo02} although outliers have been found as well, generally with flatter slopes than the Salpeter value \\citep[e.g.][]{porras99}.  Although near-infrared (NIR) spectral classification of massive stars is possible \\citep{hanson96}, in most cases determinations of the IMF from NIR data rely heavily, if not exclusively, on photometry and use spectroscopy only to obtain reliable masses of the few brightest and most massive stars in a cluster if at all.  Since these results thus depend on stellar evolutionary models as well as details of the handling of extinction, this raises concerns about to what extent the IMF depends on the methodology employed. \\citet{massey03} cites the example of NGC 6611, where two separate analyses of the same data \\citep{hillenbrand93,massey95} using different treatments of extinction produced IMFs differing by more than the formal 1$\\sigma$ errors would suggest ($\\Gamma = -1.1 \\pm 0.1$ and $\\Gamma = -0.7 \\pm 0.2$.) Similarly, the IMF for the G305+0.2 embedded cluster differs by more than the errors between the value derived from the K luminosity function (KLF; $\\Gamma = -1.5 \\pm 0.3$) and that derived using the color-magnitude diagram ($\\Gamma = -0.98 \\pm 0.2$) \\citep{paper1}.  Claims that variations in the IMF exist, whether based on individual extreme clusters such as the Arches or a general analysis of the data \\citep{scalo98}, must thus be handled with care to compare only results based on similar methodology.  The final release of the Two Micron All Sky Survey (2MASS) has fostered studies \\citep[e.g.][]{db00,db01,bdb,ivanov} which can probe a much larger portion of the Galaxy for previously unknown embedded stellar clusters and significantly increased the number of known embedded clusters.  A compilation of some of these results along with previously known embedded clusters is presented in \\citet{porras03} who find that $\\sim 80\\%$ of the stars in their sample are found in ``large clusters'' of more than 100 stars, despite the rarity of such clusters.  However, these studies are not foolproof, and compilations of ``embedded clusters'' based purely on the 2MASS data without followup must be treated with caution.  The studies based solely on stellar density criteria \\citep[e.g.][]{db00,db01} have been found in followup work by different groups \\citep[e.g.][]{db03, paper1, borissova} to have only about a 50\\% success rate toward the inner Galaxy where the stellar background is high.  We have performed an independent search of the 2MASS archive, searching the Point Source Catalog for regions of higher stellar density than the background (determined locally within a 5\\arcmin\\ radius) which are redder in $H-K$ than the local field.  This selects for embedded clusters, with the color criteria helping to eliminate chance superpositions and regions of low extinction.  A large background radius and the use of color selection are critical to the automated identification of embedded clusters, but even color selection can fail in regions of high background stellar density, where patchy extinction can mimic clusters.  In \\S\\ref{sec:obs} we present the observations and data reduction for four embedded clusters found in the 2MASS Point Source Catalog, in \\S\\ref{sec:analysis} we present the $K$-band luminosity functions (KLF) and initial mass functions (IMF) for the clusters, and in \\S\\ref{sec:compare} we address the issue of systematic differences between different methods of deriving the IMF for embedded clusters, for our clusters as well as IMFs from the literature. ",
        "conclusions": "We present NIR images of four embedded clusters in the outer Galaxy.  In the case of the cluster near IRAS 06058+2158 the number of stars with NIR excess indicates a pre-main-sequence fraction between 60\\% and 80\\% and an age of less than 3 Myr; the other three clusters show less nebular emission and fewer stars with NIR excess indicating an older age.  We compute the IMF for the three clusters dominated by main-sequence stars, in each case using both a KLF-based method relying on a single extinction value for the cluster and using only $K$ band data and a CMD-based method where an individual extinction value is calculated for each star.  We found a statistically significant difference between the two values in two of the three cases, prompting us to examine IMF values of embedded clusters from the literature to determine whether systematic effects are at work.  We found a significant difference in the mean value of the IMFs for embedded clusters derived from methods that handle extinction individually compared with those that adopt a single value for the extinction. Although a larger sample would help to make this claim more robust, since many of the results come from a single study \\citep{porras99} and methodological details of that work could affect the results, we consider it to be significant enough that IMFs obtained by different methods should not be compared in an attempt to search for variations in the IMF from region to region.  Truly reliable IMFs for embedded clusters will most likely require spectra for a large number of stars in the clusters; we are continuing to try to obtain spectra for these sources to better characterize the massive star population and the IMF of these clusters.  \\clearpage"
    },
    "0601/astro-ph0601560_arXiv.txt": {
        "abstract": "We analyze full-polarization VLBA data of ground-state, main-line OH masers in 18 massive star-forming regions previously presented in a companion paper.  From the aggregate properties of our sources, we confirm results previously seen in the few individual sources studied at milliarcsecond angular resolution.  The OH masers often arise in the shocked neutral gas surrounding ultracompact \\hii\\ regions. Magnetic fields as deduced from OH maser Zeeman splitting are highly ordered, both on the scale of a source as well as the maser clustering scale of $\\sim 10^{15}$~cm.  Results from our large sample show that this clustering scale appears to be universal to these masers.  OH masers around ultracompact \\hii\\ regions live $\\sim 10^4$ years and then turn off abruptly, rather than weakening gradually with time. These masers have a wide range of polarization properties.  At one extreme (e.g., W75~N), $\\pi$-components are detected and the polarization position angles of maser spots show some organization. At the other extreme (e.g., W51 e1/e2), almost no linear polarization is detected and total polarization fractions can be substantially less than unity.  A typical source has properties intermediate to these two extremes.  In contrast to the well ordered magnetic field inferred from Zeeman splitting, there is generally no clear pattern in the distribution of polarization position angles. This can be explained if Faraday rotation in a typical OH maser source is large on a maser amplification length but small on a single ($e$-folding) gain length.  Increasing or decreasing Faraday rotation by a factor of $\\sim 5$ among different sources can explain the observed variation in polarization properties.  Pure $\\pi$-components (in theory 100\\% linearly polarized) have long been sought, but seldom found.  We suggest that almost all $\\pi$-components acquire a signficant amount of circular polarization from low-gain stimulated emission of a $\\sigma$-component from OH appropriately shifted in velocity and lying along the propagation path. ",
        "introduction": "Hydroxyl (OH) masers are common in massive star-forming regions (SFRs).  Their small size and large Zeeman splitting coefficient allow them to serve as probes of the velocity and magnetic fields on a very small scale.  Because OH masers often cluster together in large numbers on subarcsecond scales \\citep[e.g.,][]{reidw3}, very long baseline interferometric (VLBI) techniques are required to identify individual maser features.  Due both to this close clustering of maser spots and to the tendency for the two components of a ground-state Zeeman pair to have highly unequal fluxes \\citep{cook}, the resolution afforded by VLBI is necessary in order to identify most Zeeman pairs \\citep[see, for instance,][]{fish03}.  Likewise, identifying individual maser spots in a cluster is a prerequisite to understanding the linear polarization of OH masers, since blending of the linear polarization from adjacent masers with different polarization properties can corrupt the interpretation of the polarization.  Over a quarter century has passed since the first OH maser source was observed with VLBI resolution: W3(OH) by \\citet{reidw3}. Since then, only a few more interstellar ground-state OH maser sources have been observed at VLBI resolution \\citep{haschick81,zheng97,slysh01b,slysh}.  In order to understand the range of environments probed by OH masers in massive SFRs, we have undertaken a survey of the OH masers in 18 massive SFRs with the National Radio Astronomy Observatory's\\footnote{The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc.} Very Long Baseline Array (VLBA).  The data have already been published as \\citet[hereafter Paper I]{paperi}.  In this paper we analyze the results both in terms of relevance to individual sources and as a large collection of OH masers that can shed additional light on the physical processes of OH masers and the interpretation of their polarization.  A brief overview of linear polarization theory is provided in \\S \\ref{lintheory}.  In \\S \\ref{results}, we consider the set of OH masers as a whole to derive statistical properties.  In \\S \\ref{discussion}, we present plausibility arguments to attempt to identify the physical processes responsible for the observed properties of OH masers.  Finally, we summarize our results in \\S \\ref{conclusions}.  Comments on individual sources are included in Appendix \\ref{sourcenotes}. ",
        "conclusions": "}  \\begin{itemize}  \\item Ground-state OH masers typically cluster on a scale of $10^{15}$~cm, providing evidence that their distribution is linked to a process with an inherent scale, as opposed to turbulence (which is generally scale-free).  The magnetic field strengths implied by Zeeman splitting suggest that OH masers occur in regions of density $10^5$ to several $\\times 10^7$~cm$^{-3}$.  OH masers are found preferentially near the UC\\hii\\ region in massive SFRs.  Their distribution around UC\\hii\\ regions suggest an expansion age of $\\approx 10^4$ years for typical expansion velocities.  OH masers do not appear to be systemically shifted from the velocity of the associated star by more than a few \\kms, although possible exceptions exist, as in G5.886$-$0.393 and W75~N VLA 2.  Taken together, these pieces of evidence support the theory that most OH masers occur in the shocked neutral gas between the ionization and shock fronts of UC\\hii\\ regions.  The distribution of maser fluxes with distance from the central UC\\hii\\ region suggests that OH masers turn off abruptly rather than weakening gradually after $\\sim 10^4$ years.  \\item Some OH masers are seen far from or without any associated \\hii\\ region.  It is unclear whether these masers are pumped by a star with an associated weak, undetected hypercompact \\hii\\ region or whether they are shock-excited without an ionization front.  In some sources (e.g., \\object[W75S]{W75~S}), OH masers appear to trace a collinear structure with a velocity gradient.  These formations probably trace shock fronts rather than protostellar disks.  \\item Magnetic fields are ordered in massive SFRs, lending observational support to theories that indicate that the ambient magnetic field direction may be preserved during massive star formation.  Nearly all sources show either a consistent line-of-sight magnetic field direction or a single reversal of the line-of-sight direction across the source.  Within a maser cluster of size $10^{15}$~cm, line-of-sight magnetic field direction reversals are never seen, and the field strengths deduced from Zeeman splitting are almost always consistent within $\\pm 1$~mG.  \\item We do see both $\\pi$- and $\\sigma$-components, including a ``Zeeman triplet'' in W75 N (see \\S \\ref{triplet}).  But OH maser spots that are 100\\% linearly polarized, as theoretically expected of $\\pi$-components, are extremely rare.  There is a range of sources with qualitatively different linear polarization properties.  At one extreme (as in W75~N) high linear polarization fractions are seen, and the PPAs show some correlation with observed structures and probable magnetic field directions.  In most sources the linear polarization fractions are much less than 1 and PPAs cannot be easily interpreted as magnetic field directions.  At the other extreme are sources such as W51 e1 and e2, in which little or no linear polarization is detected and the \\emph{total} polarization fraction of some maser spots is much less than unity.  \\item The wide range of polarization properties observed in OH masers may be explained by a combination of Faraday rotation and overlap of maser components.  If OH masers are indeed near or embedded in \\ion{C}{2} regions, the electron density may be high enough that the masers are near a critical point of Faraday rotation.  A typical maser spot likely has large ($> 1$~rad) Faraday rotation over the entire amplification length, but not over a single gain length of the maser. If Faraday rotation is a factor of $\\sim 5$ lower, the total Faraday rotation along the amplification path may be small enough such that the PPAs are still roughly aligned with the magnetic field lines.  On the other hand, if Faraday rotation is a factor of $\\sim 5$ larger, the Faraday rotation per gain length could exceed 1~rad, destroying linear polarization and depolarizing the maser.  Even if the Faraday rotation is small enough to allow amplification of a 100\\% linearly polarized $\\pi$-component, its polarization may be partially circularized by one of the $\\sigma$-modes of a weakly-inverted clump of OH between the maser site and the observer.  This is likely a very important effect, as only a modest inversion and a small column density of OH are required to add significant circular polarization to a $\\pi$-component.  \\item Theoretically the linear polarization fractions and directions of maser components can be used to determine the full, three-dimensional orientation of the magnetic field at masing sites.  But the interpretation of PPAs may be very difficult in sources for which the amount of Faraday rotation along the propagation path between the source and the observer is unknown.  Inferring a magnetic field orientation in the plane of the sky also requires unambiguous identification of $\\sigma$- and $\\pi$-components, due to the $90\\degr$ difference in PPA response to a magnetic field.  Zeeman pairs provide the surest method of identifying $\\sigma$-components, but their polarization properties may be too contaminated by internal Faraday rotation and maser overlap to permit interpretation of the inclination of the magnetic field to the line of sight.  \\end{itemize}"
    },
    "0601/astro-ph0601610_arXiv.txt": {
        "abstract": "We have attempted to establish an observational evidence for presence of distant companions which may have acquired and/or absorbed the angular momentum during evolution of multiple systems thus facilitating or enabling formation of contact binaries. In this preliminary investigation we use several techniques (some of them distance-independent) and mostly disregard detection biases of individual techniques in an attempt to establish a lower limit to the frequency of triple systems. While the whole sample of 151 contact binary stars brighter than $V_{max} = 10$ mag.\\ gives a firm lower limit of 42\\% $\\pm$ 5\\%, the corresponding number for the much better observed Northern-sky sub-sample is 59\\% $\\pm $8\\%. These estimates indicate that most contact binary stars exist in multiple systems. ",
        "introduction": "\\label{intro}  Formation of binary stars is a fascinating subject generating a very active research effort. Separate conferences \\citep{zinn2001} are devoted to this subject while specialized review papers \\citep{tholine2002,zinn2002,bate2004} describe difficulties in understanding of this complex process. Binaries with orbital periods of the order of a hundred of days are usually considered as close in this context. We have no idea how short-period binaries with periods much shorter than 3 to 5 days form. In fact, such binaries, particularly those with periods shorter than one day, should not exist: Indeed, even if some unknown process would form contact binaries at the T~Tauri stage, relatively large sizes of component stars would imply the resulting orbital periods to be longer than about 3 -- 5 days.  Formation in triple (or multiple) systems may alleviate the close-binary formation difficulty by allocating most of the angular momentum in the most distant component of a triple/multiple system, leaving a low angular momentum remnant. Such a scenario would have important observational consequences because -- by observing close binaries -- we would preferentially select stars with distant companions. One particular mechanism involving the Kozai cycles \\citep{koza1962} may produce close binaries by invoking a very strong tidal interaction during periastron passages when the inner, close (low angular momentum) orbit is perturbed by the third body into a very eccentric one. The close periastron interactions may then lead to a rapid angular momentum loss \\citep{kisel1998,EggKis2001}.  A confirmation of the triple body formation of very close binaries can only come through observations and careful statistical investigations. Unfortunately, proving or disapproving this hypothesis is difficult and -- as for all statistical studies -- is hindered by observational bias and selection effects. In this paper, we limit our scope to contact binaries, the most extreme objects with the lowest angular momentum content among all Main Sequence (MS) binary stars. Contact binaries \\citep{Rci1993} are made of solar-type stars with spectral types ranging from early K, through G and F to middle-A; they have short orbital periods ranging between 0.22 day to about 0.7 day at which point the statistics becomes unreliable and properties of massive contact binaries less well defined. Thus, in spite of the unquestioned existence of very massive contact binaries, it is not clear if a continuous sequence joins the solar-type contact binaries (stars also known as W~UMa type variables) with the early-type B and O-type binaries. In this paper, we limit ourselves to binaries with periods shorter than 1 day which is the traditional period limit for the solar-type (W~UMa-type) contact binaries.  \\citet{RK1982} commented on a high incidence of visual binaries among contact binaries while discussing a visual binary consisting of two W~UMa-type systems, BV~Dra and BW~Dra. Later on, \\citet{Chambliss1992} listed several cases of triple systems with contact binaries in his compilation and \\citet{HenryMki1998} detected a few cases of triple systems. However, these were fragmentary, almost anecdotal observations and there clearly exists a need of quantifying the matter of the multiplicity of very close binaries.  This paper attempts to characterize all indications of triplicity and/or multiplicity by categorizing contact systems into those with ``detected'', ``suspected'' and ``non-detected'' companions, with an additional obvious category of ``not observed'' ones. The techniques actually used are described in Sections~\\ref{direct} -- \\ref{indicators}. Not all techniques discussed by us are equally useful and not all imaginable techniques have been considered; there is still a lot of room for expansion. Also, types of companions are not explicitly discussed in the paper as our goal was simply to summarize the extant information irrespectively of what accompanies a given contact binary. Thus, a huge subject of selection effects and biases is only glanced over in this paper. Future studies, such as Paper~II of this series on spectroscopic detection in spectra used for the DDO radial velocity program \\citep{dangelo2005}, will more fully address specific limitations of various detection techniques.  We describe our sample in Section~\\ref{sample}. The main results of this paper are described in Sections~\\ref{direct} -- \\ref{indicators} and summarized in Table~\\ref{summary} where we collect all partial results and attempt to combine ``suspected'' cases into detections. Results for individual cases are summarized in a condensed form in the column FLAGS by codes specific for each technique. The codes are explained in Table~\\ref{flags} while individual binaries requiring comments are listed in Table~\\ref{notes}. Section~\\ref{conclude} summarizes the main results of this study while Section~\\ref{future} gives a discussion of possible directions for the future work. ",
        "conclusions": "\\label{future}  The approach that used is this paper was a straightforward one: We simply counted contact binaries which appeared to have companions. We totally disregarded the important matter of {\\it what are the companions and what would be the limitations in their detection\\/}. Some classes of stars, such as white dwarfs or neutron stars, would require utilization of special techniques; some classes may be entirely undetectable. We leave these issues open for future investigations of this series.  The use of several different observational techniques was certainly useful as none gives an unquestionable proof of multiplicity and each has its own biases and limitations. We used too many techniques to analyze selection effects of each; such in-depth studies of biases for each technique would inflate the study and would probably not be warranted at this preliminary stage. But we know that biases are important: Taking the astrometric data for the 50 best observed systems in the whole sky brighter than $V_{max}=10$ (Section~\\ref{direct}, Figure~\\ref{visual}), we see a frequency of 40\\% $\\pm$ 9\\%, but this number would rise to 46\\% $\\pm$ 10\\% if we disregarded the magnitude limit and argued that objects best observed astrometrically would provide an even more reliable sample (indeed, there are many bright systems with poor data). The sky location is also important: for the Northern sub-sample of astrometrically well observed stars to $V_{max}=10$, the frequency increases from 40\\% $\\pm$ 9\\% to 50\\% $\\pm$ 10\\%. Similarly, entirely independent indications for the 20 best observed systems for the LITE suggest 60\\% $\\pm$ 17\\%; among those LITE solutions, among these only two are for nominally Southern systems, although both were observed from the North (ER~Ori with $\\delta = -8.5^\\circ$ and UX~Eri with $\\delta = -6.9^\\circ$).  The most obvious observational bias caused by the brightness of the system is demonstrated in Fig.~\\ref{multi}, where the fraction of detections is plotted for all systems brighter than the given magnitude rank. While for systems brighter than $V_{max}$ = 9 mag we see an almost constant fraction of multiples, there is a small but definite decrease of detections for fainter systems.  In addition to in-depth analyzes of selection effects and biases of individual techniques, we issue the usual plea to the observers: More data are needed! What is particularly needed are photometric data in standard systems. If not for the Hipparcos mission, which led to many discoveries of contact binaries and if not for the photometric part of the TYCHO project, many contact binaries would be unknown or without most elementary data on color indices. All contact systems require also radial velocity spectroscopic support. The situation with the Southern sky is particularly grave. There exists no spectroscopic survey similar to that being conducted now for short-period Northern binaries at DDO so that even relatively bright southern contact binaries (TY~Men, DE~Oct, MW~Pav, CN~Hyi, WY~Hor, V386~Pav, etc.) still await spectroscopy or even conventional precision photometry. Very few Southern binaries are followed over time to provide LITE information. For many binaries, the last datum on the moment of minimum is the Hipparcos $T_0$\\footnote{This remark unfortunately applies also to some Northern binaries which do not happen to be popular...}.  Interesting Southern systems omitted by Hipparcos (e.g., BR~Mus, TV~Mus or ST~Ind) frequently have the last minima obtained a quarter or half a century ago. We point also that the technique of lunar occultations still remains to be fully utilized as it can provide angular resolution of about 2 mas and thus can be a powerful tool for the detection of close visual companions.  We conclude with a summary of our results: Our best estimate of the lower limit to the triple star incidence among contact binaries appears to be 59 $\\pm$ 8\\%, as based on the Northern sky to $V_{max}=10$. We know that this limit may raise as suspected cases are analyzed more thoroughly. If, in addition to unquestionable detection of 64 multiple systems among 151 targets over the whole sky, all suspected 24 cases in Table~\\ref{summary} were included as detections, then the frequency for the whole sky would increase from 42\\% $\\pm$ 5\\% to 56\\% $\\pm$ 6\\%, while that for the Northern hemisphere alone it would reach 72\\% $\\pm$ 9\\%. Our results are consistent with those for the magnitude-complete Hipparcos sample to $V_{max}=7.5$ \\citep{Rci2002b}, with 17 multiples among 35 systems ($48 \\pm 12$\\%), although for that sample a large fraction of all systems were newly discovered eclipsing binaries with currently incomplete data. Remembering that we could evaluate only a lower limit to the frequency of triple systems, and that some multiple systems are beyond reach of any current technique of detection, this number is not far from one indicating a possibility that all contact binaries originated in multiple systems."
    },
    "0601/astro-ph0601495_arXiv.txt": {
        "abstract": "Using the Gemini South 8m telescope, we obtained high resolution 11.7 and 18.3~\\mic\\ mid-IR images of SN 1987A on day 6526 since the explosion. All the emission arises from the equatorial ring. Nearly contemporaneous spectra obtained at 5--38~\\mic\\ with the {\\sl Spitzer} Space Telescope show that this is thermal emission from silicate dust that condensed out in the red giant wind of the progenitor star. The dust temperature is $166^{+18}_{-12}$~K, and the emitting dust mass is $(2.6^{+2.0}_{-1.4}) \\times 10^{-6}$ \\Msun. Comparison of the Gemini 11.7~\\mic\\ image with Chandra X-ray images, Hubble UV- optical images, and ATCA radio synchrotron images shows generally good correlation across all wavelengths. If the dust resides in the diffuse X-ray emitting gas then it is collisionally heated. The IR emission can then be used to derive the plasma temperature and density, which were found to be in good agreement with those inferred from the X-rays.  Alternatively, the dust could reside in the dense UV-optical knots and be heated by the radiative shocks that are propagating through the knots. In either case the dust-to-gas mass ratio in the CSM around the supernova is significantly lower than that in the general interstellar medium of the LMC, suggesting either a low condensation efficiency in the wind of the progenitor star, or the efficient destruction of the dust by the SN blast wave. Overall, we are witnessing the interaction of the SN blast wave with its surrounding medium, creating an environment that is rapidly evolving at all wavelengths. ",
        "introduction": "Since its explosion, SN~1987A has evolved from a supernova (SN) dominated by the emission from the radioactive decay of $^{56}$Co, $^{57}$Co and $^{44}$Ti  in the ejecta to a supernova remnant whose emission is dominated by the interaction of the supernova blast wave with its surrounding medium. The medium surrounding the SN is dominated by the well-known ``circumstellar envelope\" (CSE), which consists of an inner equatorial ring (ER) flanked by two outer rings \\citep{Bur95}, possibly part of an hour-glass structure.  The collision between the ejecta of SN~1987A and the ER predicted to occur sometime in the interval 1995-2007 \\citep{Gae97,Bor97} is now underway. At UV-optical (UVO) wavelengths, ``hot spots\" have appeared inside the ER \\citep{Pun97}, and their brightness varies on time scales of a few months \\citep{Law00}. New hot spots continue to appear as the whole inner rim of the ER lights up. The visible-light {\\it HST} image obtained in 2004 Feb 20 reveals a necklace of such hot spots which nearly fill a lighted ring. Ongoing monitoring at X-rays with the {\\it Chandra} and at radio frequencies, shows that the evolution of the emission from the ER follows a similar pattern at all wavelengths.  There exist very few mid-infrared (IR) observations of supernovae in general. Therefore SN~1987A, the closest known supernova in 400 years, gives us an opportunity to explore the mid-IR properties of supernovae and the dust in their ejecta and surrounding medium with the help of the newest generation of large-aperture telescopes and sensitive mid-IR instrumentation like the T-ReCS on the {\\it Gemini},  in combination with IR data obtained from {\\it Spitzer Space Telescope}({\\it SST}) \\citep{Wer04}. The T-ReCS observations of the mid-IR emission from SNR~1987A are part of our continuous monitoring of the SN and its surrounding medium. The first detection and analysis of mid-IR emission at the position of the supernova has been reported in \\citep{Bou04} (hereafter Paper I).  The origin of the mid-IR emission could be line emission from atomic species, synchrotron or free-free continua, or thermal emission from dust which is probably the dominant source of emission. In general, there are several scenarios for the origin and the heating mechanism of the dust giving rise to the late time mid-IR emission in Type II supernovae \\citep{Gra86,Ger00}. Thermal mid-IR emission could be: (1) the emission from SN-condensed dust that is collisionally heated by reverse shocks traveling through the SN ejecta;  (2) the emission from circumstellar/interstellar dust heated by the interaction of the expanding SN blast wave with the ambient medium; and (3) the delayed emission (echo) from circumstellar dust  radiatively heated by the early UVO supernova light curve.  Our imaging observations provide strong constraints on these possible scenarios for the mid-IR emission. They show that the bulk of the mid-IR emission is not concentrated on the center of the explosion, but arises from a ring around the SN. The morphology of the mid-IR emission can therefore be used to eliminate several scenarios for its origin. Combined with {\\it Spitzer} spectroscopy, our observations can be used to determine the dust composition and temperature distribution in the ring. The  {\\it Chandra} X-ray, and {\\it HST} UVO data provide important constraints on the physical conditions of the medium overtaken by the supernova blast wave. These constraints can be used to determine the physical association and the heating mechanism of the dust giving rise to the mid-IR emission.  The paper is organized as follows: We first describe in \\S2 the imaging observations of the SN obtained by T-ReCS, and compare its IR morphology to that at radio, X-ray, and UVO wavelengths. Lower resolution {\\it Spitzer} mid-IR imaging observations detected SNR~1987A (including the ring) as an unresolved point source, and are used in conjunction with spectroscopic data to determine the possible contribution of lines, and the composition of the dust giving rise to the continuum emission. In \\S3 we describe the procedure used for the analysis of the spectroscopic and imaging data, determining the dust composition, and presenting maps of dust temperature, IR opacities, and dust column densities. The IR image of the circumstellar medium around SN~1987A has a morphology similar to the {\\it Chandra} X-ray and {\\it HST} UVO images. The limited mid-IR resolution does not allow us to unambiguously  determine whether the dust resides in the X-ray emitting gas, or in the UV-optical line emitting knots in the ER. We therefore resort in \\S4 to an analysis of possible dust heating mechanisms and to calculating the inferred dust masses and dust-to-gas mass ratios for several possible scenarios. In \\S5 we discuss the evolution of the supernova and its environment as manifested from the observed light curves at various wavelengths.  The results of our paper are summarized in \\S6. ",
        "conclusions": "We have presented mid-IR images of SNR~1987A obtained with T-ReCS on the {\\it Gemini} South telescope on day 6526 at 11.7 and 18.3 \\mic\\ and with IRAC (5.8 and 8~\\mic) on day 6130 and MIPS (24~\\mic) on day 6187 onboard {\\it Spitzer}, together with 3 - 37~\\mic\\ spectroscopic observations of the remnant obtained with IRS on day 6190 at the same observatory. The imaging observations (Figure~\\ref{fig1}) show that the mid-IR emission arises from the  dust in the equatorial ring (ER) heated up by the interaction of the SN blast wave with its circumstellar medium. Several theoretical models predicted the presence of dust in the CSE of SN~1987A which was produced in the winds of the supergiant phase. The location of the IR emission rules out the possibility that the dust condensed out in the SN ejecta, strongly suggesting a circumstellar origin instead.  The 3 - 37~$\\micron$ spectrum (Figures~\\ref{fig10}) shows that the emission arises from a population of astronomical silicate particles. Temperature maps show that the dust temperature is fairly uniform in the ER and about 166$^{+18}_{-12}$~K, with total dust masses of $\\sim 2.6^{+2.0}_{-1.4} \\times 10^{-6}$~\\Msun.  A comparison with {\\it Chandra}  and {\\it HST} observations show an equally good correlation between the 11.7~\\mic\\ IR and the X-ray and  UV-optical images of the supernova. Because of the limited angular resolution in the IR we cannot determine the location or heating mechanism of the radiating dust.  The dust could be residing in the hot $\\sim 10^7$~K gas and collisionally heated by the X-ray emitting plasma. The dust temperature is then an excellent diagnostic of the electron density, giving a value of $\\sim $ 300-1400~cm$^{-3}$, similar to the value suggested by the {\\it Chandra} observations.  Comparison of the IR and X-ray fluxes suggests that the dust is depleted by a factor of $\\sim$ 30 in the X-ray emitting gas,  compared to its value in the local interstellar medium of the LMC. This low value could be due to its destruction by thermal sputtering in the shocked gas, requiring the initial grain radii to be below $\\sim$ 50~\\AA.  Alternatively, the dust could be residing in the UV-optical emitting knots in the ER and radiatively heated by the cooling gas that was excited by the shocks propagating through these knots. A simple calculation for a particular knot shows that the radiative energy density can heat the dust to typical temperatures of about 125~K, similar to those inferred from the IR observations. A comparison with the mass of the knots shows that the dust-to-gas mass ratio in the knots is lower by a factor of $\\sim$ 10 compared to  its value in the ISM of the LMC. This low abundance reflects the low condensation efficiency of the dust in the outflow of the progenitor star.  We stress that in order to assess the role of SNe in the production of dust in the Universe, it is clearly important to measure the presence of dust that survives into the formation of the remnant, and for this, mid-IR and sub-mm observations are critical."
    },
    "0601/astro-ph0601176_arXiv.txt": {
        "abstract": "As part of our radial velocity survey of low Galactic latitude structures that surround the Galactic disc, we report the detection of the so called Monoceros Ring in the foreground of the Carina dwarf galaxy at Galactic coordinates $(l,b)=(260\\deg,-22\\deg)$ based on VLT/FLAMES observations of the dwarf galaxy. At this location, 20 degrees in longitude greater than previous detections, the Ring has a mean radial velocity of $145\\pm5\\kms$ and a velocity dispersion of only $17\\pm5\\kms$. Based on Keck/DEIMOS observations, we also determine that the Ring has a mean radial velocity of $-75\\pm4\\kms$ in the foreground of the Andromeda galaxy at $(l,b)\\sim(122\\deg,-22\\deg)$, along with a velocity dispersion of $26\\pm3\\kms$. These two kinematic detections are both highly compatible with known characteristics of the structure and, along with previous detections provide radial velocity values of the Ring over the $120\\deg<l<260\\deg$ range. This should add strong constraints on numerical models of the accretion of the dwarf galaxy that is believed to be the progenitor of the Ring. ",
        "introduction": "With the public release of all sky surveys, many details have been gained in the structure of the outer parts of the Galactic disc. In particular, the Sloan Digital Sky Survey (SDSS) revealed the existence of a stellar structure in the anticentre direction, near the Galactic plane and slightly over the edge of the disc \\citep{newberg02}. Visible as a clear main sequence at a Galactocentric distance of $18\\kpc$ for $180\\deg<l<225\\deg$ and $|b|<30\\deg$, this structure is unlike what is expected for the Galactic disc. \\citet{ibata03} used the INT Wide Field Camera to show this structure is in fact present in the second and third Galactic quadrants, circling the disc in a ring-like fashion. Radial velocity measurements also revealed a kinematically cold population with a velocity dispersion of only $15-25\\kms$, once more unexpected for a Galactic structure, leading to the conclusion that this so-called Ring\\footnote{This structure has been called many names including the Monoceros Ring, the Galactic Anticentre Stellar Structure or GASS and the Ring. Given its extent in Galactic longitude and its presence in numerous constellations, we prefer to call it simply the Ring.} is produced by the accretion of a dwarf galaxy in the Galactic plane whose tidal arms are wrapped around the Milky Way.  Subsequently, much work has been invested in trying to determine the true extent of this Ring and where possible determine its kinematics to constrain the orbit of its progenitor. \\citet{rocha-pinto03} used the 2MASS catalogue to probe the distribution of M giant stars and found that the Ring may extend over the $120\\deg\\simlt l\\simlt270\\deg$ range in the Northern hemisphere and be present in the Southern hemisphere, with $|b|<35\\deg$. \\citet{conn05a} extended the \\citet{ibata03} INT/WFC survey to show the Ring is present within 30 degrees of the Galactic plane throughout the whole second quadrant but seems to disappear at $l\\sim90\\deg$. More constraints on the structure were gained by the \\citet{crane03} spectroscopic survey of putative Ring M-giant stars in the anticentre direction ($150\\deg<l<230\\deg$, $|b|<40\\deg$) which confirmed that most of these stars are indeed linked to this structure. Using the simple model of a population orbiting the Milky Way in a prograde, circular orbit, they determined their sample was best fitted by a population at a Galactocentric distance of $18\\kpc$ and with a rotational velocity of $220\\kms$. The resulting velocity dispersion of $20\\pm4\\kms$ around this model is in good agreement with the velocity dispersion measured by the SDSS team \\citep{yanny03}.  While tracking the Ring, two other structures were discovered that do not seem to be directly related to the Monoceros Ring. Using the 2MASS catalogue, \\citet{rocha-pinto04} reported the existence of a diffuse population of M giants in the direction of the Triangulum and Andromeda constellations, behind known detections of the Ring. Also using the 2MASS catalogue, our group presented evidence of a dwarf galaxy located closer than the Ring, just at the edge of the Galactic disc, in the Canis Major constellation \\citep{martin04,bellazzini04}. In particular, we argued that the accretion of the Canis Major galaxy onto the Galactic disc would naturally reproduce similar features as the one observed for the Ring. However, a link between these three structures remains putative at the moment, even though current models show such a scenario is highly plausible \\citep{penarrubia05,martin05,dinescu05}.  To gain more insight into the nature of these structures, we have started a radial velocity survey of regions at low Galactic latitude where the Ring structure may be located. Using the AAT/2dF multi-fibre spectrograph, we first targeted the Canis Major dwarf \\citep{martin05}. This survey also revealed the presence of the Ring behind the dwarf, with a Heliocentric radial velocity of $133\\pm1\\kms$ and dispersion of $23\\pm2\\kms$ at a Galactocentric distance of $19\\kpc$, highly compatible with previous detections \\citep{conn05b}. In this third paper of the series, we report the presence of the Ring in the foreground of the Carina dwarf galaxy from VLT/FLAMES observations of the dwarf at $(l,b)=(260\\deg,-22\\deg)$. Using Keck2/DEIMOS observations of regions around M31, we also determine the radial velocity of the Ring at $(l,b)\\sim(122\\deg,-22\\deg)$ where the Ring is known to exist but where its radial velocity has not yet been measured. Section 2 presents the Ring detection in front of Carina and Section 3 deals with the detection in front of the Andromeda galaxy. Section 4 concludes this letter.  In the following, all the magnitudes have been corrected for extinction using the maps from \\citet{schlegel98}. We also assume that the Solar radius is $R_\\odot = 8\\kpc$, that the LSR circular velocity is $220\\kms$, and that the peculiar motion of the Sun is ($U_0=10.00\\kms, V_0=5.25\\kms, W_0=7.17\\kms$; \\citealt{dehnen98}). Except when stated otherwise, the radial velocities, $v_{r}$, are Heliocentric radial velocities, not corrected for the motion of the Sun. ",
        "conclusions": "\\begin{figure} \\ifthenelse{\\UseFigs=1}{ \\includegraphics[width=\\hsize]{fig5.ps} }{caption1} \\caption{Summary of all known radial velocity detections of the Ring in the Galactic standard of rest. The dashed line corresponds to the \\citet{crane03} model of the Ring: a population orbiting the Milky Way at a Galactocentric distance of $18\\kpc$ and with a rotational velocity of $220\\kms$. The \\citet{yanny03} SDSS detections are shown as stars, the \\citet{conn05b} detection behind Canis Major is shown by a hollow circle and the two detections of this letter in front of Carina and M31 are shown as a filled triangle and a filled square respectively. Uncertainties on these values are not shown since they are smaller than the used symbols. Though near the \\citet{crane03} model, detections away from the anticentre tend to be offset from the model.} \\end{figure}  We have presented the detection of two groups of stars that lie in front of the Carina dwarf galaxy at $(l,b)=(260\\deg,-22\\deg)$ and in front of the Andromeda galaxy at $(l,b)\\sim(122\\deg,-22\\deg)$ and that cannot be satisfyingly explained by known Galactic components. The proximity with known detections of the Ring that surrounds the Galactic disc makes it highly probable that they belong to the same structure. Both detections have a low velocity dispersion ($17\\pm5\\kms$ and $26\\pm3\\kms$ respectively), a characteristic value encountered in all previous detections of the Ring.  With the Ring detection behind the Canis Major dwarf galaxy reported by \\citet{conn05b}, the radial velocity of the Ring population is now sampled throughout the $120\\deg<l<260\\deg$ range, which should provide important constraints on N-body models. However, it can be directly seen in Figure~5 that currently known radial velocity values for the Ring are not exactly reproducible by the \\citet{crane03} simple circular model. In fact, trying to fit all the detections in a single orbit of the progenitor proves unsatisfactory, whether the orbit is forced to be circular or allowed to be slightly elliptical. It would therefore seem that models where the Ring completely surrounds the Galactic disc with multiple tidal arms are following the right track (see e.g. \\citealt{penarrubia05} and \\citealt{martin05}). In addition to the new radial velocities we report in this letter, such models would highly benefit from a similar survey as the one presented in \\citet{conn05a}, but this time to higher longitudes to study in more detail the morphology of the Ring in these regions and especially to add a distance constraint to the detection in front of the Carina dwarf."
    },
    "0601/astro-ph0601389_arXiv.txt": {
        "abstract": "We study the observational constraints of CMB temperature and polarization anisotropies on models of dark energy, with special focus on models with variation in properties of dark energy with time. We demonstrate that the key constraint from CMB observations arises from the location of acoustic peaks. An additional constraint arises from the limits on $\\Omega_{NR}$ from the relative amplitudes of acoustic peaks. Further, we show that the distance to the last scattering surface is not how the CMB observations constrain the combination of parameters for models of dark energy. We also use constraints from Supernova observations and show that unlike the Gold and Silver samples, the SNLS sample prefers a region of parameter space that has a significant overlap with the region preferred by the CMB observations. This is a verification of a conjecture made by us in an earlier work \\citep{2005PhRvD..72j3503J}. We discuss combined constraints from WMAP5 and SNLS observations. We find that models with $w \\simeq -1$ are preferred for models with a constant equation of state parameters. In case of models with a time varying dark energy, we show that constraints on evolution of dark energy density are almost independent of the type of variation assumed for the equation of state parameter. This makes it easy to get approximate constraints from CMB observations on arbitrary models of dark energy. Constraints on models with a time varying dark energy are predominantly due to CMB observations, with Supernova constraints playing only a marginal role. ",
        "introduction": "Observational evidence for accelerated expansion in the universe has been growing in the last two decades \\citep{crisis2, crisis1, crisis3}. Independent confirmation using observations of high redshift supernovae \\citep{1998ApJ...509...74G, 1999ApJ...517..565P, nova_data1, nova_data2, nova_data3, snls} has made this result more acceptable to the community. Using these observations along with observations of cosmic microwave background radiation (CMB) \\citep{boomerang, wmap_params, wmap5b} and large scale structure \\citep{sdss}, we can construct a ``concordance'' model for cosmology and study variations around it (e.g., see \\citet{wmap5a, wmap5b, 2003Sci...299.1532B, 2004ApJ...606..702T}; for an overview of our current understanding, see \\citet{TPabhay, tptalk, tpessay}).  Observations indicate that the dominant component of energy density --- called dark energy --- should have an equation of state parameter $w\\equiv P/\\rho < -1/3$ for the universe to undergo accelerated expansion. Indeed, present day observations require $w \\simeq -1$. The cosmological constant is the simplest explanation for accelerated expansion \\citep{ccprob,1992ARA&A..30..499C,review3,review1,review2,review4,review5,peri_rev,sami_rev} and it is known to be consistent with observations. In order to avoid theoretical problems related to cosmological constant \\citep{ccprob,1992ARA&A..30..499C}, many other scenarios have been investigated: these include quintessence \\citep{quint1,quint2,quint3,quint4,quint5,quint6,quint7,quint8}, k-essence \\citep{2001PhRvD..63j3510A,k-essence2,k-essence3,k-essence4,k-essence5}, tachyon  field \\citep{tachyon5,2002PhRvD..66b1301P,tachyon2,tachyon3,tachyon4,tachyon6,tachyon7,tachyon8,tachyon9,tachyon10,tachyon11,tachyon12}, chaplygin gas and its generalisations \\citep{gas1,gas2,gas3,gas4}, phantom fields \\citep{phantom1,phantom2,phantom3,phantom4,phantom4a,phantom5,phantom6,phantom7, phantom9,phantom10,2004PhRvD..70d3539E,2005PhRvD..71f3004N,2007PhLB..646..105B,2007JPhA...40.6835B,2007JETPL..85....1B}, branes \\citep{brane1,brane2,brane3,brane4,brane5,brane6} and many others \\citep{water,horz1,horz2,horz3,horz4,horz5,gaussb,viscous,dynamic_new1,dynamic_new2,dynamic_new3,2006PhRvD..74f4021A,2007JHEP...09..048A,arman} In these models one can have $w \\ne -1$ and in general $w$ varies with redshift. For references to papers that discuss specific models, the reader may consult one of the many reviews \\citep{review3,review1,review2,review4,review5,alcaniz_rev}. Even though these models have been proposed to overcome the fine tuning problem for cosmological constant, most models require similar fine tuning of parameter(s) to be consistent with observations. Nevertheless, they raise the possibility of $w(z)$ evolving with time (or it being different from $-1$), which can be tested by observations.  Given that $w$ for dark energy should be smaller than $-1/3$ for the Universe to undergo accelerated expansion, the energy density of this component changes at a much slower rate than that of matter and radiation. (Indeed, $w=-1$ for cosmological constant and in this case the energy density is a constant.) Unless $w$ is a rapidly varying function of redshift and becomes $w \\sim 0$ at ($z \\sim 1$), the energy density of dark energy should be negligible at high redshifts ($z \\gg 1$) compared to that of non-relativistic matter. If dark energy evolves in a manner such that its energy density is comparable to, or greater than the matter density in the universe at high redshifts then the basic structure of the cosmological model needs to be modified. We do not consider such models, instead we confine our attention to constraints on dark energy in realistic models and choose observations which are sensitive to evolution of $w(z)$ at redshifts $z \\leq 1$.  Supernova observations permit a large variation in the equation of state \\citep{dynamic_de4,dynamic_de4a}. It has recently been argued that a cosmological constant with a small curvature can be interpreted as a dynamical dark energy model \\citep{2007JCAP...08..011C,virey_etal}. We have shown that a combination of supernova observations with CMB observations and  abundance of rich clusters of galaxies provides tight constraints on variation of dark energy  \\citep{jbp,2005PhRvD..72j3503J}. Of these CMB data provides the most stringent bounds on the allowed variation in evolution of dark energy density. In this paper the main motivation is to study how these models fare in the light of current CMB data \\citep{wmap5a,wmap5b}.  A variety of observations can be used to constrain models of dark energy, e.g. see \\S{II~B} of \\cite{2005PhRvD..72j3503J} for an overview. Observations of high redshift supernovae provided the first direct evidence for accelerated expansion of the universe \\citep{1998AJ....116.1009R,1999ApJ...517..565P}. This, coupled with the ease with which the high redshift supernova data can be compared with cosmological models has made it the favourite benchmark for comparison with models of dark energy. It is often considered sufficient to compare a model with the supernova data even though observers and theorists have pointed out potential problems with the data \\citep{astph0303428,astph0506478} as well as some peculiar implications of the data \\citep{dynamic_de1,dynamic_de2,shashikant}. Further, it has been shown that other observations like temperature anisotropies in the CMB fix the distance to the last scattering surface and are a reliable probe of dark energy \\citep{white,jbp}. In this work, we try to understand the origin of the constraint on models of dark energy from CMB anisotropy observations. We show that the angular scale of acoustic peaks provides the leading constraint on the combination of parameters. An additional constraint comes in through the degeneracy with $\\Omega_{NR}$, where a constraint on $\\Omega_{NR}$ from the relative heights of acoustic peaks translates into an indirect constraint on the equation of state parameter $w$.  In an earlier work, we compared the constraints on models of dark energy from supernova and CMB observations and pointed out that models preferred by these observations lie in distinct parts of the parameter space and there is no overlap of regions allowed at $68\\%$ confidence level \\citep{2005PhRvD..72j3503J}. Even though different observational sets are  sensitive to different combinations of cosmological parameters, we do not expect models favoured by one observation to be ruled out by another when such a divergence is not expected. This divergence may point to some shortcomings in the model, or to systematic errors in observations, or even to an incorrect choice of priors. We suggested that this may indicate unresolved systematic errors in one of the observations, with supernova observations being more likely to suffer from this problem due to the heterogeneous nature of the data sets available at the time. In Supernova Legacy Survey (SNLS) \\citep{snls} survey, a concerted effort has been made to reduce systematic errors by using only high quality observations. The systematic uncertainties are reduced by using a single instrument to observe the fields. Using a rolling search technique ensures that sources are not lost and data is of superior quality (for details see \\cite{snls}). If our claim about Gold+Silver data set were to be true, SNLS data should not be at variance with the WMAP data\\footnote{This has been shown by several authors, including ourselves in a much earlier version of this manuscript.} In this work we study constraints on dark energy models from high redshift supernova observations from the SNLS survey and also observations of the temperature and polarisation anisotropies in the CMB using the WMAP5 data.  This paper is a revised version of an earlier manuscript which became out of date after WMAP-3 and then WMAP-5 data were released.  In the interim period, the issue of systematics in inhomogeneous data sets has been accepted, therefore we do not emphasize that aspect much in this version.  The focus of the current paper is twofold: to study constraints on dark energy models in light of the SNLS and WMAP-5 data, and, to understand the combination of cosmological parameters that is constrained by the CMB observations. ",
        "conclusions": "In this work we studied the SNLS data set and compared the constraints obtained with constraints from WMAP five year data on temperature anisotropies in the CMB. We find that the parameter values favoured by the two data sets have significant overlap and the two sets can be combined to put tight constraints on models of dark energy. In an earlier work we had noted that the Gold+Silver data set does not agree with WMAP observations in that these favour distinct parts of the parameter space \\citep{2005PhRvD..72j3503J}. Constraints from WMAP and structure formation favour similar models, but ones distinct from those favoured by Gold+Silver supernova observations. This indicates some degree of inconsistency between the supernova and other observations and it led us to suggest that the Gold+Silver data set may be affected by as yet unknown systematic errors \\citep{2005PhRvD..72j3503J}. One reason for doubting the supernova data is the heterogeneity of sources from which the particular data set was collected \\citep{nova_data3}. SNLS \\citep{snls} is a homogeneous data set and should not suffer from such problems and indeed we find that there is no inconsistency between SNLS and WMAP observations.  This highlights the usefulness of CMB observations for constraining models of dark energy \\citep{white,jbp}. We believe that CMB observations should be used for testing any model of dark energy as supernova observations do not constrain models with varying $w$ effectively \\footnote{The main constraint on varying $w$ models is from the CMB data}. Thus one should use CMB observations as well and not rely only on supernova observations for constraining such models.  We would like to add a note of caution against combining the SNLS data with other data sets of high redshift supernovae in light of the very different nature of these data sets. Indeed, one should use homogeneous data sets like the SNLS in isolation  to avoid the problems mentioned above.  In terms of models, we find that the cosmological constant is favoured by individual observations (SNLS and WMAP) as well as in the combined data set with very high probability. Table~2 gives allowed values of all cosmological parameters at $95\\%$ confidence level by SNLS, WMAP as well as the combined data set. For the cases where a similar analysis has been done by others, our results are consistent with other findings \\citep{2003Sci...299.1532B,wmap_params,constraints_5,constraints_7,constraints_8,2004ApJ...606..702T,constraints_2,constraints_13,constraints_14,constraints_3,constraints_10,constraints_9,constraints_10a,constraints_11,constraints_12a,constraints_12,peri_snls,2005PhRvD..71j3515S,constraints-12,feng1,feng2,feng3, 2007PhRvD..75h3506A,2006PhRvD..74h3519C,2006PhRvD..74b3532C,2007A&A...467..421D,2008arXiv0803.1311E,mota_constr}.  We have discussed the origin of the constraint on dark energy models from CMB observations at length. We may conclude from the analysis presented here that: \\begin{itemize} \\item Location of acoustic peaks in the angular power spectrum of CMB anisotropies is the main source of constraints. \\item CMB observations only constrain $w_{eff}$, an effective value of the equation of state parameter defined in Eqn.(\\ref{eqn:weff}). This can be used to translate constraints on models with constant $w$ to models where dark energy properties vary with time. \\item We have discussed models with $\\Omega_0 = 1$ in this paper. In case this constraint is relaxed then the well known degeneracy between $w$ and $\\Omega_0$ loosens the constraints. However, it is well known that the SN and CMB data are complementary and can be combined to provide fairly tight constraints even in this case \\citep{lss_de,2001PhRvD..64l3527H}. We do not expect variations in curvature to modify our conclusions about $w_{eff}$ being the only dark energy related quantity constrained by CMB observations. \\item Integrate Sachs-Wolfe (ISW) effect due to perturbations in dark energy can, in principle, lead to variations in the CMBR angular power slectrum at small $l$. Our analysis of models with variable $w$ without perturbations does not take this into account. Various analyses have shown that ISW does not contribute significantly to constraining cosmological parameters from CMB data (see for example, \\citep{2009JCAP...09..003X}).  The main reason for this is that ISW affects power at small $l$, and due to large cosmic variance these modes do not contribute much to the overall likelihood. \\end{itemize}"
    },
    "0601/astro-ph0601340_arXiv.txt": {
        "abstract": "We solve exactly the problem of a finite slab receiving an isotropic radiation on one side and no radiation on the other side. This problem - to be more precise the calculation of the source function within the slab - was first formulated by K. Schwarzschild in 1914. We first solve it for unspecified albedos and optical thicknesses of the atmosphere, in particular for an albedo very close to 1 and a very large optical thickness in view of some astrophysical applications. Then we focus on the conservative case (albedo = 1), which is of great interest for the modeling of grey atmospheres in radiative equilibrium. Ten-figure tables of the conservative source function are given. From the analytical expression of this function, we deduce 1) a simple relation between the effective temperature of a grey atmosphere in radiative equilibrium and the temperature of the black body that irradiates it, 2) the temperature at any point of the atmosphere when it is in local thermodynamical equilibrium. This temperature distribution is the counterpart, for a finite slab, of Hopf's distribution in a half-space. Its graphical representation is given for various optical thicknesses of the atmosphere. ",
        "introduction": "\\label{sec1} In a paper considered as foundational in transfer theory, K. Schwarzschild \\cite{schwarzschild1914} modeled the solar atmosphere as a finite slab irradiated on one side by black body radiation. He set up, for the first time, the complete radiation transfer equation - including the scattering integral - and inferred the following integral equation for the source function $S$: \\begin{equation} \\label{eq1} S(\\tau) = S_0 (\\tau) + \\frac{a}{2} \\int_{0}^{b} E_{1} (| \\tau - \\tau\\p |) S(\\tau\\p) \\d\\tau\\p, \\end{equation} with \\begin{equation} \\label{eq2} S_0(\\tau) = (1-a)B^*(\\tau) + \\frac{a}{2} B(T_b)E_2(b-\\tau). \\end{equation} The first term on the right-hand side of Eq. (\\ref{eq1}) describes the contribution of the (internal and external) sources, and the integral term is the scattering term. $a \\in [0,1]$ is the albedo (supposed constant) for single scattering by a volume element, $b>0$ the optical thickness of the atmosphere, and $\\tau \\in [0, b]$ is the optical depth variable. Eq. (2) shows that the source term $S_0$ includes the thermal emission $B^*$ of the atmosphere, as well as the contribution of the external sources. In our model, they will be prescribed as black body radiation of temperature $T_b$, incident on the surface $\\tau = b$. $B(T_b)$ is the Planck function at the temperature $T_b$. There is no radiation incident on the other boundary surface $\\tau=0$. The scattering of light is supposed monochromatic and isotropic, which assumption gives rise to the exponential integral functions \\begin{equation} \\label{eq3} E_{n} (\\tau) = \\int_{0}^{1} \\exp (-\\tau / u) u^{n-2}\\d u \\qquad (n \\geq 1) \\end{equation} in the kernel and the free term of Eq. (\\ref{eq1}). Frequency dependence of quantities will not be made explicit.\\\\ The integral form of the radiation transfer equation was established before the publication of Schwarzschild's paper, by Lommel \\cite{lommel1889}, Chwolson \\cite{chwolson1890} and King \\cite{king1913}. In these articles, Eq. (\\ref{eq1}) with $S_0(\\tau) = \\exp(-\\tau/\\mu)$ ($\\mu>0$) was derived directly, with no reference to the transfer equation. Integral equations of the form (\\ref{eq1}) have been studied intensively since these pioneering works. Three general methods for solving them, which apply to the particular source term (\\ref{eq2}), are described in \\cite{rutily-bergeat1994}. Since the problem (\\ref{eq1})-(\\ref{eq2}) is linear, its solution can be written as \\begin{equation} \\label{eq4} S(\\tau) = (1-a)\\int_{0}^{b} R(a,b,\\tau, \\tau\\p)B^*(\\tau\\p) \\d \\tau\\p + B(T_b)\\xi_{0}(a, b,b-\\tau), \\end{equation} where $R$ denotes the resolvent kernel of the integral equation (\\ref{eq1}) and $\\xi_0$ the solution to the following integral equation: \\begin{equation} \\label{eq5} \\xi_0(a,b,\\tau) = \\frac{a}{2}E_2(\\tau) + \\frac{a}{2} \\int_{0}^{b} E_{1} (| \\tau - \\tau\\p |) \\xi_0(a,b,\\tau\\p) \\d\\tau\\p. \\end{equation} In view of evaluating the integral on the right-hand side of Eq. (\\ref{eq4}), the model under investigation must be further specified to make explicit the $\\tau$-dependence of the function $B^*$. When the distribution of internal sources is uniform in the slab, $B^*$ is independent of $\\tau$ and then the thermal emission of the atmosphere is characterized by the function \\begin{equation} \\label{eq6} Q(a,b,\\tau) = \\int_{0}^{b} R(a,b,\\tau, \\tau\\p) \\d \\tau\\p . \\end{equation} In the conservative case ($a=1$), both functions $\\xi_0$ and $Q$ were studied in detail by Heaslet and Warming \\cite{heaslet-warming1965}-\\cite{heaslet-warming1967} in view of some applications in heat transport theory. Their functions $\\Theta(\\tau)$ and $\\Theta_s(\\tau)$ coincide with our functions $\\xi_0(1,b,\\tau)$ and $(1/4)Q(1,b,\\tau)$, respectively. In the non-conservative case ($0\\leq a<1$), the function $\\xi_0$ is briefly evoked in Section 8 of \\cite{heaslet-warming1968}, while the literature on the function $Q$ is mentioned in a recent article devoted to this function \\cite{chevallier-rutily2003}.\\\\ In what follows, we will be interested exclusively in the second term on the right-hand side of Eq. (\\ref{eq4}), which describes the contribution of the external sources. The exact expression for the $\\xi_0$-function is given for an unspecified albedo $a$, including $a=1$. This particular case is of great importance for our understanding of grey atmospheres in radiative equilibrium. Actually, it is known \\cite{hopf1934} that the solution to Schwarzschild's problem with $a=1$ is the source function, integrated over frequency, of a grey atmosphere in radiative equilibrium. It thus yields the temperature of the atmosphere if the latter may be assumed to be in local thermodynamical equilibrium. This particular problem has been solved exactly in a half-space, that is, for $a=1$, $b=+\\infty$ and $S_0=0$ (since the sources are at infinity). The solution is due to Hopf \\cite{hopf1934}. In a finite slab ($b<+\\infty$), the conservative problem (\\ref{eq1})-(\\ref{eq2}) is equivalent to (\\ref{eq5}) with $a=1$. Yamamoto \\cite{yamamoto1956} failed to solve it exactly, as remarked by Sobolev \\cite{sobolev1962}. Heaslet and Warming's solution \\cite{heaslet-warming1965}-\\cite{heaslet-warming1967} is approximate, in contrast with that given by Kriese and Siewert \\cite{kriese-siewert1970}, who used the Case method. Four-figure tables of the conservative functions $\\xi_0$ and $Q$ are given in \\cite{heaslet-warming1967}, in good agreement with the six-figure tables of \\cite{kriese-siewert1970}.\\\\ The outline of this article is as follows: in Section 2, a recent work dealing with a problem connected with Schwarzschild's problem is summarized. Sections 3 to 5 are devoted to the calculation of the $\\xi_0$-function in the general case, in a conservative atmosphere and in a semi-infinite atmosphere, respectively. The numerical evaluation of the $\\xi_0$-function is tackled in Section 6, which contains a table of this function for $a=1$ and different values of $b$. In Section 7, an application of the conservative case to grey atmospheres in radiative equilibrium is investigated. From the exact expression of the source function integrated over frequency, we deduce a simple relation between the effective temperature of the atmosphere and the temperature of the black body that irradiates it. We also give the temperature of the atmosphere when it is in local thermodynamical equilibrium, a calculation which generalizes Hopf's calculation in a half-space to a finite slab. ",
        "conclusions": "\\label{sec8} The problem of a slab irradiated by an isotropic radiation field is by now a classical one. In the present article, its solution is expressed for unspecified values of $a$ and $b$, without thereby excluding values of astrophysical interest ($a\\sim 1$, $b\\gg 1$). We intend to generalize this solution to boundary conditions of the form $I^+(\\mu) = \\mu^n\\quad (n\\geq 0)$, since in the theory of stellar atmospheres the boundary conditions on the surface $\\tau=b$ are weakly anisotropic (diffusion approximation).\\\\ The conservative case is particularly interesting for the modeling of planetary and stellar atmospheres, because it is equivalent to finding the integrated source function of a grey atmosphere in radiative equilibrium. It thus provides a good test for models which do not presuppose LTE. It is also useful when the LTE condition holds, since it provides the temperature distribution (\\ref{eq90}) that generalizes, for a finite slab, Hopf's temperature distribution (\\ref{eq87}).  \\appendix \\renewcommand{\\theequation}{\\Alph{section}\\arabic{equation}} \\setcounter{equation}{0} \\setcounter{section}{1}"
    },
    "0601/gr-qc0601077_arXiv.txt": {
        "abstract": "We provide some considerations on the excitation of black hole quasinormal modes (QNMs) in different physical scenarios.  Considering a simple model in which a stream of particles accretes onto a black hole, we show that resonant QNM excitation by hyperaccretion requires a significant amount of fine-tuning, and is quite unlikely to occur in nature. Then we summarize and discuss present estimates of black hole QNM excitation from gravitational collapse, distorted black holes and head-on black hole collisions. We emphasize the areas that, in our opinion, are in urgent need of further investigation from the point of view of gravitational wave source modeling. ",
        "introduction": "Various authors \\cite{fhh,AG} suggested that quasinormal ringing could be resonantly excited by clumps of matter falling into the black hole at some appropriate rate. This rate should be such that the spatial separation of the clumps equals a multiple of the typical QNM wavelength. Here we show by a simple model that, in principle, the modes of a black hole can indeed be excited in this way.  We also show that a simple addition of damped sinusoids provides a good fit of the resulting gravitational waveforms: in other words, we can interpret the resonant excitation of the modes in terms of interference between gravitational waves.  For simplicity we consider clumps of mass $\\mu$ much smaller than the black hole mass, $\\mu\\ll M$, so that we can apply a perturbative analysis.  Consider first one single clump falling from infinity. This process was first analyzed in a classical paper by Davis, Ruffini, Press and Price \\cite{davis} (hereafter DRPP).  They found that the total radiated energy is $\\simeq 0.0104\\mu^2/M$, most of it ($0.0091\\mu^2/M$) being emitted as quadrupolar ($\\ell=2$) waves, and that the total energy spectrum is peaked at $\\omega \\sim 0.32/M$, very close to the fundamental $\\ell=2$ QNM frequency $\\omega \\simeq 0.3737/M$.  We can now superpose the waveform and its Fourier transform for the one-particle case to represent two or more particles falling into a black hole and to study interference phenomena (we are assuming that the clumps are non-interacting, so that there is no extra contribution to the total energy-momentum tensor).  A similar sum-over-point-particles approach has been used many times in the past \\cite{pointptcles}, but the particular process we consider here has not been studied before.  Let us first suppose that we drop two particles with a temporal separation of $T$ (which, bearing in mind that the retarded time is the measured quantity, also means a spatial separation of $T$ between the two bodies).  If $\\Psi_1(t,r)$ represents the waveform for the first body (normalized by $\\mu$, where $\\mu$ is the mass of the infalling body), then the total waveform will be \\be \\Psi(t,r)=\\mu \\Psi_1(t,r)+\\mu \\Psi_1(t-T,r)\\,. \\ee Likewise the Fourier transform ${\\tilde \\Psi}(\\omega,r)$ will be \\be {\\tilde \\Psi}=\\mu {\\tilde \\Psi_1}+\\mu e^{i \\omega T}{\\tilde \\Psi_1}\\,. \\ee For $N$ bodies dropped regularly, such that the temporal difference between the dropping of the first and of the last is $T$, we have \\be \\Psi(t,r)=\\sum_{j=0}^{N-1} \\mu \\Psi_1\\left(t-\\frac{i T\\,}{N-1}j,r\\right)\\,,\\quad {\\tilde \\Psi}(\\omega,r)=\\sum_{j=0}^{N-1} \\mu e^{\\frac{i \\omega T}{N-1}j} {\\tilde \\Psi_1}\\,. \\ee The infall of a finite-sized body can also be studied with this formalism \\footnote{ We can think of a finite body as being composed of many particles, the separation between the particles being infinitesimal. The total waveform $\\Psi$ will be a superposition of single-particle waveforms $\\Psi_1$, of the form $ \\Psi(t,r)=\\lim_{N \\to \\infty} \\sum_{j=0}^{N-1} \\frac{\\mu}{N} \\Psi_1\\left(t-\\frac{i T}{N-1}j,r\\right)\\,.$ The Fourier transform of this waveform is $ {\\tilde \\Psi}=\\lim_{N \\to \\infty} \\sum_{j=0}^{N-1}\\f{\\mu}{N} e^{\\frac{i \\omega T}{N-1}j} {\\tilde \\Psi_1}\\,.  $ By taking this limit we obtain the gravitational radiation generated by a massive thin body of length $T$ and total mass $\\mu$.  Summation of the series yields $ \\lim_{N \\to \\infty} \\sum_{j=0}^{N-1}\\frac{\\mu}{N} e^{\\frac{i \\omega T}{N-1}j}= \\mu \\frac{-1+e^{i\\omega T }}{i \\omega T}\\,,$ so that $ {\\tilde \\Psi}=\\mu \\frac{-1+ e^{i \\omega T}}{i \\omega T} {\\tilde \\Psi_1}\\,.$ }.  Incidentally, one can prove that $|\\mu \\sum_{j=0}^{N-1} e^{\\frac{i \\omega T}{N-1}j}|^2\\leq \\mu^2 N^2$, so the total energy $ E_{\\rm rad}\\sim \\omega^2 |{\\tilde \\Psi}|^2 $ radiated by a stream of plunging particles is always less than or equal to the energy radiated by a single point particle with total mass $\\mu N$.  \\begin{figure} \\includegraphics[height=.35\\textheight,angle=270]{plunge2} \\includegraphics[height=.35\\textheight,angle=270]{plunge20} \\caption{Characteristic interference pattern in the total energy radiated by two (left) and twenty (right) particles falling into a Schwarzschild black hole, as a function of the particles' initial separation. As $T\\to 0$ the particles behave as a single particle of mass $(N\\mu)$, and the total energy scales like $(N\\mu)^2$. The red, dashed horizontal line corresponds to the sum of the single-particle energies ($E\\sim N\\times 0.0091\\mu^2/M$ for $N$ particles). In the two-particle case we also overplot the prediction of the two-mode ringdown waveform, Eq.~(\\ref{twoptclesQN}).  The blue, dot-dashed line corresponds to $\\omega_r=0.3737/M$ (the fundamental QNM frequency with $\\ell=2$). To obtain the green, dashed line we choose the ``phenomenological'' value $\\omega_r=0.32/M$ (corresponding to the peak of the one-particle energy spectrum). \\label{fig:powertwotwentyptcles}} \\end{figure}  The total energy radiated in $\\ell=2$ as a function of separation for two and twenty particles is shown in Fig.~\\ref{fig:powertwotwentyptcles}. Let us first consider the two-particle case (left panel).  As we increase the distance between the two bodies the total energy decreases. At the first minimum ($T\\sim 9.9 M$) the total energy radiated is $\\sim 0.0085 \\mu^2/M$. The next maximum is attained at $T\\sim 19.8 M$, the total radiated energy being $\\simeq 0.0206 \\mu^2/M$.  Since the one-particle spectrum is highly peaked at $\\omega \\sim 0.32/M$, it is tempting to explain this behavior as an interference phenomenon. The typical dominant wavelength of the fundamental QNM is $\\lambda=2\\pi/\\omega \\sim 19.6 M$: this means that, for two particles, destructive interference should occur at $T \\sim 9.8 M$ and constructive interference at $T \\sim 19.6 M$, in good agreement with our numerical results. In the limit of large distances, interference becomes negligible and the total radiated energy tends to the sum of the individual particle energies ($\\simeq 0.0182 \\mu^2/M$ for our two particles). This interpretation of the result is confirmed by an explicit calculation of the gravitational waveforms (Fig.~\\ref{fig:waveformstwoptcles}).  \\begin{figure} \\includegraphics[height=.35\\textheight,angle=270]{WF99single} \\includegraphics[height=.35\\textheight,angle=270]{WF198single} \\caption{Two-particle waveforms (solid black lines) are compared with the single-particle waveforms (dashed red lines). In the left panel $T=9.9M$, and the ringdown is suppressed by destructive interference. In the right panel $T=19.8M$, and the ringdown is enhanced by constructive interference. \\label{fig:waveformstwoptcles}} \\end{figure}  The situation for more than two particles is similar. Notice that in the two-particle case, for $T<6 M$ there is always an enhancement on the total radiated energy (as compared to the sum of the single-particle energies).  However, as we increase the number of particles, this enhancement interval decreases.  So we should expect it will be hard to resonantly excite the black hole QNMs using a large number of particles.  \\vskip .2cm  {\\em Interpreting the results in terms of ringdown waveforms - } The DRPP signal is dominated by quasinormal ringing, so it makes sense to try and interpret the preceding results in terms of ringdown waveforms. To start with, we assume our signal has the form \\be \\psi=\\Theta(t) Ae^{-\\omega_i t}\\sin{\\omega _r t}\\,, \\label{ringdownwaveform} \\ee where $\\Theta(t)$ is the Heaviside function. For the fundamental mode of a Schwarzschild black hole $\\omega _r \\simeq 0.3737/M$ and $\\omega_i \\simeq 0.0890/M$.  The amplitude $A$ can be obtained imposing that the total radiated energy in $\\ell=2$ is equal to $0.0091 \\mu ^2/M$.  The total energy radiated in the $\\ell$-th multipole is \\be E_{\\rm rad}=\\frac{1}{64\\pi}\\frac{(l+2)!}{(l-2)!} \\int_{-\\infty}^{-\\infty} \\left(\\frac{d\\psi}{dt}\\right )^2 dt\\,, \\ee which for the waveform (\\ref{ringdownwaveform}) gives \\be E_{\\rm rad}=\\frac{A^2(l+2)!\\omega_r^2}{256\\pi(l-2)!\\omega_i}\\,. \\label{oneptcleQN} \\ee Setting $\\ell=2$ and equating this to $0.0091\\mu^2/M$ gives $A\\simeq 0.441 \\mu$, in qualitative agreement with the DRPP waveform.  Let us now consider two particles by superposing two ringdown waveforms: \\be \\psi=\\Theta(t) Ae^{-\\omega_i t}\\sin{\\omega _r t}+ \\Theta(t-T) Ae^{-\\omega_i (t-T)}\\sin{\\omega _r (t-T)}\\,. \\label{ringdownwaveform2ptcles} \\ee The total energy radiated would be \\be E_{\\rm rad}^{\\rm 2 ptcles}=\\frac{A^2e^{-T\\omega_i}\\omega_r(l+2)!}{128\\pi(l-2)!\\omega_i} \\left (e^{\\omega_i T}\\omega_r+\\omega_r\\cos {\\omega_r T}- \\omega_i\\sin \\omega_r T \\right )\\,. \\label{twoptclesQN} \\ee A plot of this function is shown in the left panel of Fig.~\\ref{fig:powertwotwentyptcles}. It is in good qualitative agreement with the full numerical calculation.  If we use a ``phenomenological'' QNM frequency $\\omega_r=0.32/M$ (corresponding to the peak of the one-particle energy spectrum, and yielding an amplitude $A\\simeq 0.515 \\mu$) instead of the ``true'' fundamental $\\ell=2$ QNM frequency $\\omega_r=0.3737/M$, the agreement between Eq.~(\\ref{twoptclesQN}) and the numerical results becomes even better.  This agreement suggests the possibility to extend our study to Kerr black holes, by simply assuming that $\\omega_r$ and $\\omega_i$ are functions of the black hole's angular momentum.  Using (\\ref{oneptcleQN}) and (\\ref{twoptclesQN}) we get \\be \\frac{E_{\\rm rad}^{\\rm 2 ptcles}}{E_{\\rm rad}}=2e^{-T\\omega_i} \\left (e^{\\omega_i T}+\\cos {\\omega_r T}- \\frac{\\omega_i}{\\omega_r}\\sin \\omega_r T \\right ) \\sim 2(1+\\cos\\omega_r T)\\,. \\ee In the last step we considered rapidly rotating (near-extremal) black holes and modes with $\\ell=m$, for which $\\omega_i \\rightarrow 0$ and $\\omega_r \\rightarrow m/(2M)$.  Therefore, for corotating modes of rapidly rotating black holes the resonance is almost perfect if we choose $T=2\\pi M/m$ (the ratio is then 4, the maximum allowed) while destructive interference results if $T=\\pi M/m$.  In summary: constructive resonance, and consequently QNM excitation, is possible (in principle) when a stream of particles accretes onto a black hole. Unfortunately, the separation between each particle required for an enhancement in the total energy becomes very small as the number of particles increases.  Unless some delicate fine tuning is at work \\cite{AG}, the separation between any two infalling particles or clumps should be driven by a random process, which means that on average there should be no resonance at all. This is confirmed by simulations of generic configurations of infalling matter, where quasinormal ringing is usually absent from the waveforms \\cite{nagar}. ",
        "conclusions": "Most numerical simulations so far consider nearly-axisymmetric situations, so they don't have much to say about the distribution of energy between modes with different values of $(l,m)$ in a realistic black hole merger. This is a crucial issue for data analysis from ringdown waveforms \\cite{BCW}. Recently there has been some remarkable progress in non-axisymmetric simulations of black hole mergers (see eg. \\cite{herrmann} and references therein), but a discussion of these developments is beyond the scope of this paper.  Some estimates of the energy radiated in the plunge phase by solar-mass binaries of rotating black holes are provided by the effective-one-body approach \\cite{EOB}. The error introduced by extrapolating to supermassive black holes should not be too large, since the energy balance of black hole binaries involves a single fundamental scale (their total mass $M$). It is useful to compare the effective-one-body estimate with earlier estimates of the energy radiated in the plunge {\\it plus ringdown} phases, as provided by the Lazarus approach \\cite{Lazarus}. For aligned (anti-aligned) spins and specific angular momenta $j_1=j_2=0.17$ ($j_1=j_2=0.25$) the effective-one-body approach predicts an energy release $\\sim (0.6-0.9)\\% M$ [$\\sim (1-3)\\% M$]. The Lazarus approach finds $\\sim (1.7-1.9)\\% M$ [$\\sim (1.9-2.1)\\% M$], respectively. In principle, the difference of these two estimates should bracket the most reasonable range of ringdown efficiencies. However, the overall error on both estimates is still too large to draw any definite conclusion. Numerical relativity should be able to provide a reliable answer to this problem within a very short time.  \\begin{theacknowledgments} We thank Leonardo Gualtieri, Luis Lehner and Alessandro Nagar for useful discussions. This work was supported in part by the National Science Foundation under grant PHY 03-53180. \\end{theacknowledgments}"
    },
    "0601/astro-ph0601030_arXiv.txt": {
        "abstract": "Observations with the ACS Wide Field Camera and G800L grism can produce thousands of spectra within a single WFC field producing a potentially rich treasure trove of information.  However, the data are complicated to deal with.  Here we describe algorithms to find and characterize spectra of emission line galaxies and supernovae using tools we have developed in conjunction with off the shelf software. ",
        "introduction": "The G800L grism combined with ACS's Wide Field Camera is a powerful combination for obtaining thousands of spectra with relatively modest outlay of HST time.  However, the resulting images are difficult to interpret due to a number of peculiarities including: (1) strong spatially varying sky background; (2) a position dependent wavelength solution; (3) the wide spectral response: a three dimensional flatfield and modeling of the wavelengths contributing to each pixel is required for precise flatfielding; (4) tilted spectra with respect to the CCD grid (the tilt varies over the field); (5) each source is dispersed into multiple orders resulting in much overlap - deep images become confusion limited; (6) zeroth order images of compact sources can easily mimic the appearance of sharp emission features; and (7) the low resolution ($R \\approx 90$ for point sources) means that only high Equivalent Width (EW) features can be discerned, while most familiar features are blends. The Space Telescope European Coordinating Facility has done an excellent job of addressing these peculiarities with the software package {\\sl aXe} (Pirzkal et al.\\ 2001).  Armed with it and a broad band detection image, users can extract 1D and 2D spectra that are sky-subtracted, wavelength-calibrated, flatfielded, and flux calibrated with minimum effort.  Here I describe complimentary techniques I have developed to analyze WFC grism images. Specifically, I describe tools geared to finding emission line sources and supernovae (SNe).  Here I concentrate on my work with the ACS GTO team to search the Hubble Deep Field North (HDFN) for Emission Line Galaxies (ELGs) and work with the PEARS team to find SNe. ",
        "conclusions": "The ACS grism produces amazingly rich datasets.  While the data are somewhat difficult to interpret, tools have been developed to make the most use of these data.  Public access tools like {\\sl aXe\\/} are readily available to remove most of these complications and extract 1D and 2D spectra.  Here I have shown how some common manipulations of the data (such as geometric correction and flatfielding) allow interesting sources such as emission line galaxies and supernovae to be efficiently found."
    },
    "0601/astro-ph0601206_arXiv.txt": {
        "abstract": "The observed radial velocities and transverse velocities of individual stars in the globular cluster M15 are implemented as inputs to a fully non-parametric code (CHASSIS) in order to estimate the equilibrium stellar distribution function and the three-dimensional mass density profile, from which are estimated the mass-to-light ratio, enclosed mass and velocity dispersion profiles. In particular, the paper explores the possibility of the existence of a central black hole in M15 via several runs that utilize the radial velocity data set which offers kinematic measurements closer to the centre of the cluster than the proper motion data. These runs are distinguished from each other in the choice of the initial seed for the cluster characteristics; however, the profiles identified by the algorithm at the end of each run concur with each other, within error bars, thus confirming the robustness of CHASSIS. The recovered density profiles are noted to exhibit unequivocal flattening, inner to about 0.0525pc.  Also, the enclosed mass profile is very close to being a power-law function of radius inside 0.1pc and is not horizontal. Simplistically speaking, these trends negate the possibility of the central mass to be concentrated in a black hole, the lower bound on the radius of the sphere of influence of which would be $\\gtrsim$0.041pc, had it existed. However, proper analysis suggests that the mass enclosed within the inner 0.01pc could be in the form of a black hole of mass $\\sim{10^3}$M$_{\\odot}$, under two different scenarios, which are discussed.  The line-of-sight velocity dispersion is visually found to be very similar to the observed dispersion profile. The enclosed mass and velocity dispersion profiles calculated from runs done with the proper motion data are found to be consistent with the profiles obtained with the radial velocity data. ",
        "introduction": "\\label{sec:intro} \\noindent Intermediate Mass Black Holes (IMBH) can serve as vincula to bridge the discontinuity that is otherwise apparent between the extremities of the black hole mass distribution. Lately, observational data, indicating the possibile existence of IMBHs, has been accruing; some of the recent analyses are due to \\citet{simonnatur}, \\citet{roberts04}, \\citet{terashima04} and \\citet{cropper04}, among others. A comprehensive review of the subject of IMBH, including similar pieces of observational evidence, is presented in \\citet{miller04}. The case for these intermediate mass black holes has also been buttressed by the predictions of a black hole of mass of the order of 10$^3$M$_{\\odot}$ in the center of the globular cluster M15 by \\citet{gebhardt97} and \\citet{gressen02}. However, the analysis by \\citet{simon03} implies a model for the central structure of M15 that does not necessarily require a central black hole.  The paradigm of globular cluster evolution that we are aware of, from works by \\citet{spitzer87} and \\citet{meylan97}, is indeed challenged if an IMBH is needed to model a globular cluster such as M15. According to conventional understanding, in a globular cluster with few primordial binaries, core collapse is stalled by the formation of binaries via three-body processes. By the time core collapse is halted, the central density in the cluster is quite high, but this happens on time scales of the order of a few gigayears, (many times the half-mass relaxation time). In other words, such clusters spend a long time with cores in their centers. Runaway growth of a massive object can occur via collisional processes, and result in an IMBH, if circumstances are right. \\citet{simon02} suggest that runaway growth is possible in clusters with initial half mass relaxation times less than 25Myr; they also indicate that the same for M15 is probably higher than this figure. Moreover, on the basis of their modelling, \\citet{simon02} also suggest that star clusters older than about 5Myrs and current half-mass relaxation times $\\lesssim$100Myrs are expected to contain an IMBH. The current half mass relaxation time of M15 is about 2.24Gyrs \\citep{harris96}, which is nearly 25 times higher than this limiting value of 100Myrs. Thus, it is unlikely that an IMBH has formed in the center of M15 by a runaway process. An alternative mechanism that could potentially form IMBHs in the centers of dense stellar clusters was suggested by \\citet{miller02}. In this picture, an initial seed is contemplated to undergo a gradual collisional growth to transform into a massive object.  Thus, if claims of an IMBH in M15 are true, then we can ratiocinate that a fresh approach needs to be invoked in order to understand globular cluster evolution. In light of such profound implications of these predictions, it is important to make an independent evaluation of the central structure of M15. This is attempted in this paper. The algorithm used in this exercise is a fully non-parametric code (CHASSIS) that was used to constrain the mass of the central black hole in our Galaxy \\citep{dalpras} and was calibrated against N-body realizations of the Hyades \\citep{simonH} and Arches \\citep{simonA} clusters \\citep{dalzwart}. The input data for this algorithm is taken from the observed radial velocities of 64 stars in M15 \\citep{gressen02} and the observed proper motions of 1764 stars in this globular cluster \\citep{mcnamara03}.  This introductory section is followed by a brief mention of the salient points of the data sets that have been used as inputs to the algorithm. In Section~\\ref{sec:algorithm}, the basic scheme of CHASSIS is briefly discussed. A desription of the procedure followed in this paper, to calculate the mass-to-light ratio, has been discussed in Section~\\ref{sec:ml}. The results obtained by the algorithm are presented in Section~\\ref{sec:results}. This is followed by Section~\\ref{sec:discussions}. Finally, the paper is rounded up by a conclusive section. ",
        "conclusions": "\\label{sec:conclu} \\noindent This paper reports the characteristics of the globular cluster M15, as estimated by the algorithm CHASSIS. To achieve this, the observed radial velocities and proper motions of individual stars in M15 are implemented as inputs to the code. The recovered three dimensional mass density distribution is found to be distinguished by a core central to about 0.0525pc. The enclosed mass profile that is calculated from this density distribution is noted to have the form of a single power-law in radius, inner to about 0.1pc. These distinct features appear to be in contradiction with a hypothesis that supports a central IMBH in M15. This argument appears to be further buttressed by the following rationale.  The mass that is found to be enclosed within the innermost radial bin (the extent of which corresponds roughly to the innermost radial location at which a velocity information is reported) at 0.0105pc is of the order of 1000M$_\\odot$. If we assume that this mass was entirely in the form of a central black hole, the radius of the sphere of influence of such a black hole ($r_s$) is found to be at least as high as 0.041pc. Here, $r_s$ is calculated as the radius at which the observed (or estimated) central velocity dispersion equals the circular velocity around the black hole, (see Equation~\\ref{eqn:r_s}). Since this value is nearly 4 times the innermost radius at which the expected trends in the density and mass profiles are noted to be absent, the inference is that there exists no IMBH at the M15 centre. However, such a conclusion can be challenged in the following ways; firstly the assumption that all the mass is in a black hole maybe wrong. Secondly, the extent of the influence of a central black hole is dependant on the phase space distribution and various forms of this could yield values of $r_s$ that may differ by a factor of few from the value that Equation~\\ref{eqn:r_s} yields.  If only a fraction of the mass that is enclosed within the first radial bin is in the form of a black hole, then the sphere of influence of such a black hole would be proportionately smaller than 0.041pc. It is possible that it is even interior to the edge of our first bin. In this case, we would not be able to spot the expected trends in our reslts, unless the resolution of the observed kinematic data sets improved. Of course the density profile then would be broken; density will rise steeply at some radius $<0.0105$pc, after staying flat 0.0525pc inwards, approximately. In this scenario, for us to miss the tell tale signs of the central IMBH, $r_s$ of this black hole, as given by Equation~\\ref{eqn:r_s}, has to be $<0.0105$pc. This implies that the mass of this IMBH will be less than about a quarter of the mass that is found enclosed within the first radial bin (presented in Table~\\ref{tab:results}).  The other source of possible concern would be the way the sphere of influence has been quantified above. It is possible to imagine orbital distributions around the assumed black hole, which are such that the influence of the central IMBH is felt at radii much lower than ($< 1/4$ times lower than) the estimated value of 0.041pc. In other words, it is possible that $r_s$ is not neccessarily the radius where the central dispersion equals the circular speed around the black hole. In this case too, the density might steepen up and the mass profile turn horizontal inner to where the edge of the first radial bin is with the currently available radial resolution. In this case, we can conclude that the upper limit on the mass of the hypothesized central IMBH will be $\\sim$1000M$_\\odot$."
    },
    "0601/astro-ph0601113_arXiv.txt": {
        "abstract": "Using the deep multi-wavelength MUSYC, GOODS, and FIRES surveys we construct a stellar mass-limited sample of galaxies at $2<z<3$. The sample comprises 294 galaxies with $M>10^{11}$\\,\\msun\\ distributed over four independent fields with a total area of almost 400\\,arcmin$^2$. The mean number density of massive galaxies in this redshift range $\\rho (M>10^{11}\\,M_{\\odot}) = (2.2 \\pm 0.6) \\times 10^{-4}\\,h_{70}^3$\\,Mpc$^{-3}$. We present median values and 25$^{\\rm th}$ and 75$^{\\rm th}$ percentiles for the distributions of observed $R_{\\rm AB}$ magnitudes, observed $J-K_s$ colors, and rest-frame ultra-violet continuum slopes, $M/L_V$ ratios, and $U-V$ colors. The galaxies show a large range in all these properties. The ``median galaxy''  is faint in the observer's optical ($R_{\\rm AB}=25.9$), red in the observed near-IR ($J-K_s = 2.48$), has a rest-frame UV spectrum which is relatively flat in $F_{\\lambda}$ ($\\beta= -0.4$), and rest-frame optical colors resembling those of nearby spiral galaxies ($U-V = 0.62$). We determine which galaxies would be selected as Lyman break galaxies (LBGs) or Distant Red Galaxies (DRGs, having $J-K_s>2.3$) in this mass-limited sample. By number DRGs make up 69\\,\\% of the sample and LBGs 20\\,\\%, with a small amount of overlap. By mass DRGs make up 77\\,\\% and LBGs 17\\,\\%. Neither technique provides a representative sample of massive galaxies at $2<z<3$ as they only sample the extremes of the population. As we show here, multi-wavelength surveys with high quality photometry are essential for an unbiased census of massive galaxies in the early Universe. The main uncertainty in this analysis is our reliance on photometric redshifts; confirmation of the results presented here requires extensive near-infrared spectroscopy of optically-faint samples. ",
        "introduction": "The properties of massive galaxies at high redshift place important constraints on galaxy formation models (see, e.g., {Kauffmann} \\& {Charlot} 1998; {Nagamine} {et~al.} 2005). The ``standard'' and most successful method for finding distant galaxies is the Lyman dropout technique, which relies on the strong break in the rest-frame ultra-violet (UV) spectra of high redshift galaxies blueward of the Lyman limit ({Steidel} {et~al.} 1996, 1999). However, it is not yet clear whether these galaxies are representative of the high redshift galaxy population, in particular at the high-mass end. As the Lyman break selection requires that galaxies are very bright in the rest-frame UV it may miss objects that are heavily obscured by dust or whose light is dominated by evolved stellar populations.  Advances in instrumentation have made it possible to select galaxies in complementary ways, and recent studies have demonstrated that the universe at $z>2$ is much more diverse than had been realized. Among recently identified ``new'' galaxy populations are submm galaxies (e.g., Smail et al.\\ 2004), distant red galaxies (DRGs) selected by the criterion $J-K_s>2.3$ ({Franx} {et~al.} 2003; {van Dokkum} {et~al.} 2003), ``IRAC Extremely Red Objects'' (IEROs; {Yan} {et~al.} 2004), and ``$BzK$'' objects (Daddi et al.\\ 2004).  The current situation is somewhat confusing, as the relative contributions of the various newly identified galaxy populations to the stellar mass budget and the cosmic star formation rate are still unclear. Furthermore, as emphasized by, e.g., {Adelberger} {et~al.} (2005) and Reddy et al.\\ (2005), there can be considerable overlap between selection techniques, which makes it difficult to interpret the plethora of windows on the high redshift Universe.  Ideally samples of high redshift galaxies are selected not by color or luminosity but by stellar mass. Whereas  luminosities and colors can vary dramatically due to star bursts and the presence of dust the mass evolution of galaxies is probably gradual. Also, galaxy formation models can predict masses with somewhat higher confidence than luminosities and colors. Stellar masses of distant galaxies are usually determined by fitting stellar population synthesis models to broad band photometry. Although there are significant systematic uncertainties, the stellar mass of a galaxy is usually better constrained than the instantaneous star formation rate, age, or dust content (e.g., {Shapley} {et~al.} 2001; {Papovich} {et~al.} 2001; {van Dokkum} {et~al.} 2004; {F{\\\" o}rster Schreiber} {et~al.} 2004).  In this Letter we explore the properties of a stellar-mass limited sample of galaxies. The main purpose is to measure ``basic'' aspects of massive galaxies to compare with simulations of galaxy formation: their density and colors. A secondary goal is to quantify selection biases introduced by two of the most widely used techniques for identifying distant galaxies: the Lyman break technique and the $J-K$ color selection of Franx et al.\\ (2003). We assume $\\Omega_m=0.3$, $\\Omega_{\\Lambda} =0.7$, and $H_0 = 70$\\,\\kms\\,Mpc$^{-1}$. All magnitudes are on the Vega system, unless identified as ``AB''.  \\begin{figure}[t] \\epsscale{1.1} \\plotone{f1.eps} \\caption{ Relation between stellar mass and observed total $K_s$ magnitude for galaxies at $2<z_{\\rm phot}<3$. The solid line shows our selection limit of $M>10^{11}$\\,\\msun. The broken line shows the approximate photometric limit of the MUSYC survey, which accounts for the majority of galaxies in our sample. There are very few galaxies in FIRES and GOODS with $M>10^{11}$\\,\\msun\\ and $K_s>21.3$, and we estimate that the completeness of the full sample of $M>10^{11}$\\,\\msun\\ galaxies is $\\approx 95\\,\\%$. \\label{kmass.plot}} \\end{figure} ",
        "conclusions": "The main result of our analysis is that massive galaxies at $z\\sim 2.5$ span a large range in rest-frame UV slopes, rest-frame optical colors, and rest-frame $M/L_V$ ratios, indicating significant variation in dust content, star formation histories, or both. This result is not surprising in the light of the recent discoveries of DRGs, IEROs, and other populations. Here we have quantified the median colors and their range for a uniformly selected, large, mass-limited sample.  The large variation in the rest-frame color distributions of our mass-limited sample implies that ``standard'' color selection techniques produce biased samples. We consider two of the two most widely used selection techniques in this redshift range: the Lyman break technique of Steidel and collaborators and the $J-K_s>2.3$ Distant Red Galaxy (DRG) selection ({Franx} {et~al.} 2003; van Dokkum et al.\\ 2003). Lyman break galaxies are identified in the following way. From the best-fitting {Bruzual} \\& {Charlot} (2003) SEDs (which include absorption due to the Ly\\,$\\alpha$ forest) we calculated synthetic colors in Steidel's $U_n G \\cal R$ system. To qualify as an LBG an object has to have $R_{\\rm AB}<25.5$ and synthetic $U_n G \\cal R$ colors which place it in the Lyman break, BX, or BM selection region (see {Steidel} {et~al.} 2003, 2004). Combined these criteria provide a continuous selection of  galaxies over the redshift range considered here.\\footnote{A LBG in this definition is therefore an object which has $R_{\\rm AB}<25.5$, $2<z_{\\rm phot}<3$, and is either a classical ``U-dropout'' or a BX/BM object.} Figure \\ref{jkr.plot} illustrates the LBG and DRG selection techniques, as applied to our sample. DRGs with $J-K_s>2.3$ are indicated by red symbols and LBGs by blue symbols. The DRG limit and the ``standard'' photometric LBG limit of $R_{\\rm AB} =25.5$ are also indicated.\\footnote{We note that not all galaxies with $J-K_s>2.3$ have redshifts in the range $2<z<3$: to $K_s=21$ we find that $\\sim 50$\\,\\% are in this redshift range with the rest about equally split between $z<2$ and $z>3$ galaxies.}  \\begin{figure}[t] \\epsscale{1.1} \\plotone{f4.eps} \\caption{ Correlation between observed $J-K_s$ color and $R_{\\rm AB}$ magnitude. The majority of the galaxies are faint in $R$ and red in $J-K_s$. Red symbols denote DRGs with $J-K_s>2.3$; blue points denote LBGs with $R_{\\rm AB}<25.5$. Red symbols with blue circles fall in both categories. Massive LBGs have blue near-IR colors and are bright in the observed $R$ band. The inset shows the optical color distribution of galaxies with $R<25.5$. Only $\\sim 50$\\,\\% of optically-bright massive galaxies have the colors of LBGs. \\label{jkr.plot}} \\end{figure}  By number, DRGs make up 69\\,\\% of the sample and LBGs 20\\,\\%. The DRG and LBG samples do not show much overlap: only 7\\,\\% of objects fall in both categories. By rest-frame $V$-band luminosity DRGs contribute 64\\,\\% and LBGs 32\\,\\%. By mass DRGs contribute 77\\,\\% and LBGs 17\\,\\%. Together, the LBG and DRG techniques identify 82\\,\\% of massive galaxies by number and 84\\,\\% by mass. Most of the remaining galaxies are optically faint, slightly bluer than the $J-K_s=2.3$ limit, and have redshifts $z<2.5$. Approximately 85\\,\\% of them fall in the ``$BzK$'' selection region (Daddi et al.\\ 2004), which is optimized for galaxies at $1.4<z<2.5$. We note that the relatively small fraction of LBGs in the sample is not solely due to the imposed $R_{\\rm AB}<25.5$ limit. As shown in the inset of Fig.\\ \\ref{jkr.plot} only $\\sim 50$\\,\\% of galaxies with $R_{\\rm AB}<25.5$ have the rest-UV colors of LBGs, and this fraction decreases going to fainter $R$ magnitudes: when no $R$ limit is imposed we find that $\\sim 1/3$ of the galaxies have the colors of LBGs. The underlying reason is the broad distribution of $\\beta$.  It is clear from Fig.\\ \\ref{jkr.plot} that the LBG selection produces very different samples of massive high redshift galaxies than the DRG selection. Both samples are biased: LBGs are too blue and DRGs are too red when compared to the median values of the full sample. This bias is shown explicitly by the blue and red histograms in Fig.\\ \\ref{mluv.plot}. The Lyman break criteria were designed to find star forming galaxies, but {Shapley} {et~al.} (2004) and {Adelberger} {et~al.} (2005) have argued that they can also be used to find massive galaxies at high redshift. This is obviously the case, but we find that the colors and $M/L_V$ ratios of massive LBGs are not representative for the full sample of massive galaxies. Surveys over the full set of optical/near-IR passbands from $U$ through $K$ are essential to obtain representative samples of massive galaxies.  The main uncertainty in this analysis  is the reliance on photometric redshifts. We estimated the effect of this uncertainty by randomly perturbing the redshifts using a Gaussian distribution with dispersion $\\Delta z/(1+z)=0.12$, and repeating the selection and analysis. Despite significant migration of galaxies in and out of the $2<z<3$ redshift range the values in Table 1 change by only $\\sim 10$\\,\\%. We note, however, that there may be subtle systematic biases which can have significant effects, in particular on the derived masses (see, e.g., Shapley et al.\\ 2005). Comprehensive tests of the techniques employed here and in other studies of infra-red selected samples (e.g., Dickinson et al.\\ 2003, Rudnick et al.\\ 2003) are urgently needed, and will become feasible with the introduction of multi-object near-IR spectrographs on 8m class telescopes."
    },
    "0601/astro-ph0601439_arXiv.txt": {
        "abstract": "{The kinematic structure of a sample of planetary nebulae, consisting of 23 [WR] central stars, 21 weak emission line stars ({\\it wels}) and 57 non-emission line central stars, is studied.  The [WR] stars are shown to be surrounded by turbulent nebulae, a characteristic shared by some {\\it wels} but almost completely absent from the non-emission line stars.  The fraction of objects showing turbulence for non-emission-line stars, {\\it wels} and [WR] stars is 7\\%, 24\\%\\ and 91\\%, respectively. The [WR] stars show a distinct IRAS 12-micron excess, indicative of small dust grains, which is not found for {\\it wels}. The [WR]-star nebulae are on average more centrally condensed than those of other stars.  On the age-temperature diagram, the {\\it wels} are located on tracks of both high and low stellar mass, while [WR] stars trace a narrow range of intermediate masses. Emission-line stars are not found on the cooling track. One group of {\\it wels} may form a sequence {\\it wels}--[WO] stars with increasing temperature.  For the other groups both the {\\it wels} and the [WR] stars appear to represent several, independent evolutionary tracks. We find a discontinuity in the [WR] stellar temperature distribution and suggest different evolutionary sequences above and below the temperature gap. One group of cool [WR] stars has no counterpart among any other group of PNe and may represent binary evolution. A prime factor distinguishing {\\it wels} and [WR] stars appears to be stellar luminosity. We find no evidence for an increase of nebular expansion velocity with time. ",
        "introduction": "Some central stars (cores) of planetary nebulae (PNe) show broad, stellar emission lines similar to Wolf-Rayet stars. Although superficially similar, they differ from classical WR stars in their degenerate structure, much lower masses, a wider range of temperatures, and being limited almost exclusively to carbon-rich stars (there are very few counterparts to the massive WN stars). Similar to their massive WC-star counterparts, they are deficient in hydrogen. They form a class named [WC]-type which historically was further divided into early [WCE] and late [WCL] groups. A less numerous class of PNe central stars show narrower and weaker emission lines than the [WR]-type stars. These are named {\\it wels} (weak emission-line stars), as defined in Tylenda et al. (1993).  The evolutionary status of the [WR]-type stars is still very uncertain, and it is unclear whether there is any evolutionary relation to the {\\it wels}. Nebular data suggest an evolutionary sequence [WC11]$\\rightarrow$[WO2] (Zijlstra et al.  1994, Acker et al. 1996, Pe\\~na et al. 2001), followed by PG1159 stars (Werner et al. 1992).  Acker \\& Neiner (2003) propose a sequence: [WC11]$\\rightarrow$[WC4]$\\rightarrow$[WO4]$\\rightarrow$[WO1].  Parthasarathy et al. (1998) have suggested that the {\\it wels} are related to the PG1159 stars: [WCL]$\\rightarrow$[WCE]$\\rightarrow${\\it wels}/PG1159$\\rightarrow$PG1159.  But this sequence is by no means proved; the precise location of the {\\it wels} in relation to the [WR] stars is under discussion (Marcolino \\& de Araujo 2003). Pe\\~na et al. (2001) also argue against too closely identifying {\\it wels} with PG1159 stars (see also Koesterke 2001).  Not all PG1159 are hydrogen-poor (Dreizler et al. 1996), showing that not all PG1159 stars will have evolved from [WR] stars. Neither are all {\\it wels} hydrogen-poor: a dual evolutionary sequence may be expected.  The lack of hydrogen is often taken to indicate that the [WR] stars have undergone a late thermal pulse (or helium flash), either during the post-AGB evolution (a so-called Late Thermal Pulse or LTP) or on the white dwarf cooling track (Very Late Thermal Pulse or VLTP: Herwig 2001, Hajduk et al. 2005). Following such a pulse, the star rejuvenates and retraces part of its earlier evolution.  This predicts that nebulae around [WR] stars should be more evolved than around non-[WR] stars, but this is not confirmed by observations; the properties of the PNe around [WR] stars do not differ from those around other central stars (G\\'orny 2001).  In this paper we discuss a much larger sample of [WR] stars and {\\it wels} than has previously been available. We apply velocity-field analysis to locate the objects on derivatives of the HR diagram. In Sect.\\,2 we describe the methods applied for analysis of PNe with specific attention towards automatic model fitting. Sect.\\,3 presents the 33 newly analyzed PNe. In Sect.\\,4 we discuss the nebular and stellar parameters for the full sample of 101 PNe. In Sect.\\,5 we discuss possible evolutionary relations.  Crowther et al. (1998) have refined the WC and the WO  schemes used for WR stars and define a unified classification for  massive WR and low-mass [WR] stars. Acker \\& Neiner (2003) developed this  scheme for a sample of 42 PN central stars, classified into [WO 1-4] and  [WC 4-11] stars. This classification (used in the present paper) is based on the ionization level of the  elements (depending on the wind-temperature of the PN core), showing essentially carbon lines for the coolest stars and oxygen lines for the hottest ones. ",
        "conclusions": "We have deduced for 33 PNe the velocity field. Since this was done in most cases by considering two nebular lines we regard as robust the mass averaged value only. For six PNe the analysis was based on three or more lines and {\\it is} therefore more reliable. The discussion of details of velocity field like e.g. turbulence should be therefore treated with caution while the analysis based on the mass averaged velocity (ages, masses) should be reliable. We combined this sample with earlier published data and discussed the full sample of 101 PNe.  We found no evidence for an increase in nebular expansion velocity with time. This correlation was checked against $T_{\\rm b-b}$ which should correspond to stellar age and on the nebular dynamical age i.e. on two independent characteristics. The [WR] stars show a tendency to low ratios of $T_{\\rm b-b}/V_{\\rm exp}$. Comparison with the Sch\\\"onberner et al. (2005a) models suggests that these objects have more centrally condensed nebulae, indicative of increasing mass-loss rate with time on the AGB. Emission-line stars are not found on the cooling track.  We find a correlation between IRAS 12-micron excess and the [WR] stars.  This excess is not shown by the {\\it wels}.  The cause of 12-micron excess is typically a population of small grains. These may form in the wind, or form in the collision between the stellar wind and the nebula. Without exception, the {\\it wels} do not show the 12-micron excess.  We confirm the stong correlation between turbulence and [WR] stars. However, the correlation is not one to one: there are a few non-emission-line stars with turbulence, and a few [WR] stars without.  The {\\it wels} appear intermediate in occurrence of turbulence.  The distribution in the age-temperature plane suggests that there are several groups of [WR] stars and {\\it wels}. One group shows evolved nebulae with cool stars and likely originates from a Very Late Thermal Pulse. The hot [WO] stars show a narrow distribution corresponding to a small range in core mass.  Some {\\it wels} fall on the same sequence but at lower temperature--they may be the [WO] progenitors. Other {\\it wels} are distributed on higher and on lower mass tracks, which lack [WR] stars. Finally, a group of cool [WC] stars with IR-bright compact nebulae have no counterpart among any other group of PNe.  The relation between the [WR] stars and the {\\it wels} is therefore a complicated one, with neither group being uniform. For the majority of the {\\it wels}, there is evidence against them being evolutionarily related to the [WR] stars. There is a noticeable apparent difference in core mass and dust properties and nebular turbulence.  A small group (5 objects) of {\\it wels} have characteristics consistent with [WO] star progenitors. Apart from this {\\it wels}--[WO] sequence, the {\\it wels} contain stars at both the high and low luminosity end of the distribution.  The former are likely stars where the strong radiation field drives the wind, but are not necessarily H-poor. The latter may be H-poor, failed [WR] stars where the luminosity is too low to maintain a strong [WR] wind. For the [WO] progenitors, the determining factor appears to be stellar temperature: the particular ionization balance may reduce the efficiency of the radiation-driven wind.  The determined discontinuity in temperatures of [WR]-type stars is in favor of the suggested separate evolutionary channels of cool and hot [WR] objects."
    },
    "0601/astro-ph0601325_arXiv.txt": {
        "abstract": "The results of a spectral analysis, using {\\it XMM-Newton} and {\\it Chandra} data of the brightest ultra luminous X-ray source in the nearby galaxy M82, are presented. The spectrum of M82 X-1, was found to be unusually hard ( photon spectral index $\\Gamma \\approx 1$ ) with a sharp cutoff at $\\approx 6$ keV. Disk black body emission model requires a  nonphysically high temperature. Instead, the spectrum is better described, with a lower reduced $\\chi^2$, as emission due to nearly saturated Comptonization of photons in an optically thick ($\\tau \\approx 10-30$, depending on the geometry) plasma having a temperature $kT \\approx 2$ keV. This is in contrast to the high energy spectra of other black hole systems, which are relatively steeper ($\\Gamma > 1.5$) and hence are modeled as un-saturated thermal and/or non-thermal Comptonization of soft photons, in an optically thin ($\\tau \\approx 1$) high temperature plasma. An Iron line emission which is marginally resolved ($\\sigma \\sim 0.2$ keV) is required to fit the data. We argue that the standard geometries for the X-ray producing region, which are an optically thin inner disk or an uniform/patchy corona on top of a cold disk, are not applicable to this source. Alternatively, the geometry of the X-ray producing region could be a large sphere surrounding a cold accretion disk or an optically thick inner disk region which cools by bremsstrahlung self-Comptonization. For the latter scenario, such an inner disk region, whose effective optical depth to absorption is less than unity, is expected in the standard accretion disk theory for near Eddington accretion rates. ",
        "introduction": "{\\it Chandra} observations of nearby galaxies have detected several non-nuclear X-ray point sources, some of which have isotropic luminosities $> 10^{39}$ ergs/s and are called Ultra Luminous X-ray sources, ULX \\citep{Mat01,Zez02}. These results have confirmed and extended the earlier evidence for the existence of such sources by the {\\it Einstein} survey \\citep{Fab89, Fab87} and the ROSAT detection of  $\\sim 100$ ULX in nearby galaxies \\citep{Col02}. Systematic studies of the {\\it Chandra} observations have revealed that ULX exist in  galaxies of different morphological types. Since, by definition, ULX emit radiation at a rate larger than the Eddington luminosity for a ten-solar mass black hole, they are believed to harbor a black hole of mass $50 \\, M_\\odot\\! < \\! M \\! <\\! 10^4 \\,M_\\odot$, where the upper limit is constrained by the fact that a more massive black hole would have settled into the nucleus due to dynamical friction \\citep{Kaa01}. If this is true, then the differences and similarities between the spectral properties of such intermediate mass black hole systems (IMBH) with those of Galactic black holes (GBH) and Active Galactic Nuclei (AGN), are expected to enhance our understanding of black hole systems in general. Alternatively, the ULX sources may be systems emitting at super-Eddington luminosities \\citep{Beg02}, or their emission may be beamed from a geometrically thick accretion disk \\citep{Rey97, Kin02}. It has been argued that in the latter case, such thick \"funnel\" shaped disks can enhance the observed flux by just a factor of few \\citep{Mis03}. In any case, the complexities involved in explaining their origin \\citep{Por02,Tan00,Mad01} and the inferred high accretion rates \\citep{Kin01} make these enigmatic sources an important class of X-ray systems, whose dominant radiative mechanism and geometry, needs to be studied and constrained.  The spectra of ULX are expected to be different from GBH and AGN, if the black hole is an IMBH and/or if their luminosity exceeds the Eddington limit. Earlier analysis of ASCA data revealed that the spectra of some ULX can be described as radiation from a black-body emitting disk with an additional high energy power-law \\citep{CM99}, which is also the standard radiative model for  the high-soft spectral state of Galactic black hole systems. The inferred inner disk temperature from these initial analysis, was found to be too high ($kT_{in} \\sim 1.1-1.8 $ keV) for an accreting IMBH \\citep{CM99,Mak00}. However, recent analysis of {\\it XMM-Newton} data have revealed the presence of a cool accretion disk component ($kT_{in} \\sim 0.1-0.5 $ keV) in   NGC 1313 X-1, X-2 \\citep{Mil03} and M81 X-9 \\citep{Mil04a} which reaffirmed that these systems may indeed harbor an IMBH. The spectra of some other ULX can be represented by a single power-law suggesting a similarity with the low-hard state of GBH, where the disk emission is either absent or weak. In fact, transitions between the two spectral states, similar to those seen in GBH, have been reported in the ULX NGC 1313 X-1\\citep{CM99} and two sources in IC342 \\citep{Kub01}. Analysis of {\\it Chandra} data for a large number of ULX \\citep[e.g.][]{Col04,Swa04} have revealed that the spectra of most of these sources can be represented by a power-law with a typical photon index, $\\Gamma \\approx 1.6-2.0$. This spectral index is similar to those of GBH in the hard state and Seyfert 1 AGN \\citep{Zdz96a}. However, most of these ULX are too faint for any further detailed spectral analysis.     One of the most intriguing ULX is the brightest X-ray source in  the starburst galaxy M82 that is located $\\sim 170 pc$ away from the dynamical center. The isotropic luminosity of this source, M82 X-1, can exceed 10$^{41}$ ergs s$^{-1}$ indicating that either it is radiating well above the Eddington limit, or it harbors a black hole $ > 100 M_\\odot$, even if some moderate beaming (a factor of few) is taken into account. ASCA observation have shown that this source is highly variable with its luminosity in $0.5 - 10$ keV band changing from 4.5$\\times$ 10$^{40}$ to 1.6 $\\times$ 10$^{41}$ ergs s$^{-1}$ \\citep{PG99}. The presence of a $54$~mHz QPO in its X-ray light curve suggests that the source is a compact object and not highly beamed \\citep{Str03}. If this QPO is related to the  $0.8-3$ Hz QPO observed in galactic sources, it would imply that its black hole mass, $M_{BH} \\sim 100-300 M_\\odot$. Its high luminosity and relative proximity makes this source an ideal ULX whose spectra can be analyzed in detail.  \\cite{Swa04} have analyzed and fitted a large number of ULX observed by {\\it Chandra} including M82 X-1. After correcting for pileup, they report that the 1999 {\\it Chandra} observation of M82 X-1 can be fitted by a power-law with spectral index $\\Gamma = 1.17^{+0.15}_{-0.14}$ and $N_H = 0.98^{+0.12}_{-0.12} \\times 10^{22}$ cm$^{-2}$. \\cite{Str03} mention that the {\\it XMM-Newton} data can be fitted both by a disk black body emission with an unphysically high temperature of $\\sim 4$ keV and a Comptonization model, but  they do not quantify the parameter values obtained nor discus the physical interpretation of the fits. Recently, \\cite{Fio04}, fitted the same {\\it XMM-Newton} data with the bulk Comptonization model and obtained a power-law index $\\Gamma \\sim 2$ indicating that this source is similar to the soft/high state of black hole binaries. However, as shown in this work, the column density required for such a fit, $N_H > 5 \\times 10^{22}$ cm$^{-2}$, is not consistent with some of the {\\it Chandra} observations. Note, that the {\\it Chandra} ACIS data essentially measures the spectrum in the $0.5-5$  keV, since the effective area of the detector decreases rapidly for energies greater than $5$ keV. On the other hand, the {\\it XMM-Newton} EPIC data measures the spectrum in the $3-8$ keV range, since below $3$ keV the counts have contribution from the diffuse background. This complementary nature of the two types of observations makes it imperative to take into account both of them in any spectral analysis. In particular the {\\it Chandra} observations give a better measure of the neutral column density, while the detailed higher energy shape of the spectrum can be constrained by {\\it XMM-Newton} observations.  Since the source is known to be variable, in principle,  any spectral analysis should be based only on simultaneous observations, which currently do not exist. However, here we take a pragmatic approach and use the non-simultaneous {\\it Chandra} observations of the source, to test whether the parameters obtained are consistent with the results of the {\\it XMM-Newton} analysis. Our motivation is to obtain the simplest physical radiative model for the source, which in turn, will highlight any differences between this ULX and other black hole systems. ",
        "conclusions": "The spectrum of M82 X-1 is unusually hard and hence the simplest radiative model which can describe it is saturated thermal Comptonization. This is in contrast, to other black hole systems, where the spectra is generally modeled as un-saturated thermal/non-thermal Comptonization. The spectra of other ULX are steeper with photon index, $\\Gamma \\sim 1.7$ \\citep{Mil03,Mil04a,Mil04b} similar to those of AGN and Galactic black hole systems. Thus either, M82 X-1 is different from these systems or at least it was in a rare and unique spectral state during the {\\it XMM-Newton observation}.   The spectral fit parameters obtained imply that the geometry of the X-ray producing region for M82 X-1, is not the same as that of other black hole systems. The large optical depth $\\tau$ \\greq $10$  would prevent the entry of any external soft photons into the thermal plasma. Hence, the standard geometry of an hot inner disk, which Comptonizies photons from an outer cold disk \\citep{Sha76} is not applicable to this source. Similarly, soft photons will not be able to enter an optically thick patchy corona or active region \\citep{Pou99} on top of a cold disk. On the other hand, the soft photons from a cold accretion disk would diffuse out through an optically thick corona if it uniformly covers the disk \\citep{Lia77}. However, such models are applicable only when the Compton amplification factor, $A \\sim 1$. Otherwise the Comptonized photons will impinge back on the disk and heat the disk to the corona temperature \\citep{Zdz96b}. Since for saturated Comptonization  the amplification factor $A \\approx T / T_s \\approx 10$, the uniform corona model is also ruled out. A possible geometry is that the inner region of an accretion disk is surrounded uniformly by a thermal plasma sphere. If the sphere is significantly larger than the soft photon producing disk, only a small fraction of the Comptonized photons will impinge back on the disk and energy balance can then be maintained. Interestingly, the fitted parameters suggest that the temperature of the accretion disk, which can be identified as the soft photon source temperature, $T_s \\sim 0.3$ keV, is expected for a disk  around a large $100 M_\\odot$ black hole. However, it is not clear whether such a large sphere would be dynamically stable or even realizable in the presence of a black hole.  Another radiative model for the source is that the thermal plasma Comptonizies photons that are internally produced by bremsstrahlung. In the standard accretion disk model \\citep{Sha73}, the disk is assumed to be effectively optical thick to absorption and hence the disk radiates as a black body locally. However, at high accretion rates, the inner region of the disk may become optically thin to absorption and the standard disk theory is no longer consistent. This transition occurs roughly when $\\tau_* \\equiv \\sqrt{\\tau_{es} \\tau_{abs}} \\approx 1$, which corresponds to a transition radius \\begin{equation} r_{tr} \\approx 40 \\; r_s \\; \\alpha^{34/93}\\; \\left (\\frac{M}{100 M_\\odot}\\right )^{2/93} \\left (\\frac{\\dot M}{\\dot M_{Edd}}\\right )^{64/93} \\end{equation} where $r_s \\equiv 2GM/c^2$ is the Schwarzschild radius, $M$ is the mass of the black hole, $\\dot M$ is the accretion rate, $\\dot M_{Edd} \\equiv L_{Edd}/0.1 c^2$, is the accretion rate corresponding to Eddington luminosity and  $\\alpha$ is the dimensionless viscosity parameter. For radii $< r_{tr}$, the disk cools inefficiently by bremsstrahlung Self Comptonization (BSC) and hence is hotter than a disk described by the standard theory. In this region, the temperature increases rapidly with decreasing radii. In certain cases, especially when the effectively optically thin region is an extended one, the temperature can vary from $\\approx 0.1$ keV at the transition radius to nearly $\\sim 50$ keV in the inner most radii. The resultant spectrum, which is a sum of local BSC spectra, can have a power-law form and such a scenario has been put forward as an alternate model to the standard unsaturated Comptonization interpretation of the high energy spectra of AGN \\citep{Mar90} and X-ray binaries \\citep{Mis98}. Since the temperature increases inwards, the high energy photons are preferentially produced in the inner region of the disk as compared to the lower energy ones. Thus the observed time-lag between different energy bands \\citep{Now99} can be explained as waves traveling from the outer to the inner regions of the disk \\citep{Mis00,Utt01}. For M82 X-1, the saturated spectrum could arise from such an effectively optically thin region. However, here the temperature variation is more modest, with a  maximum of $\\sim 2$ keV, indicating that the region is not an extended one. Detailed spectral prediction from such a model depends on the actual temperature profile, and the present data is not adequate to unambiguously constrain the physical characteristics of the region. Nevertheless, it is interesting to note that M82 X-1 is probably emitting close to its Eddington limit, even if it is a $100 M_\\odot$ black hole, and at such high accretion rates, the inner region of the standard disk is expected to become effectively optically thin.  The results obtained in this work need to be confirmed by simultaneous {\\it Chandra} and {\\it XMM-Newton} observations. It will also be interesting to find out using future {\\it XMM-Newton} observations, whether the source makes a transition to a state which has a regular steep X-ray spectrum. Such an analysis will also shed light on whether or not the temporal property of the source, especially the observed QPO, depends on its unusual spectrum. Search and analysis of other ULX that have similar hard spectra needs to be undertaken to verify whether M82 X-1 is indeed a unique source."
    },
    "0601/astro-ph0601055_arXiv.txt": {
        "abstract": "The \\HI\\ Parkes All Sky Survey (HIPASS) galaxy catalogue is cross-correlated with known low redshift, low column density (\\nhi\\ $<10^{15}$ \\cm) Lyman-$\\alpha$ absorbers from the literature. The redshift-space correlation is found to be similar in strength to HIPASS galaxy self-clustering (correlation length $s_{0,ag}=6\\pm4$ and $s_{0,gg}=3.1\\pm0.5$ \\h\\ Mpc respectively). In real-space the cross-correlation is stronger than the galaxy auto-correlation (correlation length $r_{0,ag}=7.2\\pm1.4$ and $r_{0,gg}=3.5\\pm0.7$ \\h\\ Mpc respectively) on scales from $1-10$ \\h\\ Mpc, ruling out the mini-halo model for the confinement \\lya\\ absorbers at the 99 percent confidence level. Provided that the cause of the strong cross-correlation is purely gravitational, the ratio of correlation lengths suggest that absorbers are embedded in dark matter haloes with masses log($M$/\\msun) $=14.2$ \\h, similar to those of galaxy groups. The flattening of the cross-correlation at separations less than $\\sim$ 600 \\h\\ kpc could correspond to the thickness of filaments in which absorbers are embedded. This work provides indirect statistical evidence for the notion that galaxy groups and large-scale filaments, particularly those that comprise gas-rich galaxies, are the dominant environments of low column density \\lya\\ absorbers at $z=0$. ",
        "introduction": "\\label{sec:intro}   At the present epoch less than 10 per cent of baryonic matter is located in galaxies (Cole et al. \\citeyear{Cole01}; Zwaan et al. \\citeyear{Zwaan03}). The majority of baryons therefore lie in the intergalactic medium (IGM), and knowledge of the distribution of these baryons is essential to our understanding of how matter is organised in the Universe. Lyman-$\\alpha$ (\\lya) absorption towards distant quasars reveals the presence of neutral hydrogen (\\HI) in the IGM. At low redshift, observations show that these low column density (\\nhi\\ $<10^{15}$ \\cm) absorbers exist in a variety of environments, including large-scale filaments \\citep[e.g.][]{LeBrun96, Penton02, Rosenberg03}, galaxy groups \\citep[e.g.][]{Lanzetta96, Bowen02}, and even voids \\citep{McLin02} and underdense regions \\citep{Grogin98}. A positive connection between absorbers and galaxies has been established by many studies \\citep{Morris93, Lanzetta95, Bowen96, Tripp98, Chen98, Chen01, Impey99, Ortiz-Gil99}, however most conclude that a one-to-one correspondence between absorbers and individual galaxies does not exist. \\cite{Cote05} and \\cite{Putman06} provide compelling evidence that absorbers trace the cosmic web, finding that absorbers at large radii do not follow galactic rotation curves. Thus the absorbers appear to lie in the same overdensities as the galaxies, but not in individual galaxy haloes.  Despite the absorbers' positive association with highly biased, strongly clustered objects, i.e. galaxies and large-scale structures, low-\\nhi\\ \\Lya\\ absorbers are among the weakest self-clustering objects in the low redshift Universe. The two-point correlation function in velocity separation, $\\Delta v$, shows mild absorber clustering for $\\Delta v\\lesssim500$ \\kms\\ \\citep{Tripp98, Impey99, Dobrzycki02, Penton02}, weaker than that of galaxy self-clustering \\citep[although, see ][]{Ulmer96}. Hydrodynamic simulations also measure weak clustering of small gas overdensities in velocity space \\citep{Dave03}, these overdensities, $0.5< \\log(\\rho_H/\\bar{\\rho}_H)<1.5$, correspond to \\lya\\ absorbers with $13\\lesssim\\log (\\nhi/$\\cm$)\\lesssim14$ \\citep{Dave99}. When compared with galaxies identified in the simulations, a filamentary structure emerges, where absorbers of increasing column density arise in closer proximity to galaxies.  The cross-correlation function of absorbers and galaxies can be used to measure the extent to which the two populations of objects are associated. At intermediate redshifts the strong cross-correlation between Lyman break galaxies and \\civ\\ absorbers is used to argue that these two populations are in fact the same object \\citep{Adelberger05}. The spatial co-incidence of \\civ\\ absorbers and Lyman break galaxies however can equally be interpreted as pre-enrichment from dwarf galaxies born from the collapse of $\\sim2\\sigma$ fluctuations at $6<z<12$ \\citep{Madau01,Porciani05}. The same degeneracy has the potential to plague the interpretation of the cross-correlation of absorbers and galaxies at $z=0$. How can we tell the difference between a strong cross-correlation signal due to absorbers arising in galaxy haloes and that due to a general overdensity of matter that surrounds galaxies? This issue is related to the claim that $10^{14}\\lesssim$ \\nhi\\ $\\lesssim10^{18}$ \\cm\\ absorbers arise within the haloes of galaxies, based on the anti-correlation between impact parameter and \\lya\\ absorber equivalent width \\citep{Lanzetta95, Chen98, Chen01}. Solid observational evidence \\citep{Bowen02, Cote05}, numerical simulations \\citep{Dave99} and analytic models \\citep{Lin00} have now resolved this degeneracy, showing that the anti-correlation between absorber-galaxy separation and density of absorbing material is a natural consequence of the overdensity of matter that surrounds galaxies, extending well beyond galaxy haloes. Since we know that absorbers and galaxies are largely not the same object, their cross-correlation signal can be used to measure the extent to which they are associated.  The cross-correlation signal can also, subject to interpretation, provide an indirect method to measure the bias of the absorbers, i.e. the relationship between \\HI\\ column density and underlying dark matter density.  In the Press--Schechter \\citep[PS,][]{Press74} formalism, the ratio of the bias of two populations of objects is equal to the ratio of the cross- to auto-correlation function \\citep[e.g.][]{Mo96,Mo02}.\\footnote{In this paper the function that compares the same population of objects is referred to as the auto-correlation, and the function that compares two populations is referred to as the cross-correlation.} Therefore, by (i) knowing the characteristic total (dark plus baryonic) halo mass of a galaxy population, (ii) measuring the galaxy auto-correlation function, and (iii) measuring the absorber-galaxy cross-correlation function, the mass of the haloes in which the absorbers are embedded can be inferred. This method has been successfully used by \\cite{Bouche04}. They cross-correlated \\mgii\\ absorbers with Luminous Red Galaxies (LRGs) to show that the \\mgii\\ absorbers are embedded in haloes with masses $\\sim2-8\\times10^{11}$\\msun, consistent with the expectation that $\\sim40$ per cent of \\mgii\\ absorbers arise in Damped \\lya\\ (DLA) systems \\citep{Rao06}, and that DLAs are expected to have total halo masses of that order. DLAs are gravitationally collapsed objects and at low redshift are consistent with the local galaxy population weighted by \\HI\\ cross section \\citep{Zwaan05}. Lower column density absorbers are not collapsed and may not follow the same bias trends. At intermediate redshifts the distribution of \\lya\\ absorbers reveal structures on scales up to 17 \\h\\ Mpc \\citep{Liske00}, at these redshifts the bias and clustering properties of \\lya\\ absorbers are well modelled \\citep{Cen98, McDonald02}. However the physics of the IGM increases in complexity towards $z=0$. As noted by \\cite{Dave03} the correlation function of gas overdensity no longer follows the correlation function of \\HI\\ optical depth at $z=0$, showing that the bias is not as predictable as at higher redshifts. Some models however place absorbers in individual mini-haloes, which do have a well defined bias.  Mini-halo models predict that absorbers arise in gravitationally confined gas within dark matter mini-haloes \\citep{Ikeuchi86, Rees86, Mo94}. Since mini-haloes are less massive than galaxy haloes, the absorber-galaxy cross-correlation is expected to be weaker than the galaxy auto-correlation. The mini-halo model is not an ideal description of \\lya\\ absorbers. Too many haloes per unit redshift are required to account for the observed density of absorption lines, since the mini-haloes have small spatial cross sections.  The absorber-galaxy cross-correlation function is instrumental in establishing whether absorbers are associated with individual mini-haloes, galaxies or large-scale structure. Although other statistics have been used by other authors, such as nearest neighbour and galaxy density distributions, the properties of the cross-correlation function can be related to the underlying dark matter density. The cross-correlation is also a clean statistic as it does not rely on assumptions about correlated pairs. \\cite{Morris93} is the only $z=0$ absorber-galaxy cross-correlation function published to-date that includes a full three-dimensional calculation. They used 17 absorbers along the 3C 273 line-of-sight and found that the absorber-galaxy cross-correlation is weaker than the galaxy auto-correlation on scales from $1-10$ $h_{80}^{-1}$ Mpc. \\cite{Mo94} explain this finding with a combination of $20-30$ per cent of absorbers arising in galaxies and the remainder arising in mini-haloes.  In this paper the cross-correlation function is calculated using 129 absorbers along 27 lines-of-sight from the literature and 5,317 galaxies from a blind 21-cm emission-line survey. The galaxy and absorber data are described in Section~\\ref{sec:data}. The details of the cross-correlation and halo-mass calculation methods are described in Section~\\ref{sec:method}. The subsequent results are outlined in Section~\\ref{sec:results}, including the redshift- and real-space correlation functions. The nature of the \\lya\\ absorber correlation signal and its relationship to galaxies, groups of galaxies and large-scale filaments is discussed in Section~\\ref{sec:dis}, with a final summary given in Section~\\ref{sec:summary}. To compare with correlation functions in the literature $H_0=100$ \\kms\\ Mpc$^{-1}$ is used throughout. ",
        "conclusions": "\\label{sec:dis}  The locations of low-\\nhi\\ \\Lya\\ absorbers and gas-rich galaxies have been cross-correlated at $z=0$, giving a redshift-space clustering signal similar in strength to that of galaxy self-clustering.  The projected cross-correlation removes the increased velocity dispersion along the line-of-sight and from it the real-space clustering can be derived. The cross-correlation function is a powerful statistic as it can be related to the underlying dark matter distribution. Here the absorber cross-correlation results are discussed in relation to mini-haloes, galaxy groups, large-scale filaments, numerical simulations and the environs of gas-rich galaxies.  \\subsection{Ruling out mini-haloes}  A typical mini-halo has a circular velocity, $v_c$, of 30 \\kms\\ \\citep[e.g.][]{Mo94}, which corresponds to a halo mass and bias of $6.8\\times10^9$ \\msun\\ and 0.73 respectively. Using Equation~\\ref{eqn:xiratio} and the auto-correlation result, such haloes would have a cross-correlation length of 3.3 \\h\\ Mpc. Comparing this correlation length with the absorber cross-correlation length ($r_{0,ag}=7.2\\pm1.4$ \\h\\ Mpc) rules out mini-haloes for the confinement of low-\\nhi\\ \\lya\\ absorbers at $z=0$ at the 99 percent confidence limit.  \\subsection{Comparison with galaxy groups}  The ratio of real-space clustering strengths inferred from the projected correlation functions imply that \\Lya\\ absorbers are embedded in haloes with masses in the range $13.6<$ log($M$/\\msun) $<14.5$ \\h. This halo mass range is of the same order as the median dynamical mass of galaxy groups in HIPASS, log($M$/\\msun) $\\sim13.8$ \\h\\ \\citep{Stevens05}. This similarity warrants a careful comparison between correlation properties of absorbers and groups.  The cross-correlation between groups and galaxies has been measured with both the Sloan and 2dF galaxy redshift surveys \\citep{Yang05}. The most massive groups (halo mass log($M$/\\msun) $\\le13.8$ \\h) display the most elongation along the $\\pi$ axis due to larger velocity dispersions, similar to that seen in the absorber cross-correlation. The shape of the group-galaxy cross-correlation resembles that of the absorber-galaxy cross-correlation at $\\sigma<2$ and $\\pi<10$ \\citep[][ figure 1 -- all groups and faintest galaxies]{Yang05}, although the group-galaxy correlation level is higher. No 2-D flattening signature of peculiar motions due to infall is seen in the absorber-galaxy cross-correlation, whereas this signature is apparent for groups. The projected correlation is similar in slope at $1<\\sigma<10$, but the flattening of the absorber-galaxy correlation, leads to a much stronger group-galaxy correlation at $\\sigma<1$.  The main difference between the correlation properties of absorbers and groups is their self-clustering. Absorber self-clustering (see Figure~\\ref{fig:xi_aa}) appears to be weak. The weak clustering could be due to the measurement technique, since absorber self-clustering measurements are restricted to the line-of-sight axis and the clustering signal is elongated, and therefore diluted along this axis. The cross-correlation measured along the redshift axis has almost three times the amplitude as that measured perpendicular to the line-of-sight (compare Figures~\\ref{fig:proj} and \\ref{fig:projpi}). Interestingly if the absorber auto-correlation (Figure~\\ref{fig:xi_aa}) were increased by a factor of three it would be comparable to the redshift-space galaxy auto-correlation (Figure~\\ref{fig:xi_s}). N-body simulations and analytic models can be used to measure absorber self-clustering in real-space. At least in the case of the warm-hot intergalactic medium (WHIM, slightly higher column densities than the absorbers considered in this paper) the real-space clustering of absorbers at $r=1$ \\h\\ Mpc could be a factor two less \\citep{Valageas02} or two more \\citep{Dave01} than that of galaxies.  Groups of galaxies on the other hand are strongly self-clustered, but not necessarily as strong as their constituent galaxies. More massive haloes are more strongly clustered and contain more galaxies per halo. Thus there are many more galaxies than haloes in strongly clustered environments, and a comparable number of galaxies and haloes in the field. Therefore, purely by relative numbers, galaxies in strongly clustered environments will dominate their correlation function, whereas the correlation signal of groups will be more evenly weighted. This weighting can lead to the self-clustering signal of groups being {\\it{weaker}} than that of their constituent galaxies. Naturally, group self-clustering depends on the mass of haloes considered. Padilla et al. (\\citeyear{Padilla04}) show that the redshift-space correlation length of groups identified in the 2dF galaxy redshift survey increases in amplitude from $s_0=5.5$ \\h\\ Mpc for the entire sample, to 12.6 \\h\\ Mpc for groups with a median mass of $1.1\\times10^{14}$ \\h\\ \\msun (compared with $s_0=6.82$ \\h\\ Mpc for 2dF galaxies from Hawkins et al. \\citeyear{Hawkins03}). The projected auto-correlation of this larger mass group sample is quite similar in amplitude and shape to the \\cite{Yang05} group cross-correlation function described above. Hence it appears that the self-clustering properties of groups are likely to be strong, but could be weaker than galaxies. On the other hand, the self-clustering properties of \\lya\\ absorbers are likely to be weak, but could be as strong as galaxies.  Absorbers, galaxies and groups each trace the underlying dark matter halo distribution in different ways. By comparing the clustering properties of absorbers to that of other objects, the dark matter haloes in which absorber reside can be revealed. As found by many authors (see Section~\\ref{sec:intro}), the possibility that \\lya\\ absorbers arise in individual galaxy haloes is unlikely. At a distance of 20 \\h\\ Mpc, HIPASS is sensitive to galaxies with \\mhi\\ $\\geq 2\\times10^{8}$ $h^{-2}_{100}$ \\msun. Considering only galaxies and absorbers within this distance, the nearest galaxy to 50 per cent of absorbers is greater than 1 \\h\\ Mpc in redshift-space, clearly not within the halo of an individual galaxy. Galaxy groups give a $\\xi(\\sigma,\\pi)$ diagram that is similar in shape to the absorber-galaxy cross-correlation. In addition, the ratio of real-space clustering lengths imply that \\lya\\ absorbers are embedded in haloes with masses similar to that of galaxy groups. These similarities provide some evidence for the notion that \\Lya\\ absorbers are associated with galaxy group environments.   \\subsection{Nature of the correlation signal compared with simulations}  Taking the projected absorber-galaxy cross-correlation (Figure~\\ref{fig:proj}) at face value, it can be concluded that stepping 1 to 10 \\h\\ Mpc away from an \\HI-selected galaxy, the normalised chance of finding an absorber is more likely than finding another galaxy. This likelihood is in addition to the fact that \\lya\\ absorbers are $\\sim$40 times more frequent than galaxies at zero redshift.\\footnote{The number density of absorbers with $13.1<$ log (\\nhi/\\cm) $<14.0$ is dN/dz=1.85 \\citep{Penton04}, the value for gas-rich galaxies is dN/dz=0.045 \\citep{Zwaan05}.} The exact covering factor of neutral gas 1 \\h\\ Mpc away from a galaxy however requires a value for the size of the absorbers, or an assessment of non-detections. At smaller separations the cross-correlation flattens with respect to the galaxy auto-correlation, indicating that at $<1$ \\h\\ Mpc the normalised chance of encountering a low-\\nhi\\ \\lya\\ absorber is less than that of another galaxy - perhaps giving way to higher column density systems. Flattening of the absorber auto-correlation function at a few hundred kiloparsecs is also seen in analytic models and N-body simulations of WHIM absorbers \\citep{Dave01, Valageas02}. This length ($\\sim$ 300 kpc) corresponds to the characteristic thickness of the WHIM filaments. A fairer comparison can be made with $\\xi(\\sigma,\\pi)$ diagrams for the cross-correlation of galaxies with low column density absorbers in \\cite{Dave99}. Their results suggest the correlation function turns over at small separations, but the interpretation is limited by the resolution of the simulations. The results presented here show that the cross-correlation flattens at separations less than 600 \\h\\ kpc. It is possible that this separation corresponds to a characteristic thickness of the large-scale structures in which low column density \\lya\\ absorbers are embedded.   Modelling the low redshift Universe is a complicated business. Ideally the results presented here should be compared with directly with $\\xi(\\sigma,\\pi)$ and $\\xi(r)$ from simulations and models at zero redshift. Future simulations may be able to distinguish the signature of absorbers embedded in different large-scale structures, for example that of groups or filaments. The essential effect that needs to be tested is whether the elongation along the velocity axis seen in the cross-correlation is due to absorbers that are embedded in a large gravitational potential. Another important insight offered by simulations and models is measuring absorber self-clustering in real-space. Given the large elongation of the absorber-galaxy cross-correlation in redshift-space, and the limitation that absorber self-clustering currently can only be measured along the line-of-sight, it is possible that the self-clustering strength of \\lya\\ absorbers has been understated.  \\subsection{Connection with gas-rich galaxies}  Although the HIPASS galaxies are simply used as test particles to represent the underlying density of matter, since they are gas-rich galaxies, gas-rich environments are deliberately selected. \\cite{Stevens04} have shown that the \\HI\\ content of galaxies in groups detected in HIPASS is no different to field galaxies, in contrast with optically selected groups, which are \\HI-deficient \\citep{Verdes01}. Thus the intragroup medium of \\HI-selected groups is expected to have a lower ionization fraction, leading to a higher number of \\lya\\ absorbers. This may explain why \\cite{Impey99} find that absorbers in the vicinity of the Virgo cluster lie preferentially in regions of intermediate (rather than high) galaxy density, since the outskirts of galaxy clusters are known to be in excess of \\HI-rich galaxies by a factor of three (Waugh et al. \\citeyear{Waugh02}).  \\cite{Stocke05} find that only 60 per cent of \\lya\\ absorbers with $\\log(\\nhi/$\\cm$)\\gtrsim 13.5$ (corresponding to detected \\ovi\\ absorption lines) arise in optically selected galaxy groups. Although the difference between optical and \\HI-selected groups in this case may be due to different sensitivity limits as \\cite{Stocke05} consider galaxy groups with more than one L* galaxy, whereas HIPASS is sensitive to much fainter galaxies (Doyle et al. \\citeyear{Doyle05}).  All galaxies in a representative sample of HIPASS galaxies (Meurer et al. \\citeyear{Meurer05}) have been detected in \\Ha\\ emission, and are therefore undergoing star formation. These results emphasize that interstellar hydrogen is a co-requisite for star formation and that HIPASS galaxies can be categorised as emission line galaxies. Interestingly, \\cite{Chen05} finds a stronger \\lya\\ cross-correlation for emission line galaxies, compared with absorption line galaxies, a result that suggests the environs of gas-rich galaxies are preferred by \\Lya\\ absorbers.  Despite the large number of absorbers and lines-of-sight, this study is limited by the shallow redshift coverage of HIPASS. Jackknife errors take into account cosmic variance within the HIPASS volume, but not outside the volume. Infrared galaxy counts show that the local area ($\\sim$200 \\h\\ Mpc in linear extent, to a distance of $\\sim$150 \\h\\ Mpc) around the south Galactic pole is underdense by 30 per cent \\citep{Frith03}.  \\cite{Lin00} suggest a minimum redshift path length, $\\Delta z$, of 10 for statistics that discriminate between absorbers arising in galactic haloes and those arising in the comic web. The literature sample used in this paper has $\\Delta z\\sim1$ only. HIPASS galaxies also have different clustering properties to optically selected galaxies \\citep{Meyer03}  For all these reasons it would certainly be worth repeating this study with other galaxy surveys.  Future observations designed to measure absorber-galaxy clustering to the same extent as seen here would require a survey of galaxies to projected distances of 10 \\h\\ Mpc, i.e. for nearby absorbers with velocities around 10,000 \\kms, a radius of almost 6 degrees would be required around each quasar, with a depth of at least 100 \\h\\ Mpc along the line-of-sight.     In this paper 129 low column density (\\nhi\\ $<10^{15}$ \\cm) \\lya\\ absorbers along 27 lines-of-sight have been cross-correlated with 5,317 gas-rich galaxies at $z=0$. A positive association between absorbers and galaxies is found to separations of at least 10 \\h\\ Mpc in redshift space. The redshift-space auto- and cross-correlation functions are found to be similar in strength, with correlation lengths $s_{0,gg}=3.1\\pm0.5$ and $s_{0,ag}=6\\pm4$ \\h\\ Mpc respectively. These results are contrary to that of \\cite{Morris93} based on 17 absorbers along a single line-of-sight, which found a weaker cross-correlation (the uncertainty on $s_{0,ag}$ is large however, so a weaker correlation for absorbers can only be ruled out at the 53 percent confidence limit). The cross-correlation is found to depend on the column density of the absorbers: higher column density systems (log (\\nhi/\\cm) $>13.24$) are more strongly correlated with galaxies, especially at projected separations less than 1 \\h\\ Mpc.  On a one-to-one basis, out to a distance of 20 \\h\\ Mpc, 50 per cent of \\lya\\ absorbers are located within 1 \\h\\ Mpc of a galaxy, and the other half are located $1-3$ \\h\\ Mpc from the nearest galaxy. These results agree with many other studies \\citep{Bowen96, Bowen02, Tripp98, Impey99, Penton02, Cote05} that the majority of \\lya\\ absorbers do not arise in the haloes of individual galaxies. Bearing this result in mind, the cross-correlation function can then be used to measure the extent to which absorbers and galaxies are related.  The amplitude of the absorber-galaxy cross-correlation depends on how it is integrated. The absorber-galaxy $\\xi(\\sigma,\\pi)$ diagram is very elongated along the line-of-sight. Integrating along the projected separation ($\\sigma$) axis produces auto- and cross-correlations with comparable amplitudes. Restricting this integral to 1 \\h\\ Mpc leads to a cross-correlation that is slightly weaker than the auto-correlation at small separations. Thus results from cylindrical-style cross-correlation functions \\citep{Chen05, Williger05} must be interpreted with caution.  A more conventional approach (as employed in galaxy clustering analysis) is to integrate along the line-of-sight ($\\pi$) axis to remove the increased velocity dispersion and from it derive the real-space clustering, and underlying dark matter distribution. The absorber-galaxy cross-correlation has been measured in real-space for the first time here. A correlation length of $r_{0,ag}=7.2\\pm1.4$ \\h\\ Mpc is found, compared with $r_{0,gg}=3.5\\pm0.7$ \\h\\ Mpc and slope of $\\gamma_r=1.9\\pm0.3$ for the galaxy auto-correlation. Mini-halo models on the other hand predict a weaker cross-correlation. The interpretation of these results relies on the assumption that the increased velocity dispersion seen in the cross-correlation is due to a larger gravitational potential. These results rule out mini-halo models for the confinement of low-\\nhi\\ \\lya\\ absorbers at $z=0$ at the 99 percent confidence limit. The ratio of real-space clustering strengths, together with the bias relation (Equation~\\ref{eqn:xiratio}) imply that \\Lya\\ absorbers are embedded in haloes with masses in the range $13.6<$ log($M$/\\msun) $<14.5$ \\h, similar to that of galaxy groups. The cross-correlation of massive groups with galaxies produces a $\\xi(\\sigma,\\pi)$ diagram similar in shape to that of the absorber-galaxy cross-correlation. Thus on average \\lya\\ absorbers tend to follow the same cosmic web of material that embeds both galaxies and groups. This picture is supported by careful analysis of individual systems \\citep{Bowen02}. It is also possible that absorbers arise in filaments of large-scale structure, although the cross-correlation amplitude of filaments has not been measured using observations nor predicted from models or simulations.  The flattening of the cross-correlation function at projected separation less than $\\sim$ 600 \\h\\ kpc could correspond to a characteristic thickness of the filaments in which \\lya\\ absorber are embedded. A large discrepancy exists between individual \\lya\\ absorber sizes derived from photo-ionization models using metal lines associated with each absorber \\citep[$10-70$ pc,][]{Rigby02, Tripp02} and their minimum transverse sizes derived from quasar line-of-sight pair analysis \\citep[of order 300 \\h\\ kpc,][]{Penton02, Rosenberg03}. In agreement with \\cite{Rosenberg03}, this discrepancy could be reconciled if the minimum transverse size corresponds to the thickness of the filament rather than the individual absorber size. Furthermore, \\cite{Penton02} identify a filament consisting of eight absorbers and five galaxies at least 20 $h_{70}^{-1}$ Mpc across and 1 $h_{70}^{-1}$ Mpc thick. This example fits with the $\\sim$ 600 \\h\\ kpc filament thickness inferred from the cross-correlation function.  The cross-correlation results presented in this paper provide indirect evidence that \\lya\\ absorbers arise preferentially in gas-rich galaxy groups and filaments. Over half the galaxies in HIPASS occur in groups \\citep{Stevens05} and gas with $\\log(\\nhi/$\\cm$)>13.2$ covers 50 per cent of galaxy filaments \\citep{Stocke05} -- a picture consistent with the ubiquitous nature of \\lya\\ absorbers."
    },
    "0601/gr-qc0601012_arXiv.txt": {
        "abstract": "A cosmological constant, $\\Lambda$, is the most natural candidate to explain the origin of the dark energy (DE) component in the Universe. However, due to experimental evidence that the equation of state (EOS) of the DE could be evolving with time/redshift (including the possibility that it might behave phantom-like near our time) has led theorists to emphasize that there might be a dynamical field (or some suitable combination of them) that could explain the behavior of the DE. While this is of course one possibility, here we show that there is no imperative need to invoke such dynamical fields and that a variable cosmological constant (including perhaps a variable Newton's constant too) may account in a natural way for all these features. ",
        "introduction": "The phenomenon of the accelerated expansion of the universe is presently one of the central issues of both observational and theoretical cosmology. A number of diverse cosmological observations\\,\\cite{Supernovae,WMAP03LSS} have by now established the accelerated nature of the present expansion of the universe and even provided additional information on the deceleration/acceleration transition and the redshift dependence of the expansion of the universe. From the theoretical side, the sole fact that the universe is presently accelerating, and may continue to do so, has triggered many studies. Some particularly interesting possibilities include  braneworld models of the late-time cosmic acceleration \\cite{braneworld}. The real theoretical challenge, however, lies in understanding the dynamics leading to the accelerated expansion of the universe. Despite the fact that many promising models have been proposed, the fundamental nature of the accelerating mechanism remains presently unknown. The attempts towards shedding some light on the cause of the acceleration of the universe employ a broad range of concepts in many theoretical frameworks. The most widely used and the conceptually simplest option assumes the existence of {\\em dark energy} (DE), the cosmic component with the negative pressure. Dark energy is a very useful concept since it encodes all our ignorance on the acceleration of the universe in a single cosmic component. Furthermore, DE can also be used as an effective description of other mechanisms of the acceleration of the universe\\,\\cite{Eqos}. A traditional candidate for the role of DE is the cosmological constant (CC), $\\Lambda$. In fact, the cosmological FRW model with cold dark matter and $\\Lambda$ as the DE component (the so-called $\\Lambda$CDM cosmology) fits the data reasonably well\\,\\cite{Supernovae,WMAP03LSS}. Theoretically, however, the CC as a DE candidate faces a formidable problem related to the very many ($55$ at least) orders of magnitude difference of its predicted value in quantum field theory (QFT) and the observed value\\,\\cite{CCP}. This huge discrepancy leads indeed to the cosmological challenge of the millennium and calls for a more profound treatment of the CC problem\\,\\cite{Paddy}. The many inconsistencies related to the ``$\\Lambda$ conundrum'' have led to the development of the dynamical DE models. The tacit assumption of all these models is that $\\Lambda$ does not contribute to the DE or that it vanishes, which merely circumvents the CC problem. The dynamical DE models, however, incorporate the advantage that they approach the modeling of the mysterious dark energy component in a more general way, allowing its properties to vary with the expansion. Notwithstanding this welcome feature, the necessity of a (severe) fine-tuning of the parameters of these models in order that the measured value of the DE coincides with the predicted value does persist and it is in no way less worrisome than in the CC case. The list of dynamical DE models comprise, among others, {\\em quintessence} \\cite{Quintessence}, {\\em phantom energy} \\cite{phantom},{\\em Chaplygin gas}\\,\\cite{Chaplygin} etc\\,\\cite{CCP,Paddy}. Very often these models of DE are realized in terms of dynamical scalar field(s). These fields were actually introduced on more or less phenomenological grounds to deal with the CC problem long ago \\,\\cite{Dolgov,PSW}. In particular, in the popular quintessence approach the scalar fields and their parameters/scales are totally unrelated to the known particle physics fields, and as a consequence an obvious connection to fundamental physics is lacking. Not only so, in all these models the scalar field potentials just concocted to describe an acceptable form of DE are non-renormalizable by power counting and, what is worse, they are essentially field-theoretically unmotivated.  In this paper we take the point of view that these DE models are not more fundamental than the original $\\Lambda$ model. However, we exploit the most useful property of the former, namely the dynamical character of the DE. Therefore, we generalize the CC concept allowing the variability of the $\\Lambda$ term (and possibly also of the gravitational coupling $G$) with the cosmic time, both assumptions being perfectly compatible with the cosmological principle and can be realized, as we shall see, in a fully covariant way. The support for such a generalization of the $\\Lambda$ term comes from QFT on curved space-time \\cite{JHEPCC1,Babic} and/or quantum gravity approaches \\cite{Reuter}. Here, however, we do not derive the variability of  $\\Lambda$ and $G$ from these models, but discuss the general implications of this variability. Our conclusions will therefore be valid for any specific model of the mentioned type. Let us emphasize that this approach embodies many virtues: i) the problem of the CC is dealt with instead of circumventing it; ii) the variability of $\\Lambda$ may shed some light on the observed value of $\\Lambda$; iii) a variable $\\Lambda$ is indeed a type of dynamical DE, and will be handled here in the language of the effective DE picture. We will show in a precise way how a cosmology with variable parameters $(\\Lambda,G)$ can mimic a dynamical DE model and produce a non-trivial {\\em effective} equation of state (EOS), $p_D=w_{\\rm eff}\\,\\rho_D$, beyond the naively expected one for the $\\Lambda$ term ($w_{\\Lambda}=-1$). Since the EOS parametrization is widely used in all planned experiments aiming at a precise study of the DE\\,\\cite{Eqos}, such correspondence can be useful and it may even unveil some unexpected limitations of the EOS method for characterizing the ultimate nature of the DE. ",
        "conclusions": "We have shown that a model with variable $\\Lambda$ (and possibly of $G$) generally leads to a non-trivial effective EOS, thus mimicking a canonical dynamical DE model (i.e. one with conserved DE density). That EOS can effectively appear as quintessence and even as phantom energy. Moreover, we have proven that there \\textit{always} exists a transition point $z^{*}$ near $z=0$ where $w_{\\rm eff}(z^{*})=-1$. If this point lies in our recent past (as illustrated in Fig.\\,\\ref{plot}a) there could have been a recent transition into an (effective) phantom regime $w_{\\rm eff}(z)<-1$, as suggested by several analysis of the data. If, however, experiments would unambiguously indicate $w_{\\rm eff}(z)>-1$ we could equally well interpret this in our framework, for $z^{*}$ could just be in our immediate future. We conclude that variable $(\\rho_{\\Lambda},G)$ models may account for the observed evolution of the DE, without need of invoking any combination of fundamental quintessence and phantom fields. The eventual determination of an empirical EOS for the DE in the next generation of precision cosmology experiments should keep in mind this possibility.  \\textit{Acknowledgements}. This work has been supported in part by MEC and FEDER under project 2004-04582-C02-01, and also by DURSI Generalitat de Catalunya under project 2005SGR00564. The work of HS is financed by the Secretaria de Estado de Universidades e Investigaci\\'on of the Ministerio de Educaci\\'on y Ciencia of Spain. HS thanks the Dep. ECM of the Univ. of Barcelona for the hospitality."
    },
    "0601/astro-ph0601263_arXiv.txt": {
        "abstract": "{ We examine the spatial distribution of brown dwarfs produced by the decay of small-$N$ stellar systems as expected from the embryo ejection scenario.  We model a cluster of several hundred stars grouped into 'cores' of a few stars/brown dwarfs.  These cores decay, preferentially ejecting their lowest-mass members.  Brown dwarfs are found to have a wider spatial distribution than stars, however once the effects of limited survey areas and unresolved binaries are taken into account it can be difficult to distinguish between clusters with many or no ejections.  A large difference between the distributions probably indicates that ejections have occurred, however similar distributions sometimes arise even with ejections.  Thus the spatial distribution of brown dwarfs is not necessarily a good discriminator between ejection and non-ejection scenarios. ",
        "introduction": "Brown dwarfs are observed to constitute some 15--25\\% of the objects in young star forming regions (e.g. Brice\\~no et al. 2002; Muench et al. 2003; Luhman et al. 2003; Luhman 2005).  However, the formation mechanism(s) of brown dwarfs are currently unclear.  The two most popular models are that (a) brown dwarfs form like stars but in very low-mass cores (e.g. Greaves this volume; Padoan this volume), or (b) brown dwarfs are stellar embryos that are ejected before than can accrete sufficient mass to become stars (Reipurth \\& Clarke 2001,2003). (Also see Whitworth \\& Goodwin (2005) and Whitworth \\& Goodwin (this volume) for a review of other possible mechanisms).  Probably the most popular model for brown dwarf formation (at least among theorists) is the ejection scenario (Reipurth \\& Clarke 2001).  This model appears to be supported by simulations of star formation in turbulent cores (e.g. Bate et al. 2002, 2003; Delgado Donate et al. 2004; Goodwin et al. 2004a,b).  The ejection scenario has three significant problems.  Firstly, it is not clear if ejected brown dwarfs are able to retain the significant discs which are observed around at least some young brown dwarfs (e.g. Jayawardhana et al. 2003).  Secondly, whilst ejections do produce some brown dwarf-brown dwarf binary systems (e.g. Bate et al. 2002) a large population of brown dwarf-brown dwarf binaries would be difficult to explain via the ejection scenario (this problem has recently become more acute with evidence that the brown dwarf binary fraction may be significantly higher than previously thought (Pinfield et al. 2003; Jeffries \\& Maxted this volume).  Finally, brown dwarfs are ejected with velocities of order 1 km s$^{-1}$ and it may be expected that brown dwarfs are more widely distributed than stars, in contrast to observations (Brice\\~no et al. 2002; see also Luhman 2005).  In this paper we concentrate on the final problem - the spatial distribution of brown dwarfs.  In the ejection scenario, cores form several objects in an unstable high-order multiple system.  Dynamical interactions typically eject the lowest-mass members of the system until a stable binary or hierarchical multiple remains.  We perform $N$-body simulations of a cluster of decaying small-$N$ `cores' to examine if the distributions of brown dwarfs and stars are different or if the velocity dispersion between cores effectively conceals the ejection of brown dwarfs and low-mass stars. ",
        "conclusions": "When ejections are not important ($N_* = 1$) the spatial distributions of stars and brown dwarfs are (unsurprisingly) very similar.  When ejections become important ($N_*=4$) then the spatial distributions can show significant differences.  However, these differences can disappear altogether in some clusters depending on the exact details of the initial clustering.  In addition, at the young age of most well-studied young clusters a significant difference in the spacial distributions is seen in only 1 out of 5 simulations.  Thus if brown dwarfs and stars have different spatial distributions it is probably a signature of ejections, however the lack of a difference does not necessarily exclude the ejection scenario.  We suggest that binarity is a far stronger discriminator between models since (as yet) there is no way in which to make significant numbers of brown dwarf-brown dwarf binaries from ejections.  However, we are currently investigating if a more sophisticated statistical analysis of positions may yield a more robust discriminator.  The discovery of many more brown dwarfs in Taurus extending over a wide area (Guieu et al. this volume) does suggest that ejections may have been responsible for at least some of the Taurus brown dwarfs."
    },
    "0601/astro-ph0601367.txt": {
        "abstract": "{} {We present the second campaign of the ESO/GOODS program of spectroscopy of faint galaxies in the GOODS-South field.} {Objects were selected as candidates for VLT/FORS2 observations primarily based on the expectation that the detection and measurement of their spectral features would benefit from the high throughput and spectral resolution of FORS2. The reliability of the redshift estimates is assessed using the redshift-magnitude and color-redshift diagrams and comparing the results with public data.} {807 spectra of 652 individual targets have been obtained in service mode with the FORS2 spectrograph at the ESO/VLT, providing 501 redshift determinations. The typical redshift uncertainty is estimated to be $\\sigma_z \\simeq 0.001$. Galaxies have been selected adopting three different color criteria and using the photometric redshifts. The resulting redshift distribution typically spans two redshift domains: from z=0.5 to 2 and z=3.5 to 6.2. In particular, 94 $B_{435}$-,$V_{606}$-,$i_{775}$-\"dropout\" Lyman break galaxies have been observed, yielding redshifts for 65 objects in the interval 3.4$<$z$<$6.2. Three sources have been serendipitously discovered in the redshift interval 4.8$<$z$<$5.5. Together with the previous release, 930 sources have now been observed and 724 redshift determinations have been carried out. The reduced spectra and the derived redshifts are released to the community through the ESO web page $\\it{http://www.eso.org/science/goods/}$ \\thanks{The catalog (Table~\\ref{tab:tblspec}) is available in electronic form at the CDS via anonymous ftp to cdsarc.u-strasbg.fr (130.79.128.5) or via http://cdsweb.u-strasbg.fr/cgi-bin/qcat?J/A+A/}. Large scale structures are clearly detected at $z \\simeq 0.666, 0.734, 1.096, 1.221, 1.300$, and $1.614$. A sample of 34 sources with tilted [O\\,{\\sc ii}]3727 emission has been identified, 32 of them in the redshift range 0.9$<$z$<$1.5. } {} ",
        "introduction": "The Great Observatories Origins Deep Survey (GOODS) is a public, multi-facility project that aims to answer some of the most profound questions in cosmology: how did galaxies form and assemble their stellar mass? When was the morphological differentiation of galaxies established and how did the Hubble Sequence form? How did Active Galactic Nuclei (AGN) form and evolve, and what role do they play in galaxy evolution? How much do galaxies and AGN contribute to the extragalactic background light? Is the expansion of the universe dominated by a cosmological constant? A project of this scope requires large and coordinated efforts from many facilities, pushed to their limits, to collect a database of sufficient quality and size for the task at hand. It also requires that the data be readily available to the worldwide community for independent analysis, verification, and follow-up.  The program targets two carefully selected fields, the Hubble Deep Field North (HDF-N) and the Chandra Deep Field South (CDF-S), with three NASA Great Observatories (HST, Spitzer and Chandra), ESA's XMM-Newton, and a wide variety of ground-based facilities. The area common to all the observing programs is 320 arcmin$^2$, equally divided between the North and South fields. For an overview of GOODS, see \\cite{dick03}, \\cite{renz03} and \\cite{giava04a}. In the last five years the CDF-S has been the target of several spectroscopic campaigns (\\cite{crist00}, \\cite{croom01}, \\cite{bunk03}, \\cite{stan04}, \\cite{stro04}, \\cite{vanderwel04}, \\cite{dick04}, \\cite{szo04}, \\cite{fevre05}, \\cite{vanz05}).  This is the second paper in a series presenting the results of the GOODS spectroscopic program carried out with the VLT/FORS2 spectrograph. For a full description of its aims we refer to the first paper (\\cite{vanz05}, RUN1 hereafter).  Here we recall that the ESO/GOODS spectroscopic program is designed to observe all galaxies for which VLT optical spectroscopy is likely to allow the redshift determination. The program makes full use of the VLT instrument capabilities (FORS2 and VIMOS), matching targets to instrument and disperser combinations in order to maximize the effectiveness of the observations. The magnitude limits and selection bandpasses to some extent depend on the instrumental setup being used. The aim is to reach mag~$\\sim24-25$ with adequate S/N, with this limiting magnitude being in the B band for objects observed with the VIMOS LR-Blue grism, in the V band for those observed in the VIMOS LR-Red grism, and in the z band for the objects observed with FORS2.  The second FORS2 spectroscopic campaign (17 masks, RUN2 hereafter) in the Chandra Deep Field South, was carried out in the period fall 2003 - early 2004 in service mode. New FORS2 observations were performed in December 2004 (6 masks, RUN3) mainly focused on color-selected Lyman break ``dropout'' targets and 5 more masks will be completed before February 2006 (RUN4) mainly dedicated to sources detected at 24$\\mu$m with the {\\it Spitzer Space Telescope} MIPS instrument. These data will be described in a forthcoming paper. The VIMOS spectroscopic survey in the GOODS-S field is started and will produce hundreds of redshift determinations, mainly in the redshift range 0$<$z$\\leq$3.5.  The paper is organized as follows. Sect. 2 describes the target selection, while Sect. 3 describes the observations and data reductions. The redshift determination is presented in Sect. 4. In Sect. 5 we discuss the data and in Sect. 6 the conclusions are presented. Throughout this paper the magnitudes are given in the AB system (\\cite{oke77}) (AB~$\\equiv 31.4 - 2.5\\log\\langle f_\\nu / \\mathrm{nJy} \\rangle$), and the ACS F435W, F606W, F775W, and F850LP filters are designated hereafter as $B_{435}$, $V_{606}$, $i_{775}$ and $z_{850}$, respectively. We assume a cosmology with $\\Omega_{\\rm tot}, \\Omega_M, \\Omega_\\Lambda = 1.0, 0.3, 0.7$ and $H_0 = 70$~km~s$^{-1}$~Mpc$^{-1}$.  %__________________________________________________________________ ",
        "conclusions": "As a part of the Great Observatories Origins Deep Survey, a large sample of galaxies in the Chandra Deep Field South has been spectroscopically targeted. After the RUN1 (\\cite{vanz05}) and RUN2 (present work) a total of 930 objects with $z_{850} \\mincir 26.8$ have been observed with the FORS2 spectrograph at the ESO VLT providing 724 redshift determinations. From a variety of diagnostics the measurement of the redshifts appears to be precise (with a typical $\\sigma_z \\simeq 0.001$) and reliable. The reduced spectra and the derived redshifts are released to the community ($\\it{http://www.eso.org/science/goods/}$). They constitute an essential contribution to reach the scientific goals of GOODS, providing the time coordinate needed to delineate the evolution of galaxy masses, morphologies, and star formation, calibrating the photometric redshifts that can be derived from the imaging data at 0.36-8$\\mu$m and enabling detailed studies of the physical diagnostics for galaxies in the GOODS field."
    },
    "0601/astro-ph0601505_arXiv.txt": {
        "abstract": "We have gained intermediate--resolution spectroscopy with the FORS instruments of the Very Large Telescope and high--resolution imaging with the Advanced Camera for Surveys aboard \\emph{HST} of a sample of 220 distant field spiral galaxies within the FORS Deep Field and William Herschel Deep Field. Spatially resolved rotation curves were extracted and fitted with synthetic velocity fields that take into account all geometric and observational effects, like blurring due to the slit width and seeing influence. Using these fits, the maximum rotation velocity $\\vm$ could be determined for 124 galaxies that cover the redshift range $0.1<z<1.0$ and comprise a variety of morphologies from early--type spirals to very late--types and irregulars. The luminosity--rotation velocity distribution of this sample, which represents an average look-back time of $\\sim$5 Gyr, is offset from the Tully-Fisher relation (TFR) of local low--mass spirals, whereas the distant high--mass spirals are compatible with the local TFR. Taking the magnitude--limited character of our sample into account, we show that the slope of the local and the intermediate-$z$ TFR would be in compliance if its scatter decreased by more than a factor of 3 between $z \\approx 0.5$ and $z \\approx 0$. Accepting this large evolution of the TFR scatter, we hence find no strong evidence for a mass- or luminosity-dependent evolution of disk galaxies. On the other hand, we derive stellar $M/L$ ratios that indicate a luminosity-dependent evolution in the sense that distant low--luminosity disks have much lower $M/L$ ratios than their local counterparts, while high--luminosity disks barely evolved in $M/L$ over the covered redshift range. This could be the manifestation of the ``downsizing'' effect, i.e. the succesive shift of the peak of star formation from high--mass to low--mass galaxies towards lower redshifts. This trend might be canceled out in the TF diagram due to the simultaneous evolution of multiple parameters. We also estimate the ratios between stellar and total masses, finding that these remained constant since $z=1$, as would be expected in the context of hierarchically growing structure. ",
        "introduction": "The concept of Cold Dark Matter (CDM) and its prediction of hierarchical structure growth has become an astrophysical paradigm in the last decade. Observations of the Cosmic Microwave Background \\citep[e.g.][]{Spe03} or the Large Scale Structure \\citep[e.g.][]{Bah99} strongly support the idea that the vast majority of matter is non--baryonic, non--luminous and interacts only gravitationally. In this picture, gas and stars are embedded in Dark Matter halos, and low--mass systems were the first virialized structures in the early cosmic stages, followed by an epoch of accretion and merger events during which larger systems were successively built up. Disks are destroyed during the frequent mergers of similar-mass galaxies at $z > 1$ but can regrow via accretion events at lower redshifts \\citep[e.g.][]{Aba03}.   An important tool for the quantitative test of the predictions from numerical simulations based on the hierarchical scenario are scaling relations that link galaxy parameters. Within the last years, several studies focused on the evolution of such relations up to redshifts $z \\approx 1$~-- corresponding to more than half the age of the universe~-- or even beyond. Many analyses utilized the Tully--Fisher Relation (TFR) between the maximum rotation velocity $\\vm$ and the luminosity or stellar mass \\citep[e.g.][]{Vog01, Zie02, Mil03, Boe04, Bam06, Con05, Flo06, Nak06, Wei06, Chi07, Kas07}.  Studies focusing on the stellar mass TFR consistently have found no evidence for evolution over the past $\\sim$\\,7 Gyr \\citep[e.g.][]{Con05, Kas07}. On the other hand, analyses exploiting the rest-frame $B$-band TFR~--- which is less sensitive to the stellar mass than to the fraction of young stars~--- yielded quite discrepant results. This is somewhat suprising since most recent TFR studies comprise $\\sim$\\,100\\,objects or more. Some authors derived much higher luminosities of spiral galaxies in the past \\citep[e.g.][]{Boe04, Bam06, Chi07} while others find only a very modest evolution in luminosity \\citep[e.g.][]{Vog01, Flo06}. There also still is debate on whether the $B$-band TFR changes its slope with redshift which would imply an evolution depending on $\\vm$ and hence total galaxy mass. In \\citet[][hereafter B04]{Boe04}, we found that low-mass spirals at $\\langle z \\rangle \\approx 0.5$ were brighter by up to several magnitudes than in the local universe, while high-mass spirals barely evolved in $M_B$ at given $\\vm$. To the contrary, \\citet{Wei06} observe a stronger brightening in the high-mass than in the low-mass regime towards larger redshifts. From the theoretical side, the situation is similarly unclear: in some simulations, the $B$-band TFR slope remains constant with lookback time and only the magnitude zero point changes \\citep[e.g.][]{Ste99, Por07}, while also an evolution in slope is predicted by other authors \\citep[e.g.][]{Boi01, Fer01}.  A growing number of observational studies based on other methods than the TFR have pointed towards a mass-dependent evolution of distant galaxies, either in terms of their colors \\citep[e.g.][]{Kod04}, mass--to--light ratios \\citep[e.g.][]{vdW05}, or average stellar ages \\citep[e.g.][]{Fer04}. These results indicate that the global stellar populations of distant high--mass galaxies are on average older than those of distant low--mass galaxies, similar to what has been found in the local universe \\citep[e.g.][]{BdJ00}. Since small galaxies are understood as ancient building blocks within the framework of hierarchical structure growth, and in turn should have an older stellar content than larger systems formed more recently, these observations are at variance with the straightforward expectation. They can be interpreted such that the peak of star formation shifts from high--mass to low--mass galaxies with growing cosmic age, a phenomenon that was termed ``downsizing'' by \\citet{Cow96}. It is however a non-trivial question what kind of imprint ``downsizing'' has on the TFR since the luminosity-$\\vm$ diagram most likely probes the evolution of various galaxy properties simultaneously, including stellar $M/L$ ratio, gas mass fraction, dust content etc.  In this paper, we report on an extensive observational study of spiral galaxy evolution over the last 7\\,Gyr. The new sample holds 124 intermediate-$z$ disk galaxies with reliable $\\vm$ measurements, roughly doubling the size of our previous data set from B04. The analysis given therein was limited to ground-based imaging, while we here can rely on our high-resolution imaging performed with the Advanced Camera for Surveys of the \\emph{HST}. This allows a much more accurate determination of disk inclination angles, scale lengths etc., and in turn a more robust estimate of the maximum rotation velocities. Extending the approch used in B04, we will analyse the evolution not only of the TFR but also of the stellar mass-to-light ratios and the stellar mass fractions. We stress that this is the first refereed publication on the total sample; only some brief excerpts were shown in \\citet[][]{Zie06}.  The structure of this paper is as follows. In \\S2, we describe the target selection, spectroscopic data reduction and redshift determination. \\S3 briefly introduces the \\emph{HST/ACS} imaging and morphological analysis. Details on the derivation of the intrinsic maximum rotation velocities and luminosities are given in \\S4. We present and discuss our results in \\S5, followed by a short summary in \\S6. Throughout this article, we will assume a flat cosmology with $\\Omega_\\lambda=0.7$, $\\Omega_m=0.3$ and $H_0=70$\\,km\\,s$^{-1}$\\,Mpc$^{-1}$. ",
        "conclusions": "We have used VLT/FORS spectroscopy and HST/ACS imaging to construct a sample of 220 field spiral galaxies up to redshift $z=1$. Spatially resolved rotation curves were extracted and fitted with synthetic velocity fields that take into account all geometric and observational effects, like blurring due to the slit width and seeing influence. Using these fits, the maximum rotation velocity $\\vm$ could be determined for 124 galaxies with an average look-back time of $\\sim$5 Gyr. We find that this sample is offset from the local Tully-Fisher Relation. The distant low--mass galaxies are more luminous at given $\\vm$ than their local counterparts, whereas the distant high--mass spirals are compatible with  the local TFR. Taking the magnitude limit of our sample into account, we show that the slope of local and distant relation would be in compliance if the TFR scatter decreased by more than a factor of 3 between $z \\approx 0.5$ and $z \\approx 0$. On the other hand, the $M/L$ ratios indicate a luminosity-dependent evolution in the sense that distant low--luminosity disks have much lower $M/L$ ratios than their local counterparts, while high--luminosity disks barely evolved in $M/L$ over the covered redshift range. This could be interpreted as an indication of the ``downsizing'' effect, i.e. the succesive shift of star formation from high--mass to low--mass galaxies towards lower redshifts. In terms of the Dark Matter Halos, we find evidence for a hierarchical evolution, since the fraction between stellar and total mass remained roughly constant since $z=1$."
    },
    "0601/astro-ph0601219_arXiv.txt": {
        "abstract": "Gamma ray bursts are excellent candidates to constrain physical  models which break the Lorentz symmetry. We consider deformed dispersion relations which break boost invariance and lead to an energy-dependent speed of light. In these models, simultaneously emitted photons from cosmological sources reach Earth with a spectral time delay that depends on the symmetry breaking scale. We estimate the possible bounds which can be obtained by comparing the spectral time delays with the time resolution of  available telescopes. We discuss the best strategy to reach the strongest bounds. We compute the probability of detecting bursts that improve the current bounds. The results are encouraging, depending on the model, it is possible to build a detector that within several years will improve the present limits of 0.015 $m_{pl}$. ",
        "introduction": "One of the open questions of high energy physics is how to unify gravity with quantum physics. A lot of effort has been devoted to develop a theory of quantum gravity. This theory is likely to require a drastic modification of our current understanding of the space-time.  At present there are two formal mathematical approaches : loop quantum gravity and superstring theory. Whatever might be the right description of the space-time at very short scales, there are some likely physical manifestations. It has been suggested, for instance, that such a theory would break what we believe to be basic symmetries of nature. In \\cite{Wheeler:1957mu,Hawking:1978}, it was shown that Einstein Lagrangian allows for large fluctuations of the metric and the topology of the space-time on scales of order of the Planck length, creating a foam-like structure at these scales. It has been proposed that the propagation of particles in a foamy space-time is strongly affected on short scales. The medium responds differently depending on the energy of the particle, in analogy with the propagation through a conventional electromagnetic plasma \\cite{Drummond:1979pp,Latorre:1994}. Thus space-time might exhibit a non-trivial dispersion relation in vacuum, violating therefore Lorentz invariance.  There are many different ways of breaking Lorentz symmetry; a background tensor field, like a magnetic field, breaks the vacuum rotational invariance, for instance. However, it has been shown that there are 46 different ways in which the standard model Lagrangian can be modified while remaining renormalizable, invariant under $SU(3) \\times SU(2) \\times U(1)$ and rotationally and translationally invariant in a preferred frame \\cite{Coleman:1998ti}. Among other effects, these terms cause the velocity of light to differ from the maximum attainable velocity of a particle, therefore changing the kinematics of particle decays. Modifying the dispersion relations of photons and electrons allows for new QED vertex interactions like photon splitting in vacuum, vacuum \\v{C}erenkov effect for electrons, photon decay, electron-positron annihilation to a single photon, etc. \\cite{Jacobson:2002hd,mewes}.  In this work we consider only rotationally invariant deformations of the photon dispersion relation which produce an energy-dependent speed of light. If this effect is present in nature, it has to be absent below some energy scale, $\\xi \\, m_{pl}$, (where $m_{pl}$ is the Planck energy and $\\xi$ a dimensionless constant), high enough to have been undetected so far.  Traditionally this energy is believed to be the Planck energy which, at first, seems to make hopeless any experimental attempt to test these models. However, in 1997 Amelino-Camelia et al. \\cite{Amelino-Camelia:1997gz} suggested that such models can be explored by studying the propagation of photons emitted from a distance source like a gamma ray burst (GRB) GRBs are short and intense pulses of soft gamma rays that arrive from cosmological distances from random directions in the sky. The bursts last from a fraction of a second to several hundred seconds. Most GRBs are narrowly beamed with typical energies around $10^{51}$ ergs, making them comparable to supernovae (for a recent review see \\cite{Piran:2004ba,Zhang:2003uk}). Because of the large distances traveled by the photons, these bursts are valuable tools to explore energies far beyond the reach of any laboratory on Earth.  If one considers photons with energies much smaller than the symmetry breaking scale, it is possible to expand the dispersion relation in powers of $E /\\xi m_{pl}$. The first correction produces a tiny departure from the Lorentz invariant (LI) equations and we expect that the low energy limit of the deformed dispersion relation can be written generically as: \\be E^2 - p^2  c^2 \\simeq  \\epsilon \\, E^2 \\left( {E \\over \\xi m_{pl} }\\right)^n , \\label{mod-dis-rel} \\ee where $\\epsilon = \\pm 1$ takes into account the possibility of having either infraluminal or superluminal motion, the latter appearing in some models of quantum loop gravity \\cite{Gambini:1998it,Alfaro:1999wd}. Photons simultaneously emitted from a GRB with different energies will travel at different speeds, and therefore will show on Earth a time delay.  The goal of this paper is to explore these high energy Lorentz violation models by studying such time delays. We analyze the potential of detecting observational consequences of a modified dispersion relation like Eq. \\ref{mod-dis-rel}.  The paper is organized as follows; in section \\ref{stand-th} we review cosmological photon propagation in the LI theory and show how these results are modified when a non LI term is introduced. We compute the travel time of cosmological photons and, as a check, compare it with the travel time obtained in the Newtonian approximation. In section \\ref{det-photons} we turn to the observations and show how such models can be tested using GRB observations. We compare our method with previous works in section \\ref{comparacion} and finally conclude in section \\ref{conclusion}. ",
        "conclusions": "\\label{conclusion}  Our goal was to explore the potential of GRBs to set bounds on Lorentz violation and to find optimal techniques to do so.  We modified the dispersion relation of photons by adding an extra term proportional to the photon momentum to the power $n+2$. We have shown that in models with $n=1$ it is possible to explore energies which are close to the Planck energy. When $n=2$, the energies explored are smaller, around $10^7 $ GeV.  These bounds are idealized and they do not take into account experimental limitations or the intrinsic time-structure of the $\\gamma$-ray emission. They should serve as theoretical estimations of what can be achieved. The methodology we use here can be used to design future optimal experiments for observing this effect (or setting bounds on it).  We have modelled the burst high energy emission with a power law spectrum $E^{-\\beta}$ with $\\beta \\geq 2$. This fits well the time integrated emission of most of the bursts. We found two non intuitive results: (i) The optimal redshift to set the strongest bound is less than 1.  (ii) For $n=1$, low energy, rather than high energy emission is preferred. Both results are counter-intuitive since the Lorentz violation delay increases with the distance and with the energy. However, distance or observations at high energies (where the flux is lower) dilute the photons reducing the temporal resolution achieved on Earth. It turns out that this is the dominant effect.  In the models with $n=2$, going to higher energies always improves the bounds. Here the situation will be remarkably changed when the spatial observatory GLAST will become operational.  We have also investigated the probability of improving the current experimental bounds, given a phenomenological luminosity and space distribution of bursts.  As we are discussing idealized bounds, this probability should only be trusted up to an order of magnitude."
    },
    "0601/astro-ph0601169_arXiv.txt": {
        "abstract": "{ We present the first mid-IR long baseline interferometric observations of the circumstellar matter around binary post-AGB stars. Two objects, \\object{SX Cen} and \\object{HD 52961}, were observed using the VLTI/MIDI instrument during Science Demonstration Time. Both objects are known binaries for which a stable circumbinary disc is proposed to explain the SED characteristics. This is corroborated by our N-band spectrum showing a crystallinity fraction of more than 50\\,\\% for both objects, pointing to a stable environment where dust processing can occur. Surprisingly, the dust surrounding SX\\,Cen is not resolved in the interferometric observations providing an upper limit of 11 mas (or 18 AU at the distance of this object) on the diameter of the dust emission. This confirms the very compact nature of its circumstellar environment. The dust emission around \\object{HD\\,52961} originates from a very small but resolved region, estimated to be $\\sim$35 mas at 8\\,$\\mu$m and $\\sim$55 mas at 13\\,$\\mu$m. These results confirm the disc interpretation of the SED of both stars. In \\object{HD\\,52961}, the dust is not homogeneous in its chemical composition: the crystallinity is clearly concentrated in the hotter inner region.  Whether this is a result of the formation process of the disc, or due to annealing during the long storage time in the disc is not clear. ",
        "introduction": "The fast stellar evolution connecting the Asymptotic Giant Branch (AGB) to the Planetary Nebulae (PNe) phase is still poorly understood \\citep[e.g.][]{Vanwinckel_2003}.  Many detailed studies of individual transition objects (post-AGB stars) exist, but it is not clear how these objects are related by evolutionary channels.  Moreover, there is general agreement that binary interactions must play a significant role in many well studied sources. Binarity is for instance invoked in the physical models to understand the observational characteristics of some spectacular geometries observed in PNe. More recently, also the geometries and kinematical structures around resolved post-AGB stars might be linked to binarity \\citep[][ and references therein]{Balick_2002}.  Since many uncertainties remain in our understanding of the final evolution of single stars, it is no surprise that this is even more the case when the star is a member of a binary system.  Direct detection of the binary nature of central stars of resolved nebulae is often difficult due to the high obscuration. Moreover, in crossing the HR-diagram, the stars must pass the pop\\,II Cepheid instability strip, in which pulsational instabilities occur. This makes radial velocity variations not a straightforward signature of orbital variations.  In the sample of optically bright post-AGB stars, binaries are being detected, however, and for an overview we refer to \\citet[][ and references therein]{Vanwinckel_2003}. One of the important observational characteristics of those binaries is the shape of their SED. They show a dust-excess starting near sublimation temperature, irrespective of the effective temperature of the central object and this despite the lack of a current dusty mass loss \\citep[][ and references therein]{Deruyter_2005, Deruyter_2006}.  Moreover, when available, the long wavelength fluxes show a black-body slope which indicates the presence of a component of large mm-sized grains. It is argued in \\cite{Deruyter_2006} that in all the investigated objects, gravitationally bound dust is present, likely in a Keplerian disc. Note that only for the most famous example, the \\object{Red Rectangle}, this dust emission is resolved and it shows a clear disc structure, both in the near-IR and in the visible \\citep{Menshchikov_2002,Cohen_2004}. Moreover, the gaseous component was spatially resolved using mm-interferometry \\citep{Bujarrabal_2003,Bujarrabal_2005}. The latter observations clearly demonstrated the Keplerian rotation of the disc. In all other cases, the presence of a disc is postulated.  More detailed studies of individual cases can be found: examples are \\object{89\\,Her} \\citep{Waters_1993}, \\object{HR\\,4049} \\citep{Waelkens_1991a, Dominik_2003} and \\object{IRAS08544-4431} \\citep{Maas_2003}.  Given the orbits detected so far, one of the conclusions is that it is clear that most binaries cannot have evolved along single star evolutionary tracks.  The high spatial resolution of the mid-IR instrument MIDI mounted on the VLTI interferometer of ESO makes this an ideal instrument to probe the circumstellar material around these binaries for two reasons: (i) the discs are likely compact so high spatial resolution measurements are needed to resolve the discs and (ii) the discs are shown to emit a significant part of their total luminosity in the N-band. We therefore carefully selected 2 binary post-AGB stars for which there is significant indirect evidence for the presence of a stable circumstellar dust reservoir. The data presented in this contribution are taken during Science Demonstration Time to illustrate the potential of MIDI coupled to the VLTI to study the compact circumstellar environment suspected in those evolved stars.  In Sect.~\\ref{sect:global_characteristics}, we introduce both objects and refine the orbital parameters published in our previous papers. The observational log is presented in Sect.~\\ref{sect:observations} while the reduction of the interferometric dispersed fringes is reported in Sect.~\\ref{sect:reduction}. We discuss our findings in Sect.~\\ref{sect:discussion} and come to our conclusions in Sect.~\\ref{sect:conclusions}. ",
        "conclusions": "\\label{sect:conclusions}  The main conclusion of our presented MIDI observations is that they prove the very compact nature of the circumstellar environments of \\object{HD\\,52961} and \\object{SX\\,Cen}. \\object{SX\\,Cen} is not resolved using a 45\\,m baseline, which gives an upperlimit of only 18\\,AU at the estimated distance. For the well resolved \\object{HD\\,52961}, the angular size in the N-band varies between 35 and 55\\,mas in a uniform disc approximation, which translates to a size of 50 and 80\\,AU.  Both stars have an effective temperature in the 6000\\,K range and since there is no evidence for a current dusty mass-loss we interpret these results as a very stringent proof of the existence of a stable reservoir near the star. A Keplerian disc seems the only plausible solution. The dust sublimation temperature is reached much further out than the binary orbits, hence the discs must be circumbinary. This is corroborated by the measured size of the dust-emission region around \\object{HD\\,52961}.  Given the size of the orbits, the discs were probably formed in a poorly understood phase of strong binary interaction, when the star was at giant dimensions. Both discs are O-rich and there is no evidence for a C-rich component. They were consequently formed prior to the late AGB evolution where the stars could have changed into C-stars. The mass of the companion of \\object{SX\\,Cen} (1.4 -- 1.9 M$_{\\odot}$) is probably within the range of C-star progenitors. We conclude that the normal single star AGB evolution was shortcut by the presence of a binary companion. Clearly the formation of a stable Keplerian disc is a key ingredient in the late evolution of both binaries.  \\object{SX\\,Cen} is an RV\\,Tauri star of photometric class b which shows a long term variability of the mean magnitude with a period similar to its orbital period probably due to variable circumstellar extinction in the line of sight during orbital motion. The inclination cannot be very small. Additional interferometric data on different projected angles will be necessary to probe the expected asymmetries.  The characteristics of the dust grains seem to be very different from normal single star outflows. This is shown in the mineralogy of the silicate resonance feature which shows for both objects a highly crystalline component and a size distribution with a much stronger component of large ($> 1$\\,$\\mu$m) grains than what is observed in outflows of AGB stars. It is not clear whether this reflects the formation history of the disc or this is due to the longer processing time of the dust in the Keplerian discs. Our analysis of \\object{HD\\,52961} shows that the crystallinity is clearly concentrated in the hotter inner region of the disc. Crystallisation by annealing is very temperature dependent and a similar picture arises as what is seen in the discs around some young stellar objects: the grains in the hot inner region were subject to a much stronger processing while in the outer region remained less processed. MIDI as spectrally dispersed N-band interferometer is an ideal instrument to study the chemo-physical structure of the inner regions of these discs."
    },
    "0601/physics0601012_arXiv.txt": {
        "abstract": "Diffractive/refractive optics, such as Phase Fresnel Lenses (PFL's), offer the potential to achieve excellent imaging performance in the x-ray and gamma-ray photon regimes.  In principle, the angular resolution obtained with these devices can be diffraction limited.  Furthermore, improvements in signal sensitivity can be achieved as virtually the entire flux incident on a lens can be concentrated onto a small detector area.  In order to verify experimentally the imaging performance, we have fabricated PFL's in silicon using gray-scale lithography to produce the required Fresnel profile. These devices are to be evaluated in the recently constructed 600-meter x-ray interferometry testbed at NASA/GSFC.  Profile measurements of the Fresnel structures in fabricated PFL's have been performed and have been used to obtain initial characterization of the expected PFL imaging efficiencies. ",
        "introduction": "\\label{sec:1}  The use of Phase Fresnel Lenses (PFL's) offer a mechanism to achieve superb imaging of astrophysical objects in the hard X-ray and gamma-ray energy regimes \\cite{Skinner1,Skinner2}.  In principle, PFL's can concentrate nearly the entire incident photon flux, modulo absorption effects,  with diffraction-limited imaging performance.  The impact of absorption is energy and material dependent, but can be almost negligible at higher photon energies. The performance of these diffraction optics is obtained via tailoring the Fresnel profile of each zone to yield the appropriate phase at the primary focal point.    However, PFL's have long focal lengths and are chromatic; the excellent imaging is available over a narrow energy range.  In order to demonstrate the imaging capabilities of these optics, we have fabricated ground-testable PFL's in silicon.  \\begin{figure} \\centering \\includegraphics[height=7.5cm]{Figure1.eps} \\caption{The ideal PFL profile of the first four Fresnel ridges compared to that for a stepped profile with 4 steps/ridge.} \\label{fig:1}       % \\end{figure} ",
        "conclusions": "We have fabricated ground-test PFL's in silicon using gray-scale lithography.  We have determined the imaging performance of these devices via analysis of the measured profiles of the fabricated optics.  These results indicate that the efficiencies, although less than ideal, are a significant improvement over the theoretical maximum that can be obtained with zone plates and phase-reversal zone plates. We plan on introducing these devices into the 600 m test beam to demonstrate their imaging capability and verify the anticipated efficiency determination via {\\it in situ} x-ray measurements. This material is based upon work supported by the National Aeronautics and Space Administration under Grant APRA04-0000-0087 issued through the Science Mission Directorate Office and by Goddard Space Flight Center through the Director's Discretionary Fund."
    },
    "0601/astro-ph0601396_arXiv.txt": {
        "abstract": "ASAS photometric observations of LS~1135, an O--type SB1 binary system with an orbital period of 2.7 days, show that the system is also eclipsing. This prompted us to re--examine the spectra used in the previously published spectroscopic orbit. Our new analysis of the spectra obtained near quadratures, reveal the presence of faint lines of the secondary component. We present for the first time a double--lined radial velocity orbit and values of physical parameters of this binary system. These values were obtained by analyzing ASAS photometry jointly with the radial velocities of both components performing a numerical model of this binary based on the Wilson--Devinney method. We obtained an orbital inclination $i \\sim 68^{\\circ}.5$. With this value of the inclination we deduced masses  $M_1 \\sim 30 \\pm 1 M_{\\odot}$ and $M_2 \\sim 9 \\pm 1 M_{\\odot}$;  and radii $R_1 \\sim 12 \\pm 1 R_{\\odot}$ and $R_2 \\sim 5 \\pm 1 R_{\\odot}$ for primary and secondary components, respectively. Both components are well inside their respective Roche lobes. Fixing the $T_{eff}$ of the primary to the value corresponding to its spectral type (O6.5V), the $T_{eff}$ obtained for the secondary component corresponds approximately to a spectral type of B1V. The mass ratio $M_2 /M_1 \\sim 0.3$ is among the lowest known values for spectroscopic binaries with O--type components. \\\\ ",
        "introduction": "The star LS 1135 ($\\alpha_{2000.0} = 08^{h} 43^{m} 50^{s}$,{} $\\delta_{2000.0}$ = -$46^\\circ $ 07' 09\"), a member of the OB association Bochum 7 (Vela OB 3), was discovered by Corti et al. (2003, hereafter CNM) to be a single-lined binary (SB1) system with an orbital period of 2.7532 days. The spectrum was classified as O6.5V ((f)). The rather high amplitude of the radial velocity variations and the lack of detected spectral lines of the secondary component led CNM to suggest that the secondary might be an early B type star.  Magnitudes and colors for LS~1135 from photoelectric photometry have been published by Moffat \\& Vogt (1975) as $V = 10.88, B-V = 0.4, U-B = -0.68$, and by Drilling (1991) as $V = 10.88, B-V = 0.38, U-B = -0.66$. LS~1135  was found to be a photometrically variable star in the ``All Sky Automated Survey'' (ASAS) (cf. Pojma\\'nski 2003). It is catalogued as an eclipsing system, and a period of 2.7532 days was found independently in the photometric data, in perfect agreement with the period obtained from the radial velocities by CNM. The fact that LS~1135 is an eclipsing binary, motivated us to  search for signatures of the secondary component by reinspecting the spectra used by CNM. Here we report the discovery of the spectral lines of the secondary, and calculate for the first time a double--lined radial velocity orbit and a set of values of the main physical parameters of the components of LS~1135. These are based on a simultaneous analysis of the photometric light curve and radial velocity variations of both components by means of the Wilson-De\\-vinney (W--D) Code (Wilson \\& De\\-vinney 1971, Wilson 1990, Wilson \\& Van Hamme 2004).  \\begin{figure} \\includegraphics[width=240pt]{MF1375f1.ps} \\caption{A 4' x 4' image of the Digitized Sky Survey, centered in LS~1135. The circle represents the ASAS 3-pixels (= 46\".5) photometric aperture. (North is up and East to the left).} \\label{DSS1} \\end{figure} ",
        "conclusions": "We have obtained in this work a first approximation to estimate the absolute physical parameters of the LS~1135 O-type binary system. From a simultaneous analysis of the radial velocities and the light curve with the W--D code, we find the following: \\begin{itemize}  \\item Fixing the $T_{eff}$ of the primary component to the value corresponding to its spectral type, namely O6.5V, from the best fitting W--D model we find $T_{eff}$ $25500~^{\\circ}\\!K$ for the secondary component. This value corresponds approximately to a spectral type B1V, according to the scales of different authors (B$\\ddot{o}$hm-Vitense, 1981; Schmidt-Kaler, 1982; Crowther, 1997).  \\item The inclination of the orbital plane of the system relative to the sky plane was found to be $68.5 \\pm 0.5 ^{\\circ}$ according the best fitting W--D model. With this value of $i$, we have obtained values for the individual masses of the primary and secondary components $M_{1} = 30 \\pm 1~ M_{\\odot}$ and $M_{2} =  9 \\pm 1~ M_{\\odot}$ and for the mean radii $R_{1} = 12 \\pm 1~ R_{\\odot}$ and $R_{2} =  5 \\pm 1~ R_{\\odot}$. Both components appear to be well inside their respective Roche lobes, and therefore LS~1135 is a detached binary still on the main sequence. Detached eclipsing binaries with accurate elements at the high mass end of the main sequence are exceedingly few in number, and are very much needed as the benchmarks of theoretical stellar models. Thus LS~1135 appears as an excellent candidate for an accurate determination of physical parameters in a high mass binary system. Within the uncertainties, the first values of physical parameters of the binary components of LS~1135 which we have determined are in fair agreement with recent tabulations of Galactic O star parameters based on models of stellar atmospheres (cf. Martins et al. 2005). Of course, for more accurate values a better defined light curve is needed, as well as spectra with higher signal-to-noise ratio for a more accurate determination of the radial velocities of the secondary component.  \\item It is also  interesting to mention that the mass ratio which we have obtained from the radial velocities $q = M_{2}/M_{1}\\sim 0.31$, if confirmed, would be among the lowest values recorded for O type main sequence binaries. Therefore, the physical parameters of the primary O6.5V type component would be better defined in this case, since the perturbations introduced by the secondary component on the spectral behaviour of the primary would be minor as compared to binary systems with similar components.  \\end{itemize}"
    },
    "0601/astro-ph0601603.txt": {
        "abstract": "We present 2.5D radiation-hydrodynamics simulations of the accretion-induced collapse (AIC) of white dwarfs, starting from 2D rotational equilibrium configurations of a 1.46-M$_{\\odot}$ and a 1.92-M$_{\\odot}$ model. Electron capture leads to the collapse to nuclear densities of these cores within a few tens of milliseconds. The shock generated at bounce moves slowly, but steadily, outwards. Within 50-100\\,ms, the stalled shock breaks out of the white dwarf along the poles. The blast is followed by a neutrino-driven wind that develops within the white dwarf, in a cone of $\\sim$\\,40$^{\\circ}$ opening angle about the poles, with a mass loss rate of 5-8$\\times$10$^{-3}$M$_{\\odot}$yr$^{-1}$. The ejecta have an entropy on the order of 20-50 k$_{\\rm B}$/baryon, and an electron fraction distribution that is bimodal. By the end of the simulations, at $\\ge$600\\,ms after bounce, the explosion energy has reached 3-4$\\times$10$^{49}$\\,erg and the total ejecta mass has reached a few times 0.001\\,M$_{\\odot}$. We estimate the asymptotic explosion energies to be lower than 10$^{50}$\\,erg, significantly lower than those inferred for standard core collapse. The AIC of white dwarfs thus represents one instance where a neutrino mechanism leads undoubtedly to a successful, albeit weak, explosion.  We document in detail the numerous effects of the fast rotation of the progenitors: The neutron stars are aspherical; the ``$\\nu_{\\mu}$'' and $\\bar{\\nu}_e$ neutrino luminosities are reduced compared to the $\\nu_e$ neutrino luminosity; the deleptonized region has a butterfly shape; the neutrino flux and electron fraction depend strongly upon latitude ({\\it \\`a la} von Zeipel); and a quasi-Keplerian 0.1-0.5-M$_{\\odot}$ accretion disk is formed. ",
        "introduction": "Stars can follow a few special evolutionary routes to form an unstable Chandrasekhar mass core. A main-sequence star of more than $\\sim$8\\,\\mo evolves to form either a degenerate O/Ne/Mg core (Barkat et al. 1974; Nomoto 1984,1987; Miyaji \\& Nomoto 1987) or a degenerate Fe core (Woosley \\& Weaver 1995), which, due to photodisintegration of heavy nuclei and/or electron capture, collapses to form a protoneutron star (PNS). If an explosion ensues, the event is associated with a Type II supernovae (SN). Less massive stars end their lives as white dwarfs. White dwarfs located in a binary system may accrete from a companion and achieve the Chandrasekhar mass, triggering the thermonuclear runaway of the object and leading to Type Ia SN, leaving no remnant behind.  However, a third class of objects is expected. Theoretically, massive white dwarfs with O/Ne/Mg cores, due to their high central density ($\\sgreat$10$^{10}$\\,g\\,cm$^{-3}$), experience rapid electron capture that leads to the collapse of the core. This is accretion-induced collapse (AIC), an alternative path to stellar disruption through explosive burning, currently associated with Type Ia SN (Nomoto \\& Kondo 1991). % Various paths may lead to a white dwarf with such properties. % LUC It is presently unclear what fraction of all white dwarfs will lead to AICs, but of those white dwarfs that evolve to form a Chandrasekhar-mass O/Ne/Mg core, all will necessarily undergo core collapse. One formation channel is the coalescence of two white dwarfs (Mochkovitch \\& Livio 1989), with either C/O or O/Ne/Mg cores, although few such binary systems have yet been observed with a cumulative mass above the Chandrasekhar mass. There is still uncertainty as to whether such binary systems would not undergo thermonuclear runaway rather than collapse. Since the coalescence of two white dwarfs requires a shrinking of the orbit through gravitational radiation, these systems will take many gigayears to coalesce. An alternative formation mechanism is via single-degenerate systems, through a combination of high original white dwarf mass and mass and angular-momentum accretion by mass transfer from a (non-degenerate) H/He star (Nomoto \\& Kondo 1991). Binary star population synthesis codes predict the occurence of the AIC of white dwarfs with a galactic rate of 8$\\times$10$^{-7}$ yr$^{-1}$ to 8$\\times$10$^{-5}$ yr$^{-1}$, depending, amongst other things, on the treatment of the common-envelope phase and mass transfer (Yungelson \\& Livio 1998). The set of parameters leading to a Type Ia rate of 10$^{-3}$\\,yr$^{-1}$ corresponds to an AIC rate of 5$\\times$10$^{-5}$ yr$^{-1}$. The observed Type Ia rate of $\\sim$3$\\times$10$^{-3}$\\,yr$^{-1}$ (Madau et al. 1998; Blanc et al. 2004; Manucci et al. 2005) would imply a galactic AIC rate of 1.5$\\times$10$^{-4}$\\,yr$^{-1}$. These rates are likely functions of galaxy and metallicity (Yungelson \\& Livio 2000; Belczynski et al. 2005; Greggio 2005; Scannapieco \\& Bildsten 2005). Based on r-process nucleosynthetic yields obtained from previous simulations of the AIC of white dwarfs, Fryer et al. (1999) inferred rates ranging from $\\sim$10$^{-5}$ to $\\sim$10$^{-8}$\\,yr$^{-1}$. Overall, AICs are not expected to occur more than once per 20--50 standard Type Ia events; because they are intrinsically rarer, they remain to be identified and observed in Nature.  Whatever their origin, their fundamental nature is to accrete both mass and angular momentum from a companion object. Rotation is, therefore, a key physical component. Prior to core collapse, differential rotation acts as a stabilizing agent for shell burning by widening its spatial extent through enhanced mixing and reducing the envelope density through centrifugal support (Yoon \\& Langer 2004). Mass accretion also leads to an increase of the central density, which may rise up to a few $10^{10}$ g\\,cm$^{-3}$, establishing suitable conditions for efficient electron capture on Mg/Ne nuclei. By the time of core collapse and depending on the evolutionary path followed, such white dwarfs may cover a range of masses from $\\sgreat$1.35\\,\\mo up to $\\sim$2\\,\\mo (2.7\\,\\mo) in the case of a non-degenerate (degenerate) companion, potentially well in excess of the standard Chandrasekhar limit, and possessing an initial rotational energy up to $\\sim$10\\% of their gravitational binding energies. At such values, the centrifugal potential leads to a deformation of the white dwarf from spherical symmetry, equipotentials and isopressure surfaces adopting a peanut-like shape in cross section for the largest rotation rates.  Such structures are obtained in the 2D differentially rotating equilibrium white dwarf models constructed by Yoon \\& Langer (2005, YL05; see also Liu \\& Lindblom 2001).  % Guided by 1D studies of rotating accreting white dwarfs (Yoon \\& Langer 2004), % Yoon \\& Langer (2005, YL05) constructed 2D differentially rotating equilibrium white dwarf models for a wide % range of central densities, total angular momenta and masses. In the present work, we study % two progenitors with masses at the extremes of the spectrum computed by YL05, i.e., the models with % masses 1.46\\,\\mo and 1.92\\,\\mo. Both start with a central density of 5$\\times$10$^{10}$\\,g\\,cm$^{-3}$.  In the past, the collapse of O/Ne/Mg cores originating from stars in the 8-10\\,\\mo range has been studied in 1D by Baron et al. (1987ab), Mayle \\& Wilson (1988), and Woosley \\& Baron (1992), who showed that the shock generated at core bounce stalls rather than leading to a prompt explosion. However, Hillebrandt et al. (1984) and Mayle \\& Wilson (1988) obtained delayed explosions, the former after 20-30\\,ms, and the latter after $\\sim$200\\,ms, supposedly driven by neutrino energy deposition behind the stalled shock. Woosley \\& Baron (1992) found the emergence of a sustained neutrino-driven wind, with a mass loss rate of 0.005\\,\\mo\\,s$^{-1}$ and with an ejecta electron fraction of the order of 0.45, showing promise, modulo uncertainties, for a contribution to the enrichment of the ISM in r-process elements. Using 1D/2D SPH simulations, Fryer et al. (1999) reproduced the simulations of Woosley \\& Baron (1992), confirmed that the shock stalls due to the copious neutrino losses associated with core bounce, did not find prompt explosion, and, depending on the equation of state (EOS) employed, observed a delayed explosion. They focused mostly on the early phase, prior to the neutrino-driven wind, and found an ejected mass of low-$Y_{\\rm e}$ material of $\\sim$0.05\\,\\mo, depending on adopted model assumptions. Their 2D simulations with solid-body rotation showed similar properties to their 1D equivalents, the authors attributing the small differences to the different grid resolution. This may partly stem from the essentially spherical explosion triggered just $\\la$100\\,ms after core bounce, ejecting the fast-rotating material in the outer mantle. Consequently, at the end of their 2D simulations, they obtain very slow rotation rates for the PNS, i.e., of $\\sim$1\\,s. Recently, Kitaura et al. (2005) re-inspected the collapse of the progenitor model used by Hillebrandt et al. (1984). They confirmed again the by-now well-accepted idea that no prompt explosion occurs, but instead obtain a successful, though sub-energetic, delayed explosion in spherical symmetry, powered by neutrino heating and a neutrino-driven wind that sets in $\\sim$200\\,ms after bounce.  A primary motivation for this work is to improve upon these former investigations that assumed one-dimensionality, sphericity, and/or zero-rotation, and start instead from the more physically-consistent 2D models of YL05, thereby fully accounting for the effects of rotation on the collapse, bounce, and post-bounce evolution of the white dwarf core and envelope, as well as for the strong asphericity of the progenitor. Our study uses VULCAN/2D (Livne et al. 2004; Walder et al. 2005) to perform 2D Multi-Group Flux Limited Diffusion (MGFLD) radiation hydrodynamics simulations. As we will demonstrate, rotation plays a major role in the post-bounce evolution, making 1D investigations of such objects of limited utility. By carrying out the simulations from $\\sim$30\\,ms before bounce to $\\sgreat$600\\,ms after bounce, we capture a wide range of physical processes, including the establishment of a strong, fast, and aspherical neutrino-driven wind. We model the centrifugally-supported equatorial regions and the large angular momentum budget leading to the formation of a sizable accretion disk. % LUC Moreover, an Eulerian investigation is better-suited than a Lagrangean approach to explore the neutrino-driven wind that develops after $\\sgreat$200\\,ms. Finally, despite the relative scarcity of AIC in Nature, these simulations represent interesting examples for the formation of disks around neutron stars. % , and for the collapsar model of gamma-ray bursts.  The main findings of this work are the following: We find that the AIC of white dwarfs forms $\\sim$1.4-\\mo neutron stars, expelling a modest mass of a few 10$^{-3}$\\,\\mo mostly through a neutrino-driven wind that develops $\\sgreat$200\\,ms after bounce, and that they lead to very modest explosion energies of 5-10$\\times$10$^{49}$erg. Accounting for the rotation and the asphericity of the progenitor white dwarfs reveals a wealth of phenomena. The shock wave generated at core bounce emerges through the poles rather than the equator, and it is in this excavated polar region that the neutrino-driven wind develops. The strong asphericity of the newly-formed protoneutron stars leads to a latitudinally-dependent neutrino flux, while the effects of rotation modify the relative flux magnitude of different neutrino flavors. Besides mass accretion by both the neutron star and mass ejection by the initial blast and the subsequent wind, we find a sizable component that survives and resides in a quasi-Keplerian disk, which obstructs the wind flow at low latitudes. This disk will be accreted by the neutron star only on longer, viscous timescales.  In the next two sections, we present the two selected progenitor models in more detail; we also discuss the radiation-hydrodynamics code VULCAN/2D and the various assumptions made. In \\S\\ref{sect_results}, we present the simulation results, focusing on the general temporal evolution from the start until $\\sgreat$600\\,ms, a time by which the neutrino-driven wind has reached a steady state. We then analyse in more detail the various components of these simulations. In \\S\\ref{sect_pns}, we discuss the properties of the nascent neutron stars, with special attention paid to the geometry of the neutrinospheres. In \\S\\ref{sect_nu}, we discuss the neutrino signatures, both in terms of luminosity and energy distribution. In \\S\\ref{sect_disk}, we focus on the residual material lying at low latitudes, forming a quasi-Keplerian disk. In \\S\\ref{sect_ener}, we turn to the energetics of the explosion and describe in detail the main component of the simulations at late times, i.e., the neutrino-driven wind. In \\S\\ref{sect_ye}, we analyse the electron fraction of the ejected material and address the relevance of the AIC of white dwarfs for neutron-rich element pollution of the interstellar medium. In \\S\\ref{sect_gw}, we present the gravitational-wave signal predicted for the AIC of our white dwarf models. In \\S\\ref{sect_conc}, we wrap up with a discussion of the main results of this investigation and present our conclusions.  % ~\\ref{sect_progenitor} % ~\\ref{sect_VULCAN} % ~\\ref{sect_results} % ~\\ref{sect_pns} % ~\\ref{sect_nu} % ~\\ref{sect_disk} % ~\\ref{sect_ener} % ~\\ref{sect_ye} % ~\\ref{sect_conc} ",
        "conclusions": "\\label{sect_conc}  We have presented a radiation/hydrodynamic study with the code VULCAN/2D of the collapse and post-bounce evolution of massive rotating high-central-density white dwarfs, starting from physically consistent 2D rotational equilibrium configurations (Yoon \\& Langer 2005). The main results of this study are:  \\begin{itemize} \\item The AIC of white dwarfs leads to a successful explosion with modest energy $\\sles$10$^{50}$\\,erg, thus comparable to the energies obtained through the collapse of O/Ne/Mg core of stars with $\\sim$8-11\\,\\mo main sequence mass (Kitaura et al. 2005; Buras et al. 2005b). This is, however, underenergetic, by a factor of about ten, compared with the inferred value for the core collapse of more massive progenitors leading to Type II Plateau supernovae. Although less and less likely to be the engine behind most core-collapse supernova explosions, the neutrino mechanism can successfully power explosions of low-mass progenitors and AICs due to the limited mantle mass and steeply declining accretion rate.  \\item Due to high-mass and angular-momentum accretion, white dwarf progenitors leading to AIC can have masses of up to $\\sim$2\\,\\mo and rotate fast, with rotational to gravitational energy ratios of up to a few percent prior to collapse. The asphericity of such white dwarfs allows the shock generated at core bounce to escape along the poles in just a few tens of milliseconds, opening a hole in the white dwarf along the poles. The blast expands and wraps around the progenitor and escapes the grid ($\\sim$5000 km) within a few hundred milliseconds, at which time a neutrino-driven wind has grown in the pole-excavated region of the white dwarf. Both the original blast and the wind show strong latitudinal variations, partly constrained by the obstructing uncollapsed equatorial disk regions of the progenitor, whose centrifugal support prevents it from collapsing on a dynamical timescale. Rotation in such progenitors, thus, affects both directly and indirectly the morphology of the explosion.  \\item The neutron stars formed have masses on the order of 1.4\\,\\mo, with rotation periods close to a millisecond in the rigidly rotating inner $\\sim$30\\,km. The final rotational to gravitational energy ratios, for our two test cases, cover 0.06 to 0.26, the latter being large enough to grow non-axisymmetric instabilities. At the end of our simulations, the neutron stars are oblate and pinched along the poles, with polar and equatorial radii in the ratio 1:15 for the faster rotating (1.92-\\mo) model.  \\item The morphology of the neutron star leads to a latitudinal variation of the neutrino flux, the net energy gain, and the temperature. In the faster rotating model, the ``$\\nu_\\mu$'' neutrino flux is reduced, while the anti-electron neutrino flux is a factor of ten lower than that of the electron-neutrino. This raises the electron fraction of the ejected material to values close to 0.5 along the poles, but to only $\\sim$0.25 at lower latitudes, since the corresponding neutrinosphere is more remote and, thus, the electron-neutrino flux is smaller. This introduces a latitudinal dependence of the electron fraction of the ejected material, but more importantly allows neutron-rich material, with entropy on the order of 20-40\\,$k_{\\rm B}$/baryon, to escape. Thus, a low-entropy r-process might take place under these conditions.  \\item The high original angular momentum of the progenitor follows the mass and is, thus, found mostly in the neutron star at the end of the simulation. However, rotational energy is also given to the ejecta, which uses it to gain (planar) radial kinetic energy to escape the potential well, while the rest is found in a quasi-Keplerian disk of up to 0.5\\,\\mo in the 1.92-\\mo model. This disk is an essential component of these AICs; it collimates the explosion and the neutrino-driven wind and also suggests a second stage of long-term accretion onto the compact remnant.  \\item The total ejected mass is only of the order of a few times 0.001\\,\\mo, with only a quarter in the form of $^{56}$Ni. The original blast and the short-lived neutrino-driven wind will lead to a considerable brightening of the object, but the small ejecta mass will quickly become optically-thin, gamma rays leaking out rather than depositing their energy to power a durable light curve. Therefore, these explosions should be underluminous and very short lived. Their appearance may also vary considerably with viewing angle, depending on the mass of the progenitor and the presence of a sizable disk in the equatorial regions.  \\end{itemize}  This study has shown that more consistent, rotating 2D models alter considerably our understanding of accretion-induced collapse, previously obtained under the simplifying assumptions of spherical symmetry and/or zero rotation. Further improvements will come by including a consistent temperature structure for the progenitor white dwarf and by accounting for the effects of magnetic fields.  Due to the rapid differential rotation, magnetic fields could be amplified considerably and result in MHD jets that might alter yet again our overall picture of accretion-induced collapse and the energetics of the phenomenon. Three-dimensional effects may also alter the protoneutron star properties presented here, since the faster rotating (1.92-\\mo) model is expected to experience non-axisymmetric instabilities."
    },
    "0601/hep-th0601126_arXiv.txt": {
        "abstract": " ",
        "introduction": "The quest for unification of fundamental interactions with gravity has led to the idea of extra spatial dimensions. String Theory, or some form of it, considered at present as a promising consistent theoretical framework for such a unification at the quantum level, requires extra spatial dimensions. In this higher-dimensional context, which has received a lot of attention in the last few years \\cite{ADD, RS}, gravity propagates in $D=4+n$ dimensions ({\\textit{Bulk}}), while all Standard Model degrees of freedom are assumed to be confined on a four-dimensional $D$-{\\textit{Brane}}. In theories with large extra dimensions \\cite{ADD} the traditional Planck scale $M_{Pl}\\sim 10^{18}$\\,GeV is only an effective scale, related to the fundamental higher-dimensional gravity scale $M_{*}$ through the relation $M_{Pl}^2 \\sim M_{*}^{n+2}R^n$, where $R\\sim (V_n)^{1/n}$ is the effective size of the $n$ extra spatial dimensions. If $R \\gg \\ell_{Pl}\\approx 10^{-33}\\,cm$, the scale $M_{*}$ can be substantially lower than $M_{Pl}$. If $M_{*}$ is sufficiently low, trans-planckian particle collisions could be feasible in accelerators or other environments. The product of such collisions would be objects that probe the extra dimensions. Black holes are among the possible products of trans-planckian collisions \\cite{creation}. Black holes with a horizon size $r_H$ smaller than the size of extra dimensions $R$ will be higher-dimensional objects centered at the brane and extending in the bulk. A classical treatment would require that the mass of the black hole $M_{BH}$ would have to be larger than the fundamental gravity scale $M_{*}$. Such black holes can be created in ground based colliders, although their appearance in cosmic rays is possible as well (for references, see the reviews \\cite{Kanti, reviews}). A black hole created in such trans-planckian collisions is expected to undergo a short {\\textit{``balding\"}} phase in which it will shed its additional quantum numbers, apart from charge, angular momentum and mass. Then, after a more familiar Kerr-like phase, in which it will lose its angular momentum, a Schwarzschild phase follows during which the black hole will gradually lose its mass through the emission of Hawking radiation both in the bulk and in the brane, consisting of elementary particles of a very distinct thermal spectrum. This black hole radiation has been the subject of both analytical and numerical studies. This includes lower spin degrees of freedom \\cite{kmr1, FS, HK1, BGK, Barrau, Jung} as well as graviton emission in the bulk \\cite{Naylor}. The study of graviton emission in the bulk corresponds to the study of perturbations in a gravitational background. The formalism for the treatment of gravitational perturbations of a higher-dimensional non-rotating black hole has been developed in \\cite{IK}. In the present article we consider the evaporation of $(4+n)$-dimensional non-rotating black holes into gravitons and calculate the energy emission rate for gravitons in the bulk obtaining analytical solutions of the master equation satisfied by gravitational perturbations. Our results, valid in the low-energy regime, are complementary to existing studies in the intermediate energy regime \\cite{Naylor}. These results show a vector radiation dominance for every value of $n$, while the relative magnitude of the energy emission rate of the subdominant scalar and tensor radiation depends on $n$. The low-energy emission rate in the bulk for gravitons is well below that for a scalar field, due to the absence of the dominant $\\ell=0,1$ modes from the gravitational spectrum, although higher partial waves are likely to modify this behaviour at higher energies. The calculated low-energy emission rate, for all types of degrees of freedom decreases with $n$, although the full energy emission rate, integrated over all frequencies, is expected to increase with $n$, as in the previously studied case of a bulk scalar field. ",
        "conclusions": "A higher-dimensional black hole created in the context of a brane-world theory will decay through the emission of Hawking radiation both in the bulk and on the brane. The type of particles emitted along each ``channel\" are determined by the assumptions of the particular model, with only gravitons and possibly scalars propagating in the bulk within the framework of the Large Extra Dimensions scenario. Although the emission of scalars has been studied in detail, until recently gravitons have received little attention.  In this work, we have addressed this gap in the literature by investigating the emission of tensor, vector and scalar gravitational modes in the bulk from a ($4+n$)-dimensional Schwarzschild black hole. Working in the low energy regime, we analytically solved the corresponding field equations and computed the absorption probability in each case. Both a complete analytic expression and its asymptotic low-energy simplification were studied in detail, and their dependence on the angular momentum number $l$ and number of extra dimensions $n$ was examined. Although numerically different as the energy increases, these two expressions have identical qualitative behaviours, that reveal an increase in the value of the absorption probability with increasing energy and suppression as the number of partial wave $l$ and extra dimensions $n$ increase.  The complete analytical expression for the absorption probability was then used to derive the contribution of each gravitational degree of freedom to the total graviton emission rate of the black hole in the bulk. Accounting for the rapid proliferation of state multiplicities with $l$ or $n$, a sum over the first twelve partial waves was performed to obtain the total low-energy emission rate for each gravitational degree of freedom. Our results show that vector perturbations are the dominant mode emitted in the bulk for all values of $n$. The relative emission rates of the subdominant scalar and tensor modes depend on $n$, with scalars foremost at small $n$ and tensors more prevalent at high $n$. The absence of the $l=0,1$ partial waves, dominant in the low-energy regime, from all gravitational spectra causes even the total gravitational emission rate to be subdominant to that of scalar fields in the bulk. Finally, as previously found for bulk scalar fields, the energy emission rates for all types of gravitational perturbations are suppressed with the number of extra dimensions in the entire low-energy regime.  In this work we have focused on the low-energy part of the Hawking radiation spectrum as it permitted analytical calculation and derivation of closed form expressions for the absorption probability for gravitons. Although the radiation emitted in the bulk is not directly observable, it determines the energy left for emission on our brane. In this context, our results, in addition to their theoretical interest, would be of particular use to experiments developed to detect the low-energy spectrum of radiation emitted from a higher-dimensional black hole. As the exact form of the complete gravitational spectrum is still pending, we hope to return in the near future with results from a numerical analysis that could only provide the answer to this question.  {\\bf Note added in proof.} While this work was at its last stages, two relevant papers appeared in the literature, \\cite{Park} and \\cite{Cardoso}. In particular, the latter work \\cite{Cardoso}, that studies the graviton emission rate in the bulk, overlaps with ours with respect to the analytical results for the cases of tensor and vector gravitational perturbations, which are in agreement with ours.    \\bigskip  {\\bf Acknowledgments.} S.C and O.E. acknowledge PPARC and I.K.Y. fellowships, respectively. The work of P.K. is funded by the UK PPARC Research Grant PPA/A/S/ 2002/00350. This research was co-funded by the European Union in the framework of the Program $\\Pi Y\\Theta A\\Gamma O PA\\Sigma-II$ of the {\\textit{``Operational Program for Education and Initial Vocational Training\"}} ($E\\Pi EAEK$) of the 3rd Community Support Framework of the Hellenic Ministry of Education, funded by $25\\%$ from national sources and by $75\\%$ from the European Social Fund (ESF)."
    },
    "0601/astro-ph0601661_arXiv.txt": {
        "abstract": "The rapid follow-up of gamma-ray burst (GRB) afterglows made possible by the multi-wavelength satellite {\\em Swift}, launched in November 2004, has put under a microscope the GRB early post-burst behavior, This is leading to a significant reappraisal and expansion of the standard view of the GRB early afterglow behavior, and its connection to the prompt gamma-ray emission. In addition to opening up the previously poorly known behavior on minutes to hours timescales, two other new pieces in the GRB puzzle being filled in are the the discovery and follow-up  of short GRB afterglows, and the opening up of the $z\\simg 6$ redshift range.  We review some of the current theoretical interpretations of these new phenomena. ",
        "introduction": "\\label{sec:obs}  Compared to previous satellites, Swift has made a large difference on two main accounts. First, the sensitivity of the Burst Alert Detector (BAT, in the range 20-150 keV) is a factor $\\sim$ 5 higher than for the corresponding instruments in the predecessor CGRO-BATSE, BeppoSAX and HETE-2. Second, Swift can slew in less then 100 seconds in the direction determined by the BAT instrument, positioning its much higher angular resolution X-ray (XRT) and UV-Optical (UVOT) detectors on the burst \\cite{gehr05_thisproc}  As of December 2005, at an average rate of 2 bursts detected per week, over 100 bursts had been detected by BAT, of which 90\\% were followed promptly with the XRT within $350$ s from the trigger, and about half within 100 s \\cite{burrows05}, while $\\sim 30\\%$ were detected with the UVOT \\cite{roming05}.  Of these, over 23 resulted in redshift determinations. Ten short GRB were detected, of which five had detected X-ray afterglows, three had optical, and one had a radio afterglow, and five had a redshift determination.  The new observations brings the total redshift determinations to over 50 since 1997 when BeppoSAX enabled the first one. The redshifts based on Swift have a median $z\\simg 2$, which is a factor $\\sim 2$ higher than the median of those previously culled via BeppoSAX and HETE-2, \\cite{berger05a}. This can be ascribed to the higher sensitivity of BAT and the prompt accurate positions from XRT and UVOT, making possible ground-based detection at a stage when the afterglow is much brighter. The highest Swift-enabled redshift so far is in GRB 050904, obtained with Subaru, $z=6.29$ \\cite{kawai05}, the second highest being GRB 050814 at $z=5.3$, whereas the previous Beppo-SAX era record was $z=4.5$. The relative paucity of UVOT detections versus XRT detections may be ascribed in part to this higher median redshift, and in part to the higher dust extinction at the implied  shorter rest-frame wavelenghts for a given observed frequency \\cite{roming05}, although additional effects may be at work too.  The BAT light curves show that in some of  the bursts which fall in the ``long\" category ($t_\\gamma \\simg 2$ s) faint soft gamma-ray tails can be followed  which extend the duration by a factor up to two beyond what BATSE could have detected \\cite{gehr05_thisproc}. A rich trove of information on the burst and afterglow physics has come from detailed XRT light curves, starting on average 100 seconds after the trigger, together with the corresponding BAT light curves and spectra. This suggests a canonical X-ray afterglow \\cite{nousek05} with one or more of the following:\\hfill \\lbr 1) an initial steep decay $F_X \\propto t^{-\\alpha_1}$ with a 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$ (or photon number index $2 \\siml \\Gamma=\\alpha+1 \\siml 3$), extending up to a time $300 \\s \\siml t_1 \\siml  500 \\s$;\\hfill \\lbr 2) a 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$;\\hfill \\lbr 3) a ``normal\" decay $F_X \\propto t^{-\\alpha_3}$ with $1.1 \\siml \\alpha_3 \\siml 1.7$ and $0.7 \\siml \\beta_2 \\siml 1.2$ (generally unchanged the previous stage), up to a time $t_3 \\sim 10^5 \\s$, or in some cases longer;\\hfill \\lbr 4) In some cases, a steeper decay $F_X \\propto t^{-\\alpha_4}$ with $2\\siml \\alpha_4 \\siml 3$, after $t_4\\sim 10^5 \\s$;\\hfill \\lbr 5) In about half the afterglows, one or more X-ray flares are observed, 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_{fl}}$ with $3 \\siml \\alpha_{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}[h] {\\includegraphics[height=.3\\textheight,width=.9\\textwidth]{xrt_genlc.eps}} \\caption{Schematic features seen by the XRT in bursts detected by Swift \\cite{zh05a} (see text).} \\end{figure}  Another major advance achieved by Swift was the detection of the long burst GRB 050904, which broke through the astrophysically and psychologically important redshift barrier of $z\\sim 6$. This burst was very bright, both in its prompt $\\gamma$-ray emission ($E_{\\gamma , iso} \\sim 10^{54}$ erg) and in its X-ray afterglow. Prompt ground-based optical/IR upper limits and a J-band detection suggested a photometric redshift $z>6$ \\cite{haislip05}. Spectroscopic confirmation with the 8.2 m Subaru telescope gave a $z=6.29$ \\cite{kawai05}. There are several striking features to this burst. One is the enormous X-ray brightness, exceeding for a full day the X-ray brightness of the most distant X-ray quasar know to-date, SDSS J0130+0524, by up to a factor $10^5$ in the first minutes \\cite{watson05}. The implications as a tool for probing the IGM are thought-provoking. Another feature is the extremely variable X-ray light curve, showing many large amplitude flares extending up to at least a day. A third exciting feature is the report of a brief, very bright IR flash \\cite{boer05}, comparable in brightness to the famous $m_V\\sim 9$ optical flash in GRB 990123.  The third major advance from Swift was the discovery and localization of short GRB afterglows. As of December 2005 nine short bursts had been localized by Swift, while in the same period HETE-2 discovered two, and one was identified with the IPN network. In five of the Swift short bursts an X-ray afterglow was measured and followed up, with GRB 050709, 050724 and 051221a showing an optical afterglow, and 050724 also a radio afterglow, while 040924 had an optical afterglow but not an X-ray one \\cite{fox05_thisproc}.  These are the first afterglows detected for short bursts. Also, for the first time, host galaxies were identified for these short bursts, which in four cases are early type (ellipticals) and in two cases are irregular galaxies. The redshifts of four of them are in the range $z\\sim 0.15-0.5$, while another one was initially given as $z=0.8$ but more recently has been reported as $z \\simeq 1.8$ \\cite{berger05_thisproc}. The median $z$ is $\\siml 1/3-1/2$ that of the long bursts. There is no evidence for significant star formation in any of these host environments, which corresponds to what one would expect for neutron star mergers or neutron star-black hole mergers, the most often discussed progenitor candidates (it would also be compatible with other progenitors involving old compact stars).  The first short burst seen by Swift, GRB 05059b, was a low luminosity ($E_{iso}\\sim 2\\times 10^{48}$ erg) burst with a simple power-law X-ray afterglow  which could only be followed for $\\sim 10^4$ s \\cite{gehr05_0509}. The third one, GRB 050724, was brighter, $E_{iso}\\sim 3\\times 10^{50}$ erg, and could be followed in X-rays for at least $10^5$ s \\cite{barth05_0724}. The remarkable thing about this burst's X-ray afterglow is that it resembles the typical X-ray light curves described above for long GRB -- except for the lack of a slow-decay phase, and for the short prompt emission which places it the the category of short bursts, as well as the elliptical host galaxy candidate. It also has X-ray flares, at 100 s and another one at $3\\times 10^4$ s. The first flare has the same fluence as the prompt emission, while the late flare has $\\sim 10\\%$ of that. The interpretation of these pose interesting challenges, as discussed below. ",
        "conclusions": ""
    },
    "0601/astro-ph0601382_arXiv.txt": {
        "abstract": "It has been known that \\cl\\ is not a relaxed cluster, but it is undergoing a merger. Here, we report on the analysis of the \\xmm\\ observations of this cluster.  The X-ray data give us the opportunity to reveal the complexity of the cluster, especially its temperature distribution.  The integrated spectrum is well fitted by a single temperature thermal model, indicating a mean temperature of $\\sim$7~keV. However, the cluster is not isothermal at this temperature: its eastern regions are significantly cooler, at $\\sim$5.5~keV, whilst towards the West the temperature reaches $\\sim$8.5~keV.    These temperature asymmetries can be explained if \\cl\\ has been assembled recently by the merging of smaller subunits. It is now in the phase after the cores of these subunits have collided (the `core-crossing' phase) some 0.1-0.2~Gyr ago. A comparison with numerical simulations suggests that it will settle down into a single relaxed cluster in $\\sim$(2-3)~Gyr. ",
        "introduction": "After the first X-ray observations, it became apparent that \\cl\\ is a cluster that is currently under formation, growing by the accumulation of smaller subunits. The \\rosat\\ images showed that the cluster is elongated along the East-West direction, and that the centroid of the X-ray emission does not coincide with any large cluster galaxy (Burns et all. 1995, Feretti et al. 1997). The spectral analysis of Davis \\& White (1998) found significant temperature structure in its intracluster medium (ICM), which they also attributed to a recent merger event. Abell~2255 was observed by \\chandra\\ with the ACIS-I detector for a total of 39~ksec. The \\chandra\\ data set, was used by Davis, Miller, \\& Mushotzky (2003), to investigate the X-ray properties of the cluster galaxies. It has also been observed by \\xmm\\, and Fig.~1(a) shows the \\xmm\\ mosaic. In Fig~1(b) we present an overlay of the X-ray contours onto an optical image of the central regions of \\cl\\. The X-ray contours are obtained from the \\xmm\\ observations that will be presented in the subsequent sections of this paper.    Optically, as can be seen in Fig.~1(b), a very intriguing property of \\cl\\ is that the brightest galaxies are arranged in a chain, whose orientation coincides with the major axis of the elliptical X-ray emission. The cluster has an unusually high velocity dispersion of $\\sim$1200~${\\rm km \\ s^{-1}}$, and the two brightest galaxies [galaxies A and B in Fig.~1(b)] are separated by $\\sim$2600~${\\rm km \\ s^{-1}}$ (Burns et al. 1995). Performing deep multicolour photometry in a large field around the cluster, Yuan, Zhou, \\& Jjang (2003) showed that at radii $>$(10-15)~arcmin there are small groups of galaxies, that appear to rotate around the central core of Abell~2255.    \\begin{table*} \\caption{Pointing Information}\\label{obs_info} \\begin{center} \\begin{tabular}{ccccccccc}   \\hline \\hline   (I)          & (II)            & (III)           & (IV)            & (V)             & (VI)            & (VII) \\\\  Rev             & Obs             & $\\alpha$~(2000)         & $\\delta$~(2000)         & Instr.                  & $Exp$           & $Exp_{corr}$            &  \\\\ & & & & & ksec            & ksec            &  \\\\ \\hline  525            & 0112260501      & 17 12 59.920        & +64 03 25.00        & MOS1                    & 8.558  & 4.307  & \\\\  & & & & MOS2    & 8.645  & 4.466  & \\\\  & & & & PN      & 4.326   & 1.665   & \\\\  548            & 0112260801      & 17 12 58.460                & +64 04 49.20              & MOS1                            & 16.468          & 10.810          & \\\\  & & & & MOS2    & 16.504  & 10.657  & \\\\  & & & & PN      & 11.918  & 4.068  & \\\\  \\hline  \\end{tabular} \\vspace{0.2cm} \\begin{minipage}{16cm} \\small NOTES : (I)-revolution number; (II)-observation number; (III)-pointing $Right Ascension$; (IV)-pointing $Declination$; (V) EPIC Instrument; (VI) Exposure time (live time for the central CCD); (VII) Reduced Exposure time, after the subtraction of the bright background flares (see text for more details). \\end{minipage} \\end{center} \\end{table*}   In the radio \\cl\\ contains a central radio halo (see, for example, Giovannini, Tordi \\& Feretti 1999, and references there-in), and a number of tailed radio galaxies. More recently, Govoni et al. (2005) presented a high sensitivity radio image of \\cl\\ which reveals the detailed structure of its radio halo and a possible radio relic. Radio halos are rare radio sources, and they have been found in the inner, 1 Mpc of X-ray bright and hot clusters [see Giovannini \\& Feretti (2000) for some recent examples]. They locate the site of relativistic electrons and magnetic fields in clusters. The presence of such a halo in Coma, in conjunction with its X-ray morphology, motivated the suggestion by Burns et al. (1994) that radio halos are fuelled by cluster collisions, and are associated with clusters undergoing disturbances from recent or on-going merging events.  Since the first discoveries, significant advances have been made in understanding their origin (e.g. Buote 2001).  Thus, the X-ray, optical and radio data provide evidence that \\cl\\ is currently active, and other surrounding structures might be interacting with it. This impression is not very surprising if one thinks that \\cl\\ is a member of the rich North Ecliptic Pole supercluster, that contains at least 21 galaxy clusters, as was revealed by the analysis of the \\rosat\\ All-Sky Survey data by Mullis et al. (2001).  We have observed \\cl\\ with \\xmm\\ in order to uncover its dynamical state, decide on its past history and future evolution, and derive vital information that would help us to test the results of the numerical simulations of merging clusters. In this paper, we present the analysis and results of the \\xmm\\ observations: the observations are described in Section~2; Sections 3 and 4 are devoted to the presentation of the X-ray properties of the cluster, as found from the \\xmm\\ data analysis; in Section~5 we compare the \\xmm\\ and \\rosat\\ results; finally, in Section~6 we discuss a possible dynamical scenario that can describe well the data.  The redshift of \\cl\\ is z=0.0806 (NED), and throughout this paper we use $H_0=71~{\\rm km \\ s^{-1} \\ Mpc^{-1}}$, $\\Omega_{\\rm M}$=0.3, and $\\Omega_{\\rm \\Lambda}$=0.7, giving a scale of 1.499 kpc/\". The Galactic hydrogen column for the direction of \\cl\\ is $N_{\\rm H,G} = 2.6\\times 10^{20}\\ {\\rm cm^{2}}$. ",
        "conclusions": "The \\xmm\\ data and analysis, presented in the previous sections show strong evidence that \\cl\\ is far from our idealized picture of a `relaxed' cluster. It has suffered a merger event in its recent past, and the signatures of such a turmoil are still visible: the X-ray emission is elongated along the East-West direction, aligned with a chain of galaxies; the temperature distribution does not follow the universal temperature profile of Loken et al. (2002); the X-ray emission is not centered at any galaxy.  Especially, the temperature structure shows the kind of disturbances expected during/after a merger event. Although we derive a `global' temperature for the cluster of $\\sim$6.9~keV, it is not isothermal at this temperature. Its eastern regions are cooler at $\\sim$(5-6)~keV, and towards the West the temperature reaches $\\sim$8.5~keV. Following the temperature asymmetry, there is a pressure imbalance between East and West (see Fig.~\\ref{P}), which is mainly driven by the temperature inequality.  Over the last years, thanks to the \\xmm\\ and \\chandra\\, we have witnessed the complicated structure of the ICM in merging/evolving clusters. Additionally, numerical simulations of merging clusters have advanced to such an extent that can show us the details of the evolution of a merger event. The challenge now is to match the observations with the simulations and understand how the final cluster is formed.  For unequal mass mergers, simulations show (e.g., Takizawa 1999, Roettiger et al. 1997) that after the cores of the two subclusters collide, the smaller and cooler cluster continues travelling way from the collision location, leaving a cool trail along its trajectory. Two shocks are generated during the collision, that travel along the collision axis towards opposite directions. One `front' shock is leading the smaller subcluster, and it is diffuse. The `back' one is more compact, and propagates towards the opposite direction.  The X-ray properties of \\cl\\ support the idea that it is a merger remnant after the core collision phase.  Specifically, there is a correspondence with the results, for example, presented by Takizawa (1999). A comparison of Fig.~\\ref{Tmap} with the temperature structure in the merger remnant of his fig.~8 at $t=4.75$~Gyr shows a striking similarity.  This correspondence can be also found with other numerical work.  The comparison of the temperature distributions along sector-E and -W of Fig.~\\ref{A2255T_prof_pies} shows that there is a temperature increase from the cluster average ($\\sim$6.9~keV) up to $\\sim$8.5~keV. A possible explanation for this is that it is due to a shock wave.  If this is the case, using eqn.~(2) from Markevitch, Sarazin \\& Vikhlinin (1999), we find that the Mach number of the relative motion is $M \\sim 1.24$, which implies a velocity of $\\simeq 2400 \\ {\\rm km \\ s^{-1}}$ (the sound speed in Abell~2255 is $c_s=1940\\ {\\rm km \\ s^{-1}}$), and a compression factor of 1/$x \\sim 1.36$. For the above calculation we used $\\gamma = 5/3$, a preshock temperature of k$T_0\\sim6.9$~keV, and postshock temperature of k$T_1\\sim8.5$~keV.  If this is the `back' shock seen in the simulations, and it has been travelling at $v=2400 \\ {\\rm km \\ s^{-1}}$, and if the core crossing happened where the current cluster centre is, we find that the cores collided some $t = s/v \\sim (4~{\\rm arcmin}) / (2400 \\ {\\rm km / s^{-1}}) \\sim 0.15~{\\rm Gyr}$ ago.  Such a shock wave would have increased the flux by an amount of $\\Sigma_1/\\Sigma_0=1.36^{2} \\sim 1.85$. If we compare the flux at a distance of $\\sim5$~arcmin from the cluster centre in Sector-W with the flux along sector-S and -N, we find that the measured $\\Sigma_1/\\Sigma_0$ is not more than $\\sim$2.  Of course, the details of the structure depends on the initial conditions and projection effects.  Being guided again by the numerical work [e.g., Takizawa (1999), Randall, Sarazin, Ricker (2002)] we note that the collision of the subclusters' cores during the `core-crossing' phase results in a sharp and brief increase of the luminosity and temperature. Afterwards, the cluster expands adiabatically, resulting in a decrease of its temperature and luminosity. The luminosity drops below its initial value, which is defined as the sum of the initial luminosities of the two subunits. The temperature on the other hand, remains at the same levels as the initial temperature (before the collision) until a later stage of the merging process, at which it increases more mildly, due to the collapse and accumulation of the cluster material towards the new cluster centre. Thus, if a cluster is at a stage after the dramatic `core-crossing' phase, its luminosity should be high, but not as high as during the violent sub-clusters collision, and lower than the initial total luminosity of the system.  In order to locate \\cl\\ on the L-T relation, we compare its properties with the L-T relation derived by Markevitch (1998). We calculated the cluster bolometric luminosity, and extrapolated it out to 1$h^{-1}$~Mpc, in the same manner as in Markevitch (1998). Although \\cl\\ does not host a traditional `cooling flow', we exclude the inner regions of the cluster to be consistent with the analysis of Markevitch (1998). The derived luminosity, $L_{\\rm bol} \\simeq 12.13 \\times 10^{44} {\\rm erg \\ s^{-1}}$, is almost double the expected $L_{\\rm bol}$ of a $\\sim$7~keV cluster, according to the L-T relation, if we compare with fig.~2 in Markevitch (1998).  During the next stages of the merging process, the luminosity of the remnant will remain almost unaltered, while its temperature will increase as explained earlier. This temperature increase may be such that its temperature and luminosity come into agreement with the L-T relation, and will stay at that condition until the final merger remnant is formed.  Being guided again by the the numerical simulations we find that \\cl\\ will settle down to a single remnant in some $\\sim$(2-3)~Gyr."
    },
    "0601/astro-ph0601457_arXiv.txt": {
        "abstract": "This paper provides a complete theoretical treatment of the point-mass lens perturbed by constant external shear, often called the Chang--Refsdal lens. We show that simple invariants exist for the products of the (complex) positions of the four images, as well as moment sums of their signed magnifications. The image topographies and equations of the caustics and critical curves are also studied. We derive the fully analytic expressions for pre-caustics, which are the loci of non-critical points that map to the caustics under the lens mapping. They constitute boundaries of the region in the image domain that maps onto the interior of the caustics. The areas under the critical curves, caustics and pre-caustics are all evaluated, which enables us to calculate the mean magnification of the source within the caustics. Additionally, the exact analytic expression for the magnification distribution for the source in the triangular caustics is derived, as well as a useful approximate expression. Finally, we find that the Chang--Refsdal lens with additional convergence greater than unity (the `over-focusing case') can exhibit third-order critical behaviour, if the `reduced shear' is exactly equal to $\\sqrt3/2$, and that the number of images for $N$ point masses with non-zero constant shear cannot be greater than $5N-1$. ",
        "introduction": "The so-called Chang--Refsdal lens was put forward to describe the lensing effects of stars in a background galaxy \\citep{CR79,CR84}. The star acts as a point-mass lens, while the galaxy provides a background perturbation field, which can be approximated by a tidal term to lowest order. The Chang--Refsdal lens is therefore the simplest description of lensing by a tidally perturbed point-mass. As such, it has found widespread astronomical applications -- e.g. in the modelling of binary microlensing lightcurves \\citep{GL92,GG97l}, in the analysis of the statistics of high-magnification events caused by stars in foreground galaxies \\citep{Sc87}, and in the study of microlensing of stars around the black hole in the Galactic Centre \\citep{Al01}, and so on.  The Chang--Refsdal lens is of outstanding physical importance, but it also has a number of elegant mathematical properties, primarily because the `deflection function' is a rational function of position in the lens plane. Although there has already been much work done on the structure of the caustics and critical curves \\citep{CR84,Su85,Sc92}, many of the properties of the Chang--Refsdal lens do not appear to be in the literature. The purpose of this paper is to give a compendium of those results.  The paper is organized as follows. Section \\ref{sec:LE} presents the Chang--Refsdal lens utilizing the complex notation popularized by \\citeauthor{Wi90} (\\citeyear{Wi90}; see also \\citealt{BK75}). In section \\ref{sec:IP}, the lens equation is converted into the imaging polynomial, the roots of which include the image locations. Provided that the source lies within the caustic -- so that the number of images is four -- then there exist invariants, such as the products of the image positions or the sums of the magnifications weighted by powers of the positions. Section \\ref{sec:SSA} considers two cases when the lens equation is solvable via elementary means, namely when the source lies on one of the axes of the symmetry. In section \\ref{sec:CCC}, the pre-caustics of the Chang--Refsdal lens are isolated, together with the critical curves and the caustics. This enables us to calculate in section \\ref{sec:LEM} the mean magnification of the 4-image configurations, and the exact form for a certain conditional distribution of magnifications. Finally, sections \\ref{sec:CB} and \\ref{sec:NPM} consider two generalizations of the Chang--Refsdal lens, namely the cases of a convergent background and of $N$ point masses with shear. We find that, with the convergence greater than unity, the number of images is either nil or two if the `reduced shear' is smaller than $\\sqrt3/2$, nil, two, or four if it is in between $\\sqrt3/2$ and the unity, and two or four if it is greater than unity. As for $N$ point masses with shear, we establish that the maximum number of possible images are bounded by $5N-1$. ",
        "conclusions": "This paper has presented a complete mathematical treatment of the Chang--Refsdal lens. This simple and flexible model has found widespread applications in astrophysics. The lens equation for the Chang--Refsdal lens can be converted to a polynomial of the fourth degree, the roots of which provide the image locations (together with possibly spurious roots). If the source lies on one of the symmetry axes with respect to the external shear, this quartic readily factorizes into a product of quadratics, which enables the lens equation to be solved via elementary means for these special cases. Simple lensing invariants exist for the products of the (complex) positions of the four images, as well as for sums of the moments of their signed magnifications. In other words, provided the source lies within the astroid or deltoid caustics, the sums of signed magnification, multiplied by moments of the complex positions, are invariant.  Of the results in this paper, we highlight just two here. First, we have provided the equations for the pre-caustics. Only one other lens model has been known for which the pre-caustics can be written down, namely the singular isothermal sphere with external shear \\citep{Fi02}. Second, the exact analytical expression (eq.~\\ref{eq:distdel}) in terms of hypergeometric functions (or elliptic integrals) has been calculated for the distribution of magnifications for the source within the deltoid caustics of the Chang--Refsdal lens. For many practical purposes, however, we have shown that a truncated power-law (eq.~\\ref{eq:imp}) is an excellent approximation to this distribution.  There are a number of applications of these two results. For example, the Chang--Refsdal lens is often used as an approximation to the binary lens equation, especially in the case when the lens is a star plus planet. The areas under the critical curves and pre-caustics are directly related to the mean magnifications of the 2 and 4 image geometries. and hence the detectability of a planet through the gravitational microlensing effect. Equally, the distribution of magnifications is applicable to the statistics of high-magnification microlensing events in the lightcurves of quasars."
    },
    "0601/astro-ph0601005.txt": {
        "abstract": "To study how supernova feedback structures the turbulent interstellar medium, we construct 3D models of vertically stratified gas stirred by discrete supernova explosions, including vertical gravitational field and parametrized heating and cooling. The models reproduce many observed characteristics of the Galaxy such as global circulation of gas (i.e., galactic fountain) and the existence of cold dense clouds in the galactic disk. Global quantities of the model such as warm and hot gas filling factors in the midplane, mass fraction of thermally unstable gas, and the averaged vertical density profile are compared directly with existing observations, and shown to be broadly consistent. We find that energy injection occurs over a broad range of scales. There is no single effective driving scale, unlike the usual assumption for idealized models of incompressible turbulence. However, $>$90\\% of the total kinetic energy is contained in wavelengths shortward of 200 pc. The shape of the kinetic energy spectrum differs substantially from that of the velocity power spectrum, which implies that the velocity structure varies with the gas density. Velocity structure functions demonstrate that the phenomenological theory proposed by Boldyrev is applicable to the medium. We show that it can be misleading to predict physical properties such as the stellar initial mass function based on numerical simulations that do not include self-gravity of the gas. Even if all the gas in turbulently Jeans unstable regions in our simulation is assumed to collapse and form stars in local freefall times, the resulting total collapse rate is significantly lower than the value consistent with the input supernova rate. Supernova-driven turbulence inhibits star formation globally rather than triggering it. ",
        "introduction": "The interstellar medium (ISM) is observed to be turbulent. The electron density spectrum in the warm ionized medium shows a nearly Kolmogorov spectrum over at least six orders of magnitude in length (Armstrong et al. 1995). Supposing the small density perturbations measured to behave as a passive scalar, this can be taken as evidence for a Kolmogorov velocity spectrum in that gas. Giant molecular clouds, which contain most of the molecular gas, show supra-thermal emission linewidths (Zuckerman \\& Palmer 1974), indicating the presence of supersonic random motions.  What drives this turbulence? Massive stars are a major structuring agent in the ISM (McCray \\& Snow 1979). Once formed inside molecular clouds, they influence the surrounding medium via intense ionizing radiation and stellar winds until they end their lives in catastrophic explosions as supernovae, heating and stirring the ISM throughout their lives. Among these energetic processes, supernovae (and their collective effect, superbubbles) are likely to be the dominant contributor to the observed supersonic turbulence (Norman \\& Ferrara 1996; see Mac Low \\& Klessen 2004 for a recent review). This turbulence changes the balance between pressure and gravity on a wide range of scales, from galactic (Boulares \\& Cox 1990) to sub-cloud scales (e.g., Larson 1981), and in turn may regulate subsequent star formation, hence the term ``feedback.'' With strong enough feedback, a galactic-scale superwind may be driven out of the galaxy (Heckman et al. 2000, 2001).  Our understanding of how the radiative and hydrodynamic feedback affects spatio-temporal structure of the ISM and star formation is only qualitative. It remains difficult to incorporate the feedback into a consistent theory of star formation and the ISM, partly because star formation is an inherently complex phenomenon that involves an interplay between gravity, magnetohydrodynamics, and thermodynamics of compressible gas. Yet it is a crucial physical process for correct predictions of, e.g., the thickness, disk stability, and star formation rate of galaxies. Many of the serious problems that cosmological models are faced with today are also believed to be associated with such feedback.\\footnote{The standard CDM paradigm predicts basic properties of the luminous component of matter on galactic and sub-galactic scales which are seriously inconsistent with observations. Outstanding astrophysical problems include incorrect predictions for the shape of the galaxy luminosity function and for the fraction of gas that should have cooled to form galaxies (the overcooling problem) (Somerville \\& Primack 1999; Cole et al. 2000).}  Since no analytic theory based on first principles exists for turbulence in a compressible medium (see, however, a heuristic model by Sasao 1973), most previous work has relied heavily on numerical methods. Recent numerical simulations in periodic boxes yielded key insights into the role of turbulence in star formation. First, supersonic turbulence in an isothermal or an adiabatic medium was found to decay quickly, i.e. within a sound crossing time, whether the medium was magnetized or not (Stone et al. 1998; Mac Low et al. 1998; Padoan \\& Nordlund 1999), unless it was constantly stirred by an external source of driving (Mac Low 1999). Second, supersonic turbulence delays or sometimes prevents collapse (Klessen et al. 2000, hereafter KHM00), and is hence thought to be responsible for the observed low overall star formation efficiencies (V\\'{a}zquez-Semadeni et al. 2005). Overall, turbulence impedes collapse in unstable regions. Third, in contrast to the second point, it also creates converging flows that may enhance density and thus promote collapse (Ballesteros-Paredes et al. 1999). Even when turbulence can support a cloud globally against gravitational collapse, it produces density enhancements that allow local collapse in both non-magnetized (KHM00) and magnetized media (Heitsch, Mac Low, \\& Klessen 2001, hereafter HMK01). However, the net effect of turbulence is to inhibit collapse (Mac Low \\& Klessen 2004).  In some of these previous simulations, turbulence was seeded as the initial condition and left to decay for the rest of the run (Porter et al. 1992; Porter et al. 1998; Stone et al. 1998; Mac Low et al. 1998; Padoan \\& Nordlund 1999). In other simulations, turbulence was driven constantly but the driving was included only in a general way as ubiquitous Fourier forcing, i.e., a random forcing in a prescribed narrow range of wavenumbers applied uniformly to the box (Mac Low 1999; KHM00; HMK01). These models only partly represent the real ISM, where forced and decaying regimes may coexist (Avila-Reese \\& V\\'{a}zquez-Semadeni 2001) and where an appropriate forcing function is probably broad-band (Norman \\& Ferrara 1996). It has been demonstrated that the slope of the power spectrum for the forcing function profoundly affects the statistics of turbulent fluctuations (Bonazzola et al. 1987; V\\'{a}zquez-Semadeni \\& Gazol 1995; Biferale et al. 2004). Hence, it is of great interest to determine the wavelength distribution of the kinetic energy in a more realistic medium driven by \\textit{discrete physical space forcing}.  Our goal is to examine physical characteristics of the ISM driven by realistic, discrete physical space forcing in 3D. In particular, we study density and velocity structures and determine the kinetic energy power spectrum. If, as some previous works suggest, the interaction of blast waves from supernovae structures the interstellar gas, we must resolve the individual explosions in order to correctly reproduce the main statistical properties of the ISM. For an accurate simulation of supernova driving, a wide range of length scales from small clouds ($\\sim$1 pc) to at least the ``driving scale'' (using the language derived from incompressible turbulence research) must be resolved. In addition, to predict the cloud structure, the model must take into account appropriate heating and radiative cooling processes.  There have been modeling efforts toward more realistic supernova driving by including its explosive nature. Rosen, Bregman, \\& Norman (1993), Rosen \\& Bregman (1995), and Wada \\& Norman (2001) performed two-dimensional hydrodynamic simulations for the ISM in a galactic disk including feedback from supernovae and/or stellar winds. Results from two-dimensional models must be interpreted with caution, since turbulence behaves differently in 2D than in 3D: in 2D, an inverse cascade of energy occurs towards large scales, while the enstrophy, the square of vorticity, cascades to small scales (Kraichnan 1967; Frisch 1995). More recently, three-dimensional models have been studied by several groups. These authors find that blast waves from supernovae sweep up the ISM into relatively thin shells that collide with one another to form cold dense clouds. However, none of the goals raised above has been adequately addressed. Korpi et al. (1999) used resolution of 10 pc, insufficient to follow the details of turbulent interactions, and only covered (0.5 kpc)$^2 \\times$ 2 kpc, although they did include magnetic field and galactic shear. Their main interest was to simulate a turbulent galactic dynamo. Slyz et al. (2005) included self-gravity of the gas but suffer from low spatial resolution ($\\ge$10 pc) and the absence of vertical density stratification. Kim et al. (2001) and Mac Low et al. (2005) also included magnetic field, but not galactic stratification. Avillez (2000) and Avillez \\& Berry (2001) used a stratified model similar to ours, but their studies focused on the properties of a Galactic fountain in the Milky Way. All the aforementioned 3D models employed minimum gas temperatures $T_{min} = 100$--310 K, so no cold gas could form. The difference between Avillez \\& Breitschwerdt (2004a, b) and the present work is more subtle, and will be discussed in \\S\\ \\ref{midplane}.  We simulate a small patch of our Galaxy, (0.5 kpc)$^2$ in area. The computation box is elongated in $z$ and extends from $-5$ kpc to $+5$ kpc to study the vertical structure. We choose the length scales in our simulations ($2$ pc $\\le l \\le$ $10$ kpc) to lie between those of star-forming molecular clouds and those of large-scale galactic outflows (Heckman et al. 2000, 2001). Our box size represents an optimal choice to improve our understanding of this enigmatic yet relatively unexplored regime. Expanding shells and their interactions are well-resolved at this scale. We attempt to build local, high-resolution models of the ISM based on which subgrid models for turbulent pressure can be developed, that will provide insight into how supernova feedback should be treated in global, cosmological simulations. We will continue to pursue this problem in a companion paper (M. Joung \\& M. Mac Low 2005, in preparation; hereafter Paper II). %in future work. Here, we Here, in Paper I, we describe our numerical model (\\S\\ \\ref{model}) and present the basic results (\\S\\ \\ref{results}). In \\S\\ \\ref{clouds}, we explore where gravitational collapse would occur in the presence of explosion-driven turbulence. Finally, we summarize our findings and discuss their implications in \\S\\ \\ref{discuss}. ",
        "conclusions": "\\label{discuss}  In this work, we have constructed 3D models of the ISM including vertical stratification and discrete physical space forcing by supernova explosions. We included parametrized heating and cooling as well as both isolated and correlated supernovae. This work supports models that treat supernova-driven turbulence as a major structuring agent in the ISM. Our main results are as follows:  \\begin{itemize} \\item Cold dense clouds naturally form as blast waves sweep the gas and collide. The intercloud region is filled with warm or hot gas having relatively low densities. The resulting density and velocity structures adequately match the observed filamentary and clumpy structure of the ISM.  \\item A galactic fountain reaching a height of several kiloparsecs forms, with most of the volume moving away from the midplane. Fragmented shells rain back down as small cold clumps, which may be identified as intermediate velocity clouds.  \\item Global quantities in our model are broadly consistent with existing observations. Specifically, our results are consistent with observations of the mass fraction of thermally unstable gas and the warm gas filling factor in the Galactic plane derived from H I absorption lines (Heiles 2001; Heiles \\& Troland 2003).  \\item The diffuse heating rate $\\Gamma$ influences the global structure of our model. For example, in our numerical experiments, increasing $\\Gamma$ by 4 lowered the occupation fraction of the hot gas in the Galactic plane by $\\sim$20\\%. This occurs because high background heating rates effectively shift the thermal equilibrium curve in the phase diagram so that the cold (thermally stable) branch lies mostly outside the range of pressure occupied by the bulk of the gas.  \\item The vertical distribution of gas in our model deviates from the profile inferred from observations. Specifically, the gas is excessively concentrated in the disk midplane by factors of 2--3, while it is deficient at $0.1 \\lesssim |z| \\lesssim 0.5$ kpc. We attribute this discrepancy to neglected components of pressure such as those from the magnetic field and cosmic rays.  \\item The density power spectrum has a broad peak close to the dissipation scale $\\lambda \\sim 20$ pc. It turns over and falls off at large scales.  \\item There is no single effective driving scale; instead, we see energy injection over a range of scales. As a result, the kinetic energy power spectrum possesses a slope much flatter than Kolmogorov's spectrum at scales between $\\sim$20 and $\\sim$500 pc. The predominance of small scale density structure causes the kinetic energy power spectrum to look remarkably different from the often-used velocity power spectrum. About 90\\% of the total kinetic energy is contained in wavelengths shortward of 190 pc.  \\item The velocity structure functions demonstrate that the phenomenological theory proposed by Boldyrev (2002) is applicable even to a medium driven by discrete point explosions and subject to nonlinear heating and cooling processes.  \\item Even if all the gas in turbulent Jeans unstable subboxes in our simulation collapses and forms stars in local freefall times, the resulting collapse rate is significantly lower than the value consistent with the input supernova rate. Supernova-driven turbulence inhibits, but does not prevent, star formation. From this, we infer that the density PDF must be significantly affected by self-gravity. For predicting physical properties of star-forming regions (such as the stellar initial mass function), the statistics of turbulent fluctuations generated by models without self-gravity can be misleading. %[Say something about log-normal?] \\end{itemize}  Applying high supernova rates to our model galaxy naturally leads to galactic outflows that have already been observed at both low (Heckman et al. 2000) and high redshifts (Pettini et al. 2001, 2002; Shapley et al. 2003). In the future, we will use these local models to measure the efficiency of energy transfer into superwinds out of a stratified disk for a given star formation rate, and thereby construct subgrid models for stellar energy feedback in cosmological simulations."
    },
    "0601/astro-ph0601277_arXiv.txt": {
        "abstract": "% I review recent progress from $N$-body simulations in our understanding of the secular evolution of isolated disk galaxies.  I describe some of the recent controversies in the field which have been commonly attributed to numerics.  The numerical methods used are widely used in computational astronomy and the problems encountered are therefore of wider interest. ",
        "introduction": "\\label{sec:intro}  The fragility of disk galaxies suggests that a significant fraction of their history was spent experiencing mild external perturbations. Thus their evolution is likely to have been driven partly by internal processes.  Dissipational gas physics probably plays a prominent role in this evolution but the effects of collisionless evolution very near equilibrium are also significant.  Collisionless secular evolution is driven largely by non-axisymmetric structures, especially by bars, which represent large departures from axisymmetry, and are thus efficient agents of mass, energy and angular momentum redistribution. Moreover bars occur in over $50\\%$ of disk galaxies \\citep{knapen_99, esk_etal_00} and simulations starting from the 1970s \\citep{mil_etal_70, hohl_71} have shown that they form readily via global instabilities \\citep{kalnaj_72,toomre_81}.  Here I review recent progress in understanding the secular evolution of isolated disk galaxies. ",
        "conclusions": ""
    },
    "0601/astro-ph0601088_arXiv.txt": {
        "abstract": "The use of Type Ia supernovae as cosmological tools has reinforced the need to better understand these objects and their light curves.  The light curves of Type Ia supernovae are powered by the nuclear decay of $^{56}Ni \\rightarrow ^{56}Co \\rightarrow ^{56}Fe$.  The late time light curves can provide insight into the behavior of the decay products and their effect of the shape of the curves.  We present the optical light curves of six ``normal\" Type Ia supernovae, obtained at late times with template image subtraction, and the fits of these light curves to supernova energy deposition models. ",
        "introduction": "Type Ia Supernovae (SNe Ia) are thought to be the thermonuclear explosion of a white dwarf \\citep[see][and references therein]{2000A&ARv..10..179L}.  The light curves of SNe Ia are powered by deposition in the SN ejecta of the $\\gamma$-ray and positron products of the $^{56}Ni\\rightarrow ^{56}Co\\rightarrow ^{56}Fe$ decay \\citep*{1969ApJ...157..623C}. The extreme brightness and seemingly uniform light curves of SNe Ia make them good candidates for use as standard candle distance indicators. In more recent years, it has been shown that Type Ia supernovae do not have uniform light curve magnitudes, shape or spectra.  The light curves can, however, be normalized to account for this inhomogeneity, thus allowing these objects to be used at standard candle distance indicators \\citep[e.g.][]{1993ApJ...413L.105P,1996ApJ...473...88R}.    Between 100-200 days after the explosion the ejecta become transparent to the $\\gamma$-rays and the light curve is powered by the deposition of the positron kinetic energy into the ejecta.  The escape of a fraction of these positrons from the ejecta has been suggested as a possible source of the Galactic 511 keV annihilation radiation \\citep*{1999ApJS..124..503M}.  There are currently two methods of modeling the late emission of SNe Ia.  One is radiation transport with complete and instantaneous trapping of the positrons.  \\cite{1980PhDT.........1A} showed, by comparing a model to the late time spectra of SN 1972E, that the ejecta will cool leading to an increased fraction of the emission coming out in the infrared, the so named ``infrared catastrophe\" (IRC).  Other studies of radiation transport \\citep[e.g.][]{1996ssr..conf..211F,2004A&A...428..555S}, have reproduced the IRC, but they also predict the abrupt fall off of the optical light curves as the emission shifts into the NIR and ultimately into the IR, which is not seen in observed light curves.  The other method consists of positron energy deposition modeling without radiation transport.  In this type of modeling \\citep[e.g.][]{1980ApJ...237L..81C,1997A&A...328..203C,1998ApJ...500..360R,1999ApJS..124..503M} , optical band light curves are used as tracers of the bolometric luminosity and fit to model energy deposition curves.  The results show model curves with varying degrees of positron escape fitting the light curves.  One weakness in this model fitting technique is in using the optical bands as tracers of bolometric.  \\citet{2001ApJ...559.1019M} constructed bolometric curves using BVRI bands and showed those curves roughly fitting the positron escape energy deposition curves. ",
        "conclusions": "Figure \\ref{radmodel} shows the combined light curves of the six SNe plotted on the radiation transport models of \\cite{2004A&A...428..555S}.  The V-band model light curve has been normalized to be zero magnitude at 200d along with the data.  The B, R, \\& I band model light curves have been adjusted so that the colors of the model are preserved.  The dotted curve is the model including photoionization and the dot-dashed curve is the model without photoionization.   Figure \\ref{posmodel} shows the combined light curves of the six SNe plotted on the positron energy deposition curves of \\cite{2001ApJ...559.1019M}.  The model curves have been normalized to be zero magnitude at 200d.  The solid curve is the energy deposition with the positron kinetic energy trapped and deposited into the ejecta and the dashed curve is the energy deposition curve with a radially combed magnetic field allowing a fraction of the positrons to escape the ejecta without depositing their kinetic energy.  As can be seen in these figures, the shape of the B, V, \\& R band light curves could be explained by either the color evolution in the radiation transport model or the escape of positrons from the ejecta, while the I-band has a slower decline rate than both models.  This suggests that a model combining radiation transport with positron transport would be preferred to attempt to explain the late light curves of SNe Ia.  One thing is clear from the light curves; these SNe show very little deviation from each other in a given band, implying that within this class of normally-luminous SNe Ia there is only one answer for the question of positron escape from SNe Ia ejecta.     \\begin{figure} \\centerline{\\includegraphics[scale=0.6]{fig2.ps}} \\caption{BVRI Light curves of 6 SNe Ia plotted on the radiation transport model light curves of \\citet{2004A&A...428..555S}.\\label{radmodel}} \\end{figure}  \\begin{figure} \\centerline{\\includegraphics[scale=0.6]{fig3.ps}} \\caption{BVRI Light curves of 6 SNe Ia plotted on the postiron transport energy deposition curves of \\citet{1999ApJS..124..503M}\\label{posmodel}} \\end{figure}"
    },
    "0601/astro-ph0601549.txt": {
        "abstract": "Far-infrared continuum data from the {\\it COBE}/{\\it DIRBE} instrument were combined with Nagoya 4-m $\\cOone$ spectral line data to infer the multiparsec-scale physical conditions in the Orion$\\,$A and B molecular clouds, using 140$\\um$/240$\\um$ dust color temperatures and the 240$\\um$/$\\cOone$ intensity ratios. In theory, the ratio of far-IR, submillimeter, or millimeter continuum to that of a $\\cO$ (or $\\Co$) rotational line can place reliable upper limits on the temperature of the dust and molecular gas on multi-parsec scales; on such scales, both the line and continuum emission are optically thin, resulting in a continuum-to-line ratio that suffers no loss of temperature sensitivity in the high-temperature limit as occurs for ratios of CO rotational lines or ratios of continuum emission in different wavelength bands.  Two-component models fit the Orion data best, where one has a fixed-temperature and the other has a spatially varying temperature.  The former represents gas and dust towards the surface of the clouds that are heated primarily by a very large-scale (i.e. $\\sim 1\\,$kpc) interstellar radiation field.  The latter represents gas and dust at greater depths into the clouds and are shielded from this interstellar radiation field and heated by local stars.  The inferred physical conditions are consistent with those determined from previously observed maps of $\\COone$ and $\\Jtwo$ that cover the entire Orion$\\,$A and B molecular clouds.  The models require that the dust-gas temperature difference is 0$\\pm 2\\,$K.  If this surprising result applies to much of the Galactic ISM, except in unusual regions such as the Galactic Center, then there are a number implications.  These include dust-gas thermal coupling that is commonly factors of 5 to 10 stronger than previously believed, Galactic-scale molecular gas temperatures closer to 20$\\,$K than to 10$\\,$K, an improved explanation for the N(H$_2$)/I(CO) conversion factor (a full discussion of this is deferred to a later paper), and ruling out at least one dust grain alignment mechanism.  The simplest interpretation of the models suggests that about 40--50\\% of the Orion clouds are in the form of cold (i.e. $\\sim 3$-10$\\,$K) dust and gas, although alternative explanations are not ruled out. These alternatives include the contribution to the 240$\\um$ continuum by dust associated with atomic hydrogen and reduced $\\cO$ abundance towards the clouds' edges.  Even considering these alternatives, it is still likely that cold material with temperatures of $\\sim 7$-10$\\,$K still exists.  If this cold gas and dust are common in the Galaxy, then mass estimates of the Galactic ISM must be revised upwards by up to 60\\%.  The feasibility of submillimeter or millimeter continuum to $\\cO$ line ratios constraining estimates of dust and molecular gas temperatures was tested.  The model fits allowed the simulation of the necessary millimeter-continuum and $\\cOone$ maps used in the test.  In certain ``hot spots\" --- that have continuum-to-line ratios above some threshold value --- the millimeter continuum to $\\cO$ ratio can estimate the dust temperature to within a factor of 2 over large ranges of physical conditions.  Nevertheless, supplemental observations of the $\\cOtwo$ line or of shorter wavelength continuum are advisable in placing lower limits on the estimated temperature.  Even without such supplemental observations, this test shows that the continuum-to-line ratio places reliable upper limits on the temperature.  %I'll put something really revolutionary here after I've done the rest. ",
        "introduction": "}  While interesting in themselves, molecular clouds provide insights into star formation. Since stars form in and from molecular clouds, knowing the physical conditions within these clouds is essential for a complete understanding of star formation.  As mentioned in Paper~I \\citep{W05}, the warm (i.e., $\\gsim 50$--100$\\,$K) molecular gas associated with star formation is often identified and diagnosed from observations of different rotational lines of CO \\citep[e.g.,][]{Wilson01, Plume00, Howe93, Graf93, Graf90, Boreiko89, Fixsen99, Harris85, Harr99, W91, Gus93, Wild92, Harris91}.   Molecular gas, and the interstellar medium in general, can also be observed in the millimeter, submillimeter, and far-IR continuum, which trace the emission of the dust grains associated with interstellar gas.  Continuum surveys can probe the structure and excitation of the ISM \\citep[see, for example,][]{Dupac01, W96, Bally91, Zhang89, Werner76, Heiles00, Reach98, Boulanger98, Lagache98, Goldsmith97, Sodroski94, Boulanger90, Sellgren90, Scoville89, Sodroski89, Leisawitz88}.  Estimating physical parameters like temperature, and sometimes density, requires using the ratios of intensities of spectral lines or of the continuum at different wavelengths.  Given that each of these ratios is dependent on the ratio of two Planck functions at two different wavelengths, they often lose temperature sensitivity at higher temperatures. While there are methods of addressing this shortcoming, having two tracers of molecular gas with different dependences on the temperature would complement other methods of tracing warm dust or molecular gas.  This is especially true if the tracers are optically thin, because low opacity emission is more sensitive to the physical parameters of the bulk of the gas, rather than in just the surface layers.  One such pair of tracers is a rotational line of an isotopologue of CO, such as that of $\\cO$ or $\\Co$, and the submillimeter continuum.  Both of these tracers are optically thin on the scales of many parsecs, which are the scales of interest for the current work.  \\citet{Schloerb87} and \\citet{Swartz89} showed that the intensity ratio of an optically thin isotopic CO line emission to submillimeter continuum emission can estimate the temperature of gas and dust in molecular clouds. The \\citet{Schloerb87} expression for this ratio goes roughly like T$^2$ in the high-temperature limit.  Accordingly, the $\\rm I_\\nu(submm)/I(\\Coone)$ and $\\rm I_\\nu(submm)/I(\\cOone)$ ratios are actually {\\it more sensitive to temperature as that temperature increases.\\/}  This is in stark contrast to ratios of rotational lines of a given isotopologue of CO and to ratios of continuum intensities at different frequencies, which {\\it lose\\/} sensitivity to temperature in the high-temperature limit. The $\\rm I_\\nu(submm)/I(\\cO)$ ratio can then serve as the needed diagnostic of high gas/dust temperatures, provided that the shortcomings and complications of these tracers can be overcome or at least mitigated.  These complications include variations in the $\\cO$-to-dust mass ratio, non-molecular phases of the ISM along the line of sight, variations of gas density, variations in dust grain properties, appreciable optical depth variations in the $\\cO$ line used, and others (see the Introduction of Paper~I for more details).  Such complications are often reduced in the case of observations on multi-parsec scales, because spatial gradients on such scales are generally smaller than the extremes that occur on very small scales.  Consequently, testing the reliability of the $\\rm I_\\nu(submm)/I(\\cO)$ ratio as high-temperature diagnostic is best carried out with observations of a molecular cloud, or of clouds, on multi-parsec scales.  The Orion~A and B molecular clouds were chosen as the clouds for testing the\\hfil\\break $\\rm I_\\nu(submm)/I(\\cO)$ ratio's diagnostic ability.  They have been mapped in the $\\cOone$ line \\citep[see, e.g.,][]{Nagahama98} and in the far-IR by {\\it IRAS\\/} \\citep{Bally91} and {\\it COBE/DIRBE\\/} \\citep[][W96 hereafter]{W96}.  Avoiding the complication of the emission of stochastically heated dust grains requires far-IR observations at wavelengths longer than 100$\\um$ \\citep[e.g.][W96]{Desert90}. Accordingly, the far-IR observations of {\\it COBE/DIRBE\\/} were used instead of {\\it IRAS\\/} because the former has two bands --- $\\lambda = 140\\um$ and 240$\\um$ --- longward of 100$\\um$, whereas the latter does not.  The Orion~A and B clouds were chosen for this study because they have the advantages that they are bright in $\\cOone$ and at far-IR wavelengths, are out of the Galactic plane to avoid confusion with foreground and background emission, are several degrees in size so as to accommodate many {\\it DIRBE\\/} beams, and have the best range of dust temperatures at the {\\it DIRBE\\/} resolution of $0\\degree.7$ \\citep[see the Introduction of Paper~I and][for more details]{dirbex}. Therefore, the $\\Ib/\\Ic$ ratio, hereafter called $\\rd$, was plotted against the 140$\\um$/240$\\um$ dust color temperature, or $\\Tdc$, to test the $\\rd$'s ability to recover molecular cloud physical conditions.  Physical models were applied to these data and physical conditions were inferred in Paper~I.  The reliability of the model results were tested with simulated data in Paper~II \\citep{W05a}.  The next section summarizes the model results (i.e. Paper~I) and the results of the simulations (i.e. Paper~II).   Section~\\ref{ssec38} then discusses general systematic effects that had not been treated previously.  Section~\\ref{sec4} gives the scientific implications of the results. ",
        "conclusions": "}  Far-infrared continuum data from the {\\it DIRBE\\/} instrument aboard the {\\it COBE\\/} spacecraft were combined with $\\cOone$ spectral line data from the Nagoya 4-m telescope to infer the large-scale (i.e. $\\sim 5$ to $\\sim 100\\,$pc) physical conditions in the Orion molecular clouds.  The 140$\\um$/240$\\um$ dust color temperatures, $\\Tdc$, were compared with the 240$\\um$/$\\cOone$ intensity ratios, $\\rd$, to constrain dust and molecular gas physical conditions.  In addition, such a comparison provides valuable insights into how the ratio of FIR/submillimeter/millimeter continuum to that of a $\\cO$ (or $\\Co$) rotational line can constrain temperature estimates of the dust and molecular gas.  For example, ratios of rotational lines or ratios of continuum emission in different wavelength bands often cannot place realistic upper limits on gas or dust temperature, whereas the continuum-to-line ratio can place such limits.  Two-component models fit the Orion data best.  One component has a fixed-temperature and represents the gas and dust towards the surface of the clouds and is heated primarily by a kiloparsec-scale interstellar radiation field, referred to here as the general ISRF.  The other component has a spatially varying temperature and represents gas and dust towards the interior of the clouds that can be both shielded from the general ISRF and heated by local stars.  The model results and their implications are as follows: \\begin{itemize} \\item[1)] The inferred physical conditions are consistent with those derived from the large-scale observations of the $\\Jtwo$ and $\\Jone$ lines of $\\CO$ by \\citet{Sakamoto94}. \\item[2)] At least two gas (dust) temperatures are needed on the majority of sightlines through molecular clouds for reliably estimating column densities.   This is supported by the work of \\citet{Schnee06}. \\item[3)] The dust-gas temperature difference, $\\Td-\\Tk$ or $\\DT$, is 0$\\,$K to within 1 or 2$\\,$K.  If this result applies more generally to the Galactic-scale molecular ISM, except for unusual regions such as the Galactic Center, then there are a number of implications: \\begin{itemize} \\item Dust-gas thermal coupling is factors of 5 to 10 stronger than has been previously assumed.  Such factors may be due to the distribution of dust grain sizes and grains with larger cross-section to volume ratios than that of a simple sphere. \\item Galactic-scale molecular gas temperatures are closer to 20$\\,$K than to 10$\\,$K, because the emission from the CO rotational lines, even the optically thick $\\COone$ line, does not fill the beam within the velocity interval about the line peak. \\item This CO emission that does not fill the beam provides a better explanation of the N(H$_2$)/I(CO) conversion factor or X-factor.  Discussion of this is deferred to a later paper \\citep{W05b}. \\item Having $\\DT$ nearly 0 constrains which mechanisms explain dust grain alignment in the ISM.  A negligible dust-gas temperature difference rules out the Davis-Greenstein alignment mechanism, but not other possible mechanisms \\citep[see][and references therein]{Lazarian97, Abbas04}. \\end{itemize} \\item[4)] Roughly 40--50\\% of the ISM in Orion is cold (i.e. 10$\\,$K) to very cold (i.e. down to 3$\\,$K) dust and gas.  Accordingly, there is roughly 60\\% more gas and dust in Orion than inferred from simple one-component models.  This {\\it may\\/} also imply a similar increase in the estimated mass of entire Galactic ISM.  Fractal dust grains \\citep[see][and references therein]{Reach95} or iron needles \\citep{Dwek04} may explain the low temperatures of this gas and dust  and, at the same time, may account for the high dust-gas thermal coupling needed to explain $\\DT\\simeq 0\\,$K.  Nevertheless, alternative explanations that do {\\it not\\/} require cold dust and gas cannot be ruled out; the least unlikely of these other explanations is a contribution to the 240$\\um$ continuum emission of the dust associated with atomic hydrogen.  The data suggest that the effect of the H$\\,$I-associated dust is negligible, but still might permit raising the lower temperature limit of this cold gas and dust from 3 to 5$\\,$K. %% even though cold dust and gas represent the simplest interpretation of the models. \\end{itemize}  The model parameter values derived from the fits to the $\\rd$ versus $\\Tdc$ plots were used to create simulated 1300$\\um$ continuum and $\\cOone$ line maps.  These simulated maps tested whether the millimeter continuum to $\\cOone$ line intensity ratio could constrain temperature estimates of the dust and molecular gas. The ratio $\\rm I_\\nu(1300\\um)/I(\\cOone)$, or $\\rdl$, was found to estimate the dust temperature to within a factor of 2 in most cases, provided that $\\rdl$ was higher than a threshold level of 0.5$\\MJkk$.  Supplemental observations of the $\\cOtwo$ line and shorter wavelength continuum would confirm the high temperatures in these high-$\\rdl$ ``hot spots\".   The results here can be easily generalized to other continuum wavelengths and other rotational lines, even permitting interpretation of millimeter and submillimeter molecular line {\\it equivalent widths.\\/}  And this is entirely possible with only ground-based observations.  The full potential of using millimeter continuum and $\\cO$ (or $\\Co$) rotational line comparisons has yet to be realized.     %% The \\notetoeditor{TEXT} command allows the author to communicate %% information to the copy editor.  This information will appear as a %% footnote on the printed copy for the manuscript style file.  Nothing will %% appear on the printed copy if the preprint or %% preprint2 style files are used.  %% The eqnarray environment produces multi-line display math. The end of %% each line is marked with a \\\\. Lines will be numbered unless the \\\\ %% is preceded by a \\nonumber command. %% Alignment points are marked by ampersands (&). There should be two %% ampersands (&) per line.    %% Putting eqnarrays or equations inside the mathletters environment groups %% the enclosed equations by letter. For instance, the eqnarray below, instead %% of being numbered, say, (4) and (5), would be numbered (4a) and (4b). %% LaTeX the paper and look at the output to see the results.   %% If you wish to include an acknowledgments section in your paper, %% separate it off from the body of the text using the"
    },
    "0601/astro-ph0601107_arXiv.txt": {
        "abstract": "This paper generalises the hybrid power spectrum estimator developed in Efstathiou (2004a) to the estimation of polarization power spectra of the cosmic microwave background radiation.  The hybrid power spectrum estimator is unbiased and we show that it is close to optimal at all multipoles, provided the pixel noise satisfies certain reasonable constraints. Furthermore, the hybrid estimator is computationally fast and can easily be incorporated in a Monte-Carlo chain for {\\it Planck}-sized data sets. Simple formulae are given for the covariance matrices, including instrumental noise, and these are tested extensively against numerical simulations. We compare the behaviour of simple pseudo-$C_\\ell$ estimates with maximum likelihood estimates at low multipoles. For realistic sky cuts, maximum likelihood estimates reduce very significantly the mixing of $E$ and $B$ modes. To achieve limits on the scalar-tensor ratio of $r \\ll 0.1$ from sky maps with realistic sky cuts, maximum likelihood methods, or pseudo-$C_\\ell$ estimators based on unambiguous $E$ and $B$ modes, will be essential.   \\vskip 0.1 truein  \\noindent {\\bf Key words}: Methods: data analysis, statistical; Cosmology: cosmic microwave background, large-scale structure of Universe   \\vskip 0.3 truein ",
        "introduction": "In an earlier paper (Efstathiou 2004a, hereafter E04) a hybrid power spectrum estimator was developed for temperature anisotropies of the cosmic microwave background (CMB) anisotropies. This estimator combined quadratic maximum likelihood (QML) estimates at low multipoles with a set of pseudo-$C_\\ell$ (PCL) estimates  at higher multipoles to produce a near optimal power spectrum estimate over the entire multipole range for realistic sky coverages, scanning patterns and instrument noise.  We used analytic arguments and large numbers of numerical simulations to demonstrate that the method was near optimal and that the covariance matrix over the full range of multipoles could be estimated simply and accurately.  A full maximum-likelihood estimator would require the inversion and multiplication of $N_d \\times N_d$ matrices, where $N_d$ is the size of the data vector.  For {\\it WMAP} (Bennett \\etals 2003a) or {\\it Planck} (Bersanelli \\etals 1996; The Planck Consortia 2005) sized data sets, with $N_d \\simgt 10^6 -10^7$, a brute force application of ${\\cal O} (N_d^3)$ methods is impossible computationally.  This has been portrayed as a major computational challenge in CMB data analysis ({\\it e.g.} Bond \\etals 1999; Borrill 1999) and various approximate, but computationally demanding, techniques for solving this problem have been proposed ({\\it e.g.} Oh, Spergel and Hinshaw 1999, Dor\\'e, Knox and Peel (2001), Pen 20003). One of the major motivations for E04 was to show that an ${\\cal O} (N_d^3)$ computation is unnecessary and that a near-optimal estimator (essentially indistinguishable from an exact ${\\cal O} (N_d^3)$ maximum-likelihood solution), with a calculable covariance matrix, could be estimated simply by combining PCL and QML estimates. Furthermore, the analysis of realistic experiments must deal with uncertainties such as correlated instrument noise, beam calibrations, foreground separation, point sources {\\it etc}.  It is simply not worth performing an exact ${\\cal O} (N_d^3)$ power spectrum analysis if the assumptions under which it is optimal are violated by `real world' complexities, and if these complexities cannot be folded into the error estimates.  With a fast hybrid estimator it is feasible to assess such `real-world' complexities using a hybrid estimator within a Monte-Carlo chain and to quantify any corrrections to the covariance matrix, or likelihood function. The ability to analyse such `real-world' effects is likely to be much more important than any hypothetical marginal improvement that might be gained by applying a full ${\\cal O} (N_d^3)$ method. We refer the reader to E04 for a discussion of the rationale for a hybrid estimator. The arguments will not be repeated here. Instead, we will focus on generalising E04 to the estimation of power spectra of the CMB polarization anisotropies.   As is well-known, the Thomson scattering of an anisotropic photon distribution leads to a small net linear polarization of the CMB anisotropies (for an introductory review and references to earlier work see Hu and White 1997). This polarization signal can be decomposed into scalar $E$-modes and pseudo-scalar $B$-modes. The separation of a polarization pattern into $E$ and $B$ modes is of particular interest since scalar primordial perturbations generate only $E$ mode while tensor perturbations generate $E$ and $B$ modes of roughly comparable amplitudes (Zaldarriaga and Seljak 1997; Kamionkowski, Kosowsky and Stebbins 1997). The detection of an intrinsic $B$-mode signal in the CMB would provide incontrovertible evidence that the Universe experienced an inflationary phase and would fix the energy scale of inflation (see {\\it e.g.} Lyth 1984).  An $E$-mode polarization signal was first discovered by DASI\\footnote{ Degree Angular Scale Interferometer} (Kovac \\etals 2002, Leitch \\etals 2004).  Exquisite measurements of the temperature-$E$-mode (TE) cross power spectrum have been reported by the WMAP\\footnote{Wilkinson Microwave Anisotropy Probe} team (Kogut \\etals 2003). Measurements of the $E$-mode power spectrum have been reported by the CBI\\footnote{Cosmic Background Imager} experiment (Readhead \\etals 2004) and by the 2003 flight of Boomerang (Montroy \\etals 2005). Primordial $B$-mode anisotropies have not yet been detected in the CMB. The detection of $E$-mode and possibly $B$-mode anistropies is one of the main science goals of the {\\it Planck} mission and of ground based experiments such as {\\it Clover} (Taylor \\etals 2004).  Associated with these experimental developments, there have been many investigations of techniques for analysing $E$ and $B$ modes from maps of the CMB sky. These analyses can be grouped, approximately, into the following catagories:  \\begin{itemize} \\item[(i)] {\\it PCL estimators and correlation functions:} Fast methods of estimating $E$ and $B$-mode power spectra using correlation functions are described by Chon \\etals (2004). Statistically equivalent PCL estimators are described by Kogut \\etals (2003), Hansen and Gorski (2003), Challinor and Chon (2005, hereafter CC05) and Brown \\etals (2005).  In particular, CC05 present analytic approximations to the covariance matrices of PCL estimates for the case of noise-free data, while Brown \\etals (2005) develop a Monte-Carlo method for calibrating covariance estimates from incomplete maps of the sky. PCL estimators can be evaluated using fast spherical transforms and hence scale as ${\\cal O} (N_d^{3/2})$.  \\item[(ii)] {\\it Maximum Likelihood Methods:} The generalization of the iterative maximum likelihood power spectrum estimation methods of Bond \\etals (1998) to the analysis of polarisation is straightforward and will not be discussed further here. Tegmark and de Oliveira-Costa (2001, hereafter TdO01) define a quadratic estimator which is based on assumed forms for the temperature and polarization power spectra, generalising earlier work of Tegmark (1997). This method, which we will refer to as QML, is equivalent to a maximum likelihood solution if the guesses for the power spectra are close to their true values. As explained above, these methods involve matrix inversions and multiplications which scale as ${\\cal O} (N_d^{3})$.  \\item[(iii)] {\\it Harmonic $E$ and $B$ Mode Decomposition:} The Stokes' parameters $Q$ and $U$ describing linear polarization define a rank two symmetric trace-free polarization tensor on sphere. Over the complete sky, the polarization tensor can be decomposed uniquely into $E$ and $B$ modes. However, since the $E$ and $B$ decomposition is non-local, it is non-unique in the presence of boundaries. In any realistic situation, a sky cut must be imposed to exclude contamination of the CMB signal by high levels of Galactic emission at low Galactic latitudes. Various authors have discussed ways of detecting pure $E$ and $B$ modes from $Q$ and $U$ maps on an incomplete sky (Lewis \\etals 2002; Bunn \\etals 2003; Bunn 2003; Lewis 2003). A key motivation for these analyses has been for diagnostic purposes ({\\it e.g.} checking for a frequency dependent Galactic $B$-mode signal). However, as this article was nearing completion an interesting paper appeared by Smith (2005) describing how to construct pseudo-$C_\\ell$ estimators from unambiguous $E$ and $B$ modes on a cut sky. This type of analysis can effectively eliminate the severe $E$ and $B$ mode mixing that afflicts simple pseudo-$C_\\ell$ estimators at low multipoles (see Section 3.4). (The method is not completely straightforward, nor unique, because it requires a `pre-estimation' step to define unambiguous modes. We will throughout this paper use the abbreviation PCL to refer to the simple pseudo-$C_\\ell$ estimators as defined in Section 2.)  \\end{itemize}  As we will show in this paper, these three analysis techniques are closely related. To illustrate how these methods are interelated it is useful to consider separately the cases of noise-free and noisy data:  \\vskip 0.05 truein  \\noindent $\\bullet$ {\\it Comparison of PCL and QML estimators for noise-free data:} On a complete sky with noise-free data, PCL and QML estimates are identical. However, a sky cut will couple CMB modes over a range of multipoles $\\Delta L$, and this will lead to mixing of $E$ and $B$ pseudo-multipoles defined over the incomplete sky (see Section 2). Although PCL power spectrum estimators can be defined which give unbiased estimates of the true $E$ and $B$ mode power spectra {\\it in the mean}, the mixing of $E$ and $B$ modes is reflected in large variances and couplings between the estimated power spectra at multipoles $\\ell \\simlt \\Delta L$ (Section 2). A QML estimator applied to noise-free data unscrambles the $E$- and $B$-mode multipoles, in much the same way as the direct modal decompositions described by Lewis \\etals (2002), Lewis (2003) and Smith (2005). The QML estimator returns almost optimal power spectrum estimates with smaller variances than a PCL estimator (Section 3). In particular, for realistic sky cuts an estimator that minimises $E$ and $B$ mode mixing, such as the QML estimator, is essential if one wants to probe low amplitude $B$-mode signals (tensor-scalar ratios $r \\ll 0.1$). At high multipoles, $\\ell \\gg \\Delta L$, $E$ and $B$ mode mixing becomes unimportant and, in the case of noise free data, QML and PCL estimators become statistically equivalent (Section 3).  \\vskip 0.05 truein  \\noindent $\\bullet$ {\\it Comparison of PCL and QML estimators for noisy data:} If the following conditions are satisfied: (a) the $Q$ and $U$ maps have identical noise properties; (b) the noise is uncorrelated with variance per pixel $\\sigma^2_i$; (c) the estimates of $E$ and $B$ mode power spectra are for multipoles higher than some characterstic multipole $L_N$, where noise dominates; then one can show that an inverse-variance weighted PCL estimator is statistically equivalent to a QML estimator (Section 4.2). For the intermediate multipoles, $ L_N \\simlt l \\simlt \\Delta L$, the optimal weighting for PCL estimates is intermediate between equal weight per pixel and inverse variance weighting. By combining a set of PCL estimates with different weights, it is possible to define a fast polarization estimator (in analogy with the temperature estimator discussed by E04) that is statistically indistinguishable from a maximum likelihood estimator over a wide range of multipoles.  Dealing with strongly correlated noise is more problematic. Fortunately, for {\\it Planck}-type scanning strategies, the noise pattern (in $Q$, $U$ and $T$) should be accurately white at high multipoles (Efstathiou 2005).  A detailed correlated noise model should therefore only be required at low multipoles. It is straightforward to include a model for correlated noise in a QML estimate at low multipoles, provided the pixel noise covariance matrix can be estimated for low resolution maps (Section 4.2).  The layout of this paper is similar to that for the temperature analysis presented in E04. PCL polarization estimates are discussed in Section 2 together with analytic approximations for covariances matrices in the absence of instrumental noise. The relationship of this work to the results presented in CC05 is discussed. Section 3 discusses QML polarization estimators for noise free data. A large set of numerical simulations is used to illustrate the effects of $E$ and $B$ mode mixing on an incomplete sky for both PCL and QML estimators. Section 4 discusses PCL and QML estimators including instrumental noise and Section 5 discusses a hybrid polarization power spectrum estimator. To keep the discussion simple, most of the analytic and numerical results presented in this paper refer to $E$ and $B$ mode power spectrum estimation. Our discussion of the temperature-$E$ mode cross power spectrum (denoted $C^X$ in this paper) is, intentionally, less complete since no new concepts are required for its analysis. Our conclusions are summarized in Section 6. ",
        "conclusions": "This paper generalises the analysis of power spectrum estimators presented in E04 to the case of polarisation. Simple analytic expressions involving the power spectrum of the square of the window function (the scalar approximation) are given for the covariance matrices of PCL estimates, including the effects of uncorrelated instrument noise. For noise free data, the scalar approximation is shown to give accurate error estimates for the $E$-mode power spectrum at high multipoles for realistic (Kp0-like) sky-cuts, and also for the $B$-mode power spectrum provided the Q and U maps are apodised appropriately and the amplitude of the $B$-mode is high enough. For noise dominated data, the scalar approximation povides extremely accurate covariance estimates for both the $E$ and $B$ modes at high multipoles.  The results presented in Section 3 establish certain relationships between QML and PCL estimators, in particular, the statistical equivalence of QML and PCL estimators in the noise-free and noise-dominated limits.  This analysis parallels the discussion of temperature power spectrum estimates given in E04. However, in the noise-dominated limit, an optimal PCL polarization estimator can only be constructed for the special case of uncorrelated noise and identical noise levels in the Q and U maps, $(\\sigma^Q)^2_i = (\\sigma^U)^2_i$. Fortunately, for a {\\it Planck}-type experiment, both of these assumptions should be reasonably accurate at high multipoles.  An analytic discussion of the noise properties of CMB temperature maps for a {\\it Planck}-like scanning strategy with realistic `$1/f$' noise is given by Efstathiou (2005). For a {\\it Planck}-type experiment, the noise should be accurately white at high multipoles (see also Stompor and White, 2004). Low frequency `$1/f$' noise introduces striping in the maps with a characterstic spectrum that varies approximately as $\\Delta C_\\ell \\propto 1/\\ell$ at low multipoles. This low frequency behaviour for the temperature power spectra has been verified in many simulations of {\\it Planck}-like experiments ({\\it e.g.} Burigana \\etals 1997, Maino \\etals 1999; Keih\\\"anen \\etals 2004; Efstathiou 2005). It is also seen in the $X$, $E$ and $B$ power spectra from full-scale simulations of the polarised {\\it Planck} $217$GHz detectors (Ashdown \\etals 2006).  For the knee-frequencies\\footnote{The characteristic frequency above which the detector noise is approximately white.} and noise levels expected of {\\it Planck} detectors, the noise power spectra for both temperature and polarization will be accurately white-noise above multipoles of $\\ell \\simgt 20$.  For the temperature power spectrum, the effect of residual striping errors on the power spectrum at multipoles $\\ell \\simlt 20$ should be negligible, and so it should be an excellent approximation to treat the noise as pure white-noise. However, the simulations of Ashdown \\etals (2005) show that the effects of striping errors on the $E$ and $B$-mode power spectra at low multipoles will be comparable to (or will dominate) the intrinsic signal. These errors should therefore be folded into estimates of the power spectrum covariance matrices. This can be done by evaluating the QML estimator on  low resolution maps, provided one can adequately approximate the noise covariance matrices $N_{ij}$ for these low resolution maps. In Section 5 we showed that it is easy to compute $N_{ij}$ for a low resolution map if the noise is strictly uncorrelated. However, the analogous problem for {\\it Planck}-like correlated noise has not yet been solved and is currently under investigation.  Since the noise for a {\\it Planck}-like experiment is accurately white at some characteristic multipole, a hybrid estimator based on PCL estimates at high multipoles and QML estimates at low multipoles, is well motivated. Furthermore, by applying a QML estimator at low multipoles, it is possible to eliminate almost completely the troublesome effects of $E$ and $B$ mode mixing associated with sky cuts. We have previously argued (E04, Efstathiou 2004b) that maximum likelihood estimators should be used to determine the temperature power spectrum at low multipoles on a cut sky to eliminate the large `estimator-induced' variance inherent to PCL estimators. The problem of estimator-induced variance is even more acute for PCL estimators of the polarization power spectra ({\\it cf} CC05). In particular, for the PCL estimators discussed here, mode-mixing places significant limits on the amplitude of the tensor-to-scalar ratio, $r$, that can be probed even in a noise-free experiment. For example, the tests described in Section 3.4 show that it is difficult to probe below $r \\sim 0.1$ using a PCL estimator applied to maps with a Kp0 sky-cut. As Smith (2005) has shown, it is possible to define fast estimators based on unambiguous $E$ and $B$ modes on a cut sky. However, applying a QML estimator to low (or intermediate) resolution maps is also fast and is very close to optimal. In fact, if one wants to incorporate QML power spectrum estimation into a Monte-Carlo chain, then as equation (\\ref{MLN1}) shows, the time taken to evaluate a QML power spectrum is no greater than the time taken to evaluate a fast spherical transform.      \\vskip 0.1 truein   \\noindent {\\bf Acknowledgements:} I thank members of the Cambridge Planck Analysis Centre, especially Mark Ashdown and Anthony Challinor, for helpful discussions."
    },
    "0601/astro-ph0601476_arXiv.txt": {
        "abstract": "{ We have derived abundances of 33 elements and upper limits for 6 additional elements for the metal-poor ($\\mathrm{[Fe/H]} = -2.42$) turn-off star HE~0338$-$3945 from high-quality VLT-UVES spectra. The star is heavily enriched, by about a factor of 100 relative to iron and the Sun, in the heavy $s$-elements (Ba, La, ..).  It is also heavily enriched in Eu, which is generally considered an $r$-element, and in other similar elements. It is less enriched, by about a factor of 10, in the lighter $s$-elements (Sr, Y and Zr). C is also strongly enhanced and, to a somewhat lesser degree, N and O.  These abundance estimates are subject to severe uncertainties due to NLTE and thermal inhomogeneities which are not taken into detailed consideration. However, an interesting result, which is most probably robust in spite of these uncertainties, emerges: the abundances derived for this star are very similar to those of other stars with an overall enhancement of all elements beyond the iron peak.  We have defined criteria for this class of stars, r+s stars, and discuss nine different scenarios to explain their origin. None of these explanations is found to be entirely convincing. The most plausible hypotheses involve a binary system in which the primary component goes through its giant branch and asymptotic giant branch phases and produces CNO and $s$-elements which are dumped onto the observed star. Whether the $r$-element Eu is produced by supernovae before the star was formed (perhaps triggering the formation of a low-mass binary), by a companion as it explodes as a supernova (possibly triggered by mass transfer), or whether it is possibly produced in a high-neutron-density version of the $s$-process is still unclear. Several suggestions are made on how to clarify this situation. ",
        "introduction": "Elements with atomic numbers $\\mathrm{Z} > 30$ are believed to be almost exclusively synthesized in neutron-capture ($n$-capture) reactions. In the most metal-poor stars the overall abundance of these elements varies from star to star, by more than a factor of 100 at a given metallicity (McWilliam et~al.~\\cite{mcwilliam1995}, Ryan et~al.~\\cite{ryan1996}). Also, the different abundance ratios vary, e.g. the Ba/Eu ratio tends to decline with decreasing [Fe/H]\\footnote{We use the standard notations $\\mathrm{A/B} = N_\\mathrm{A}/N_\\mathrm{B} $ and $\\mathrm{[A/B]} = \\log ( N_\\mathrm{A} / N_\\mathrm{B} )_\\star - \\log (N_\\mathrm{A} / N_\\mathrm{B})_\\mathrm{\\sun}$ where $N_\\mathrm{X}$ are number densities.} (McWilliam~\\cite{mcwilliam1998}, Burris et~al.~\\cite{burris2000}). Eu is often thought to be synthesized in conditions of very high neutron fluxes (the {\\it r-process}) while Ba is most easily made in conditions of much lower neutron fluxes (the {\\it s-process}). The decline of Ba/Eu may show the relative dominance of $r$-process sites in the early history of the Galaxy.  The first Extremely Metal Poor (EMP, Beers \\& Christlieb~\\cite{beers2005}) star found to be enriched in $n$-capture elements was the giant HD\\,115444 with $\\mathrm{[Fe/H]} \\sim -3$ (Griffin et~al.~\\cite{griffin1982}). Westin et~al.~(\\cite{westin2000}) found marginal overabundances of Ba and other $s$-elements relative to iron in this star as compared with the Sun. This still makes the star rich in these elements relative to ``normal'' Population\\,{\\sc ii} stars of the same metallicity. However, it has a high abundance of europium, $\\mathrm{[Eu/Fe]} = 0.85$, even compared to the Sun.  The heavy elements of HD\\,115444 show good agreement in relative abundances with scaled solar $r$-process abundances, as does the EMP giant CS\\,22892-052 which is even more $r$-element rich (McWilliam et~al.~\\cite{mcwilliam1995}, Sneden et~al.~\\cite{sneden2003}). This agreement suggests that a similar process is responsible for generating the heavy $r$-elements, both for the most metal-poor stars and for the Sun. However, several studies have found a deviation from the scaled solar $r$-process abundance pattern for the light $r$-process elements (e.g.\\ Sneden et~al.~\\cite{sneden2000}, \\cite{sneden2003}, Cowan et~al.~\\cite{cowan2002}, Aoki et~al.~\\cite{aoki2005}). Also, similar stars like CS\\,31082-001 (Hill et~al.~\\cite{hill2002}) and CS\\,30306-132 (Honda et~al.~\\cite{honda2004a}) seem to have had an ``actinide boost'', which has selectively affected the abundances of the most heavy $r$-elements. These variations in behaviour of the $r$-elements suggest that multiple astrophysical sites for the $r$-process may exist.  The discovery of CS\\,22892-052 was a result of the compilation of metal-poor HK survey objects of Beers et~al.~(\\cite{beers1992}). Similarly, the large programme of Cayrel et~al.~(\\cite{cayrel2004}) at ESO, ``First Stars'', was founded on that survey. In the First Stars project, high-resolution studies of the very metal-poor giant CS\\,31082-001 revealed its high Eu abundance. It was possible to identify a \\ion{U}{ii} line in this star, and this opened up the possibility to use the U/Th ratio as a new chronometer. More recently, Honda et~al.~(\\cite{honda2004a, honda2004b}) reported on two new Eu-rich stars, again taken from the HK survey, namely CS\\,30306-132  with $\\mathrm{[Eu/Fe]} = +0.85$ and CS\\,22183-031 with $\\mathrm{[Eu/Fe]} = +1.2$. According to the classification suggested by Beers \\& Christlieb (\\cite{beers2005}) the latter of these is an r-II star (having $\\mathrm{[Eu/Fe]} > +1.0$ and $\\mathrm{[Ba/Eu]} < 0.0$) while the former is classified as r-I ($+0.3 \\le \\mathrm{[Eu/Fe]} \\le +1.0$ and $\\mathrm{[Ba/Eu]} < 0.0$). Another interesting object which is very overabundant in $n$-capture elements, the carbon-enhanced metal-poor (CEMP) star with $s$-element enhancement CS\\,31062-050, was recently studied by Johnson \\& Bolte~(\\cite{johnson2004}). Interestingly, they concluded that the abundance profile of the star, in spite of its high Eu abundance, suggests that the $r$-process was not responsible in this case. These indications will be further discussed in Sect.\\,\\ref{discussion}.  Although the contribution of the $s$-process to the gas in the very early Galaxy may in general be tiny, as one would expect due to the small fraction of iron to start the nucleosynthesis from, some stars with [Fe/H] as low as about $-3$ have been found to have clear $s$-element signatures (Johnson \\& Bolte~\\cite{johnson2002a}, Simmerer et~al.~\\cite{simmerer2004}, Sivarani et~al.~\\cite{sivarani2004}). The implications of the presence of these stars are still not fully understood. It was predicted by Gallino et~al.~(\\cite{gallino1998}) (see also Goriely \\& Mowlavi~\\cite{goriely2000}), that the fraction of the heaviest $s$-elements should be enhanced in metal-poor environments with respect to the lighter $s$-elements, a result of the scarcity of seed nuclei with respect to neutrons. This prediction was qualitatively confirmed by Aoki et~al.~(\\cite{aoki2000}) who made the first discovery of Pb in a $s$-element rich very metal-poor star, LP625-44. Somewhat more iron-rich ``lead stars'' were also discovered by Van Eck et~al.~(\\cite{van_eck2001}) (see also Sivarani et~al.~\\cite{sivarani2004}). We also note that the abundances of the ``lighter'' $n$-capture elements (Sr, Y, Zr) do not scale very well with those of the heavier elements (Ba etc, Z $\\ge 56$). This may be due to the fact that the lighter $s$-elements (the so-called ``weak component'', see Prantzos et~al.~\\cite{prantzos1990}) received important contributions from seed nuclei reacting with neutrons created in the $^{22}$Ne($\\alpha$,n)$^{25}$Mg reaction in massive AGB stars, while the heavy elements were essentially contributed by neutrons from the $^{13}$C($\\alpha$,n)$^{16}$O reaction in less massive stars.  Another class of stars was discovered by Barbuy et~al.~(\\cite{barbuy1997}) and Hill et~al.~(\\cite{hill2000}). Two carbon-rich Population\\,{\\sc ii} giants, the CEMP stars CS\\,22948-027 and CS\\,29497-034 from the HK survey, were found not only to be rich in the commonly enhanced $s$-elements such as Sr, Y, Ba and La, but also the $r$-element Eu was significantly enhanced. Preston \\& Sneden~(\\cite{preston2001}) found another star of this class in the HK survey (CS\\,22898-027), Aoki et al. (\\cite{aoki2002}) added two more (CS\\,29526-110 and CS\\,31062-012) and Sivarani et~al.~(\\cite{sivarani2004}) another (CS\\,29497-030) (see also Ivans et~al.~\\cite{ivans2005}). The stars of this type presently known are listed in Table\\,\\ref{Tstars}. We shall denote these stars {\\it r+s stars}, and we postpone a detailed discussion of the definition of this class of stars to Sect.\\,\\ref{discussion}.  The origin of the abundance peculiarities of the r+s stars is not clear, and many scenarios have been proposed (see Sect.\\,\\ref{discussion}). However, this is not the only reason for studying them further. In fact, the astrophysical sites of the $r$-process are still unclear. The precise sites of the $s$-process among the most metal-poor stars are also still debatable, as well as the fraction of nuclei provided by neutrons from the $^{13}$C($\\alpha$,n)$^{16}$O reaction as compared with the $^{22}$Ne($\\alpha$,n)$^{25}$Mg neutron source. One might hope that a clarification of the origin of the r+s stars may shed some light on the general questions concerning the sites of the $r$- and $s$-processes.  The Hamburg/ESO objective-prism survey for bright quasars (HES; Wisotzki et~al.~\\cite{wisotzki2000}) has been used for a search for metal-poor stars (Christlieb~\\cite{christlieb2003}). One of these was the r+s star HE~2148$-$1247, found in the Keck pilot programme on EMP stars from the HES (Cohen et~al.~\\cite{cohen2003}). In a systematic approach to exploring the $n$-enhanced HES stars, the Hamburg/ESO R-process Enhanced Star survey (HERES) has been carried out (Christlieb et~al.~\\cite{christlieb2004}, hereafter Paper\\,{\\sc i}). For several hundred candidate stars from HES, ``snapshot'' spectra at moderate resolving power ($R \\sim 20\\,000$) and low $S/N$ ($\\sim 50$) were obtained with the main goal to find $r$-element enriched stars, as disclosed by strong Eu lines. These spectra were analysed with an automatic procedure based on MARCS model atmospheres and synthetic spectra (Barklem et al.~\\cite{barklem2005}, Paper\\,{\\sc ii}). As a result of the HERES survey, a number of new detections of $n$-capture enriched EMP stars were made. In particular, 8 new r-II stars and 35 r-I stars were found.  One of the stars found by the HERES survey is HE~0338$-$3945. This object is a star located close to the main-sequence turnoff of an old halo population with an overall metallicity of $\\mathrm{[Fe/H]} \\sim -2.4$, and enriched in both Ba and Eu. We decided to obtain spectra of higher quality of this star and perform a detailed analysis which is presented here. In Sects.\\,\\ref{observations} and \\ref{analysis} the observations and analysis respectively are described.  In Sect.\\,\\ref{resultsA} the abundance results are presented.  In Sect.\\,\\ref{comparisons} the abundances are compared with those of other stars.  This reveals HE~0338$-$3945 to be very similar to another r+s star, HE~2148$-$1247 (Cohen~et~al.~\\cite{cohen2003}).  This result led us to survey the known r+s stars in the literature, and this survey, along with a comparison with HE~0338$-$3945 is also presented.  In Sect.\\,\\ref{discussion} the results for HE~0338$-$3945 are discussed, as is both the classification and formation of r+s stars in general.  Finally, in Sect.\\,\\ref{conclusions} the conclusions are presented. ",
        "conclusions": "\\label{conclusions}   In this paper we have analysed a high-quality spectrum of the $r$- and $s$-element rich star HE~0338$-$3945, deriving abundances for 33 elements and upper limits for an additional 6 elements.  We found the abundances of this star to be very similar to those of HE~2148$-$1247, which is supported by an homogeneous analysis of the two spectra.  In fact the abundance patterns among known r+s stars in the literature seem to be quite similar considering the differences between analyses.  Based on our results for HE~0338$-$3945 and our survey of other r+s stars, we have discussed a range of scenarios for the formation of the r+s stars.   Some scenarios have considerable merit, but all have also drawbacks or problems. Scenario I is already dismissed. Scenarios II and III need the ISM to be highly inhomogeneous to produce such high $r$-process abundances as detected in these stars.  We note that the [Eu/Fe] ratios for the r+s stars are systematically higher than those of the r-II stars which might suggest that a process with two independent steps, one of which is also responsible for the r-II stars is not very probable.  However, this is uncertain due to the possibly significant contribution of Eu by the $s$-process, and the possible relative effects of convective mixing between the r+s stars, which are predominantly TOP stars, and the r-II stars, which are all giants (see also discussion below). Scenario II is also improbable as it requires all r+s stars to have passed at least the RGB phase.  Scenario III presently seems unlikely in view of the frequency of r+s stars and their apparent chemical homogeneity. The discussed modification of scenario III to include supernova-triggered binary formation seems plausible, though it is presently unclear whether such events can provide the observed frequency of r+s stars. Scenario IV is improbable as it requires strict stellar orbit constraints to produce the chemical homogeneity of the r+s class.   Scenarios V and VI may be considered, but as they depend heavily on physical processes and parameters that are poorly understood at present, these scenarios at best have to be regarded as uncertain. The observed chemical uniformity is not an obvious consequence of them.  Scenario VII still needs a site which can produce the high neutron densities at suitable exposures; the scenario also has to be supported by renewed and more realistic $s$-process calculations. It is also not clear how well this scenario would explain the chemical homogeneity. A probable merit of it is, however, that the observed great but varying amounts of Pb may result naturally from the high neutron fluxes. The absence of observed radial velocity variations for several stars (see Preston \\& Sneden~\\cite{preston2001} and Aoki et~al.~\\cite{aoki2002}) is a problem, although not very severe as yet, for all scenarios III-VIII. Scenario VIII seems hypothetical and it is also questionable whether it is efficient enough. The available astrometric data, though admittedly meager, does not support Scenario IX.   The majority of the r+s stars in Table~\\ref{Tstars} are TOP stars, while all the r-II stars are giants. One should note the great difference between the TOP stars and the red giants as regards two important aspects:  First, the convective zone of the TOP stars is less than 5\\% of the stellar mass, while it is typically 75\\% of the mass for the giants. Thus, the dilution of enriched material transferred to the stellar surface of a TOP star is at least a factor of 10 less than for a giant. So, if the enhancements of $n$-capture elements are due to mass transfer from a companion, the amount transferred has to be at least one order of magnitude greater for the r-II stars than for most of the TOP stars. Second, the space volume surveyed in apparent-magnitude limited surveys is at least a factor of 200 greater for the giants than for the r+s stars. This suggests that the space density of r+s stars is considerably higher than that of r-II stars. In spite of their scarcity, it is interesting that no r-II stars have been found at the TOP, since a much smaller transfer of $r$-elements would raise the atmospheric abundances to considerable values.  More statistically controlled surveys are desirable.  Future detailed studies, with accurate analysis of all known and suspected stars, are worthwhile. Spectra with high $S/N$ and high resolution, also in the the UV, are important to obtain better limits on Ag, Th and U and to get results for $r$-elements that have lines in the UV such as Os and Ir. Attempts to measure Eu isotopic ratios for these stars should be continued. The analysis of certain important elements, such as Eu and Ag, should be scrutinised for NLTE and 3D effects.  More detailed calculations of the $s$-process at high neutron densities should be carried out. Finally, all known r+s stars should be monitored in an attempt to disclose their possible binarity."
    },
    "0601/astro-ph0601195_arXiv.txt": {
        "abstract": " ",
        "introduction": "\\label{sec1}  Taken literally, galaxy clusters must be comprised of an overdensity of galaxies. Almost as soon as the debate was settled on whether or not the ``nebulae'' were extragalactic systems, it became clear that their distribution was not random, with regions of very high over- and under-density. Thus, from a historical perspective, it is important to discuss the detection of galaxy clusters through their galactic components. Today we recognize that galaxies constitute a very small fraction of the total mass of a cluster, but they are nevertheless some of the clearest signposts for detection of these massive systems. Furthermore, the extensive evidence for differential evolution between galaxies in clusters and the field (discussed at length elsewhere in these proceedings) means that it is imperative to quantify the galactic content of clusters.  Perhaps even more importantly, optical detection of galaxy clusters is now inexpensive both financially and observationally. Large arrays of CCD detectors on moderate telescopes can be utilized to perform all-sky surveys with which we can detect clusters to $z\\sim0.5$. Using some of the efficient techniques discussed later in this section, we can now survey hundreds of square degrees for rich clusters at redshifts of order unity with 4-meter class telescopes, and similar surveys, over smaller areas but with larger telescopes are finding group-mass systems to similar distances.  Looking to the future, ever larger and deeper surveys will permit the characterization of the cluster population to lower masses and higher redshifts. Projects such as the Large Synoptic Survey Telescope (LSST) will map thousands of square degrees to very faint limits (29th magnitude per square arcsecond) in at least five filters, allowing the detection of clusters through their weak lensing signal (\\ie mass) as well as the visible galaxies. Ever more efficient cluster-finding algorithms are also being developed, in an effort to produce catalogs with low contamination by line-of-sight projections, high completeness, and well-understood selection functions.  This chapter provides an overview of past and present techniques for optical detection of galaxy clusters. It follows the progression of cluster detection techniques through time, allowing readers to understand the development of the field while explaining the variety of data and methodologies applied. Within each section we describe the datasets and algorithms used, pointing out their strengths and important limitations, especially with respect to the characterizability of the resulting catalogs. The next section provides a historical overview of pre-digital, photographic surveys that formed the basis for most cluster studies until the start of the twenty-first century. Section three describes the hybrid photo-digital surveys that created the largest current cluster catalogs. The fourth section is devoted to fully digital surveys, most specifically the Sloan Digital Sky Survey and the variety of methods used for cluster detection. We also describe smaller surveys, mostly for higher redshift systems. The fifth section gives an overview of the different algorithms used by these surveys, with an eye towards future improvements. The concluding section discusses various tests that remain to be done to fully understand any of the catalogs produced by these surveys, so that they can be compared to simulations. ",
        "conclusions": "It is clear that there exist many methods for detecting clusters in optical imaging surveys. Some of these are designed to work on very simple, single--band data (AK, Matched Filter, VT), but will work on multicolor data as well. Others, such as maxBCG and the CRS method, rely on galaxy colors and the red sequence to potentially improve cluster detection and reduce contamination by projections and spurious objects. Very little work has been done to compare these techniques, with the exceptions of Kim \\etal 2001, Bahcall \\etal (2003) and Lopes \\etal (2004), each of whom compared the results of only two or three algorithms. Even from these tests it is clear that no single technique is perfect, although some (notably those that use colors) are clearly more robust. Certainly any program to find clusters in imaging data must consider the input photometry when deciding which, if any, of these methods to use.  One of the most vexing issues facing cluster surveys is our inability to compare directly to large scale cosmological simulations. Most such simulations are N-body only, but have perfect knowledge of object masses and positions. Thus, it is possible to construct algorithms to detect overdensities based purely on mass, but it is {\\em not} possible to obtain the photometric properties of these objects! Recent work, such as the Millennium Simulation (Springel \\etal 2005), is approaching this goal. It is necessary to extract from these simulations the magnitudes of galaxies in filters used for actual surveys, and run the various cluster detection algorithms on these simulated galaxy catalogs. The results can then be compared to that of pure mass selection, and the redshift-, structure- and mass-dependent biases understood. Ideally, this should be done for many large simulations using different cosmologies, since the galaxy evolution and selection effects will vary. Such work is fundamental if we are to use the evolution of the mass function of galaxy clusters for cosmology. As deeper and larger optical surveys, such as LSST, and other techniques such as X-ray and Sunyaev--Z'eldovich effect observations become available, the need for these simulations becomes ever greater."
    },
    "0601/astro-ph0601640_arXiv.txt": {
        "abstract": "An approximate Riemann solver for the equations of relativistic magnetohydrodynamics (RMHD) is derived. The HLLC solver, originally developed by  Toro, Spruce and Spears, generalizes the algorithm described in a previous paper (Mignone \\& Bodo 2004) to the case where magnetic fields are present. The solution to the Riemann problem is approximated by two constant states bounded by two fast shocks and separated by a tangential wave. The scheme is Jacobian-free, in the sense that it avoids the expensive characteristic decomposition of the RMHD equations and it improves over the HLL scheme by restoring the missing contact wave.  Multidimensional integration proceeds via the single step, corner transport upwind (CTU) method of Colella, combined with the contrained tranport (CT) algorithm to preserve divergence-free magnetic fields. The resulting numerical scheme is simple to implement, efficient and suitable for a general equation of state. The robustness of the new algorithm is validated against one and two dimensional numerical test problems. ",
        "introduction": "Strong evidence nowadays supports the general idea that relativistic plasmas may be closely related with most of the violent phenomena observed in astrophysics. Most of these scenarios are commonly believed to involve strongly magnetized plasmas around compact objects. Accretion onto super-massive black holes, for example, is invoked as the primary mechanism to power highly energetic phenomena observed in active galactic nuclei, \\citep{Macchetto99,ERZ02,McK05,Shapiro05}. In this respect, the formation and propagation of relativistic jets and the accretion flow dynamics pose some of the most challenging and interesting quests in modern theoretical astrophysics. Likewise, a great deal of attention has been addressed, in the last years, to the darkling problem of gamma ray bursts \\citep[see for example][]{MR94,McFW99,KG02,RRD03}, whose models often appeal to strongly relativistic collimated outflows \\citep{Aloy_etal00, Aloy_etal02}. Other attractive examples include pulsar wind nebulae \\citep{BZAV05}, microquasars \\citep{Meier03,McKG04}, X-ray binaries \\citep{VRT02} and stellar core collapse in the context of general relativity \\citep{Bruenn85, DFM02}.  Theoretical investigations based on direct numerical simulations have paved a way towards a better understanding of the rich phenomenology of relativistic magnetized plasmas. Part of this accomplishment owes to the successful generalization of existing shock-capturing Godunov-type codes to relativistic magnetohydrodynamics (RMHD) \\citep[see][and reference therein]{K99, Balsara01,dZBL03}. Implementation of such codes is based on a conservative formulation which requires an exact or approximate solution to the Riemann problem, i.e., the decay of a discontinuity separating two constant states \\citep{Toro97}. In terms of computational cost, employment of exact relativistic Riemann solvers may become prohibitive due to the high degree of intrinsic nonlinearity present in the equations. This has focused most computational efforts towards the development of approximate solvers which, nevertheless, require knowledge of the exact solution, at least on some level \\citep{MM03}. The presence of magnetic fields further entangles the solution, since the number of decaying waves increases from three to seven \\citep{AP87,Anile89}. An exact analytical approach to the solution (which does not allow compound waves) has been recently presented in \\cite{GR05}, while \\cite{RMPIM05} derived a special case where the velocity and magnetic field are orthogonal.  The trade-off between efficiency, accuracy and robustness of such approximate methods is still a matter of research. Solvers based on local linearization have been presented in \\cite{K99} (KO henceforth), \\cite{Balsara01} (BA henceforth) and \\cite{KKU02}. Despite the higher accuracy in reproducing the full wave structure, these solvers rely on rather expensive characteristic decompositions of the Jacobian matrix. Conversely, the characteristic-free formulation of Harten-Lax-van Leer (HLL) of \\cite{HLL83} has gained increasing popularity due to its ease of implementation and robustness. The HLL approach has been successfully applied to the RMHD equations by \\citealt{dZBL03} (dZBL henceforth) as well as to the general relativistic case \\citep[see for example][]{GmKT03,DLSS05} and to the investigation of extragalactic jets, see \\cite{LAAM05}.  Besides the computational efficiency, however, the HLL formulation averages the full solution to the Riemann problem into a single state, and thus lacks the ability to resolve single intermediate waves such as Alfv{\\'e}n, contact and slow discontinuities. In \\cite{MB05} (paper I henceforth) we proposed an approach that cured this deficiency by restoring the missing contact wave. The resulting scheme generalized the HLLC approximate Riemann solver by \\citet{TSS94} to the equations of relativistic hydrodynamics without magnetic fields. Here, along the same lines, we propose an extension of the HLLC solver to the relativistic magnetized case. Similar work has been presented in the context of classical MHD by \\cite{Gurski04} and \\cite{Li05}.  The new HLLC Riemann solver is implemented in the framework of the corner transport upwind (CTU) method of \\cite{Colella90}, coupled with the constrained transport (CT) evolution \\citep{EH88} of magnetic field. The algorithm naturally preserves the divergence-free condition to machine accuracy and is stable up to Courant number of $1$.  The paper is organized as follows. The relevant equations are given in \\S\\ref{sec:equations}. In \\S\\ref{sec:hllc} we derive the new HLLC Riemann solver. Numerical tests, together with the implementation of the CTU-CT method are shown in \\S\\ref{sec:test}. ",
        "conclusions": "An HLLC approximate Riemann solver has been developed for the relativistic magnetohydrodynamic equations. The new approach improves over the single state HLL solver in the ability to capture exactly isolated tangential and contact discontinuities. Several test problems in one and two dimensions demonstrate better resolution properties and a reduced amount of the numerical diffusion inherent to the averaging process of the single state HLL scheme. The solver is well-behaved for strictly two-dimensional flows, although applications to genuinely three-dimensional problems may suffer from a pathological singularity when the component of magnetic field normal to a zone interface approaches zero. This feature does not persist in the classical limit.  Multidimensional integration has been formulated in a versatile and efficient way within the framework of the corner transport upwind (CTU) method. The algorithm is stable up to Courant numbers of $1$ and preserves the divergence-free condition via constrained transport evolution of the magnetic field. The additional computational cost and the numerical implementation in an existing relativistic MHD code are minimal."
    },
    "0601/astro-ph0601689_arXiv.txt": {
        "abstract": "Spectra of the spreading layers on the neutron star surface are calculated on the basis of the Inogamov-Sunyaev model taking into account general relativity correction to the surface gravity and considering various chemical composition  of the accreting matter. Local (at a given latitude) spectra are  similar to the X-ray burst spectra and are described by a diluted black body. Total spreading layer spectra are integrated accounting for the light bending, gravitational redshift, and the relativistic Doppler effect and aberration. They depend slightly  on the inclination angle  and on the luminosity. These spectra also can be fitted by a diluted black body with the color temperature depending mainly on a neutron star compactness. Owing to the fact that the flux from the spreading layer is close to the critical Eddington, we can put constraints on a neutron star radius without the need to know precisely the emitting region area or the distance to the source. The  boundary layer spectra observed in the luminous low-mass X-ray binaries, and described by a black body of color temperature $\\Tc=2.4 \\pm 0.1$ keV, restrict  the neutron star radii   to $R=14.8 \\pm 1.5$ km (for a 1.4-$\\msun$ star and solar composition of the accreting matter), which corresponds to the hard equation of state. ",
        "introduction": "Matter accreting on to a weakly magnetized neutron star (NS) in low mass X-ray binaries (LMXRBs) can form an accretion disc which extend down to the NS surface.  A boundary layer (BL) is formed between the accretion disc and the NS surface, where a rapidly rotating matter of the disc is decelerated down to the NS angular velocity.  The amount of the energy, which is generated during this process, is comparable with the energy generated in the accretion disc \\citep{SS86,SS98}.  There is no generally accepted theory of the BL. Two different approaches to the BL description are considered. First of them, which we will call a `classical model', considers the BL between a central star (a white dwarf or a NS) as a part of the accretion disc \\citep{P77,PS79,T81,SS88,BK94,PN95,PS01}. In this model the component of velocity normal to the accretion disc plane is zero. The half-thickness of the BL is determined by the same relation, as for the accretion disc: \\be \\label{eq:Hbl} \\Hbl\\sim \\frac{c_{\\rm s}}{v_{\\rm K}} R, \\ee where $c_{\\rm s}$ is a sound speed in the BL and $v_{\\rm K}$ is the Keplerian velocity close to the NS surface of radius $R$. The radial extension of the BL is determined by the relation \\citep{P77} \\be \\label{eq:hbl} \\hbl\\sim \\frac{c^2_{\\rm s}}{v^2_{\\rm K}} R \\sim \\Hbl\\frac{\\Hbl}{R}. \\ee In this classical model the accreting matter in the BL is decelerated in the accretion disc plane, along radial coordinate only, due to the viscosity operating within the differentially rotating BL, similarly to the accretion disc. From the observational point of view, the classical BL is a bright equatorial belt close to the NS surface.  The effective temperature of the BL is higher than the maximum accretion disc effective temperature, because the BL is smaller than the accretion disc, while their luminosities  are comparable.  Another approach was suggested by  \\citet[][ hereafter IS99]{IS99}. The BL is considered as a spreading layer (SL) on the NS surface. The accreting matter diffusing along the radial direction in the accretion disc  and reaching the neutron star surface gains the velocity component normal to the accretion disc plane due to the ram pressure from  the accretion disc. Then the matter spirals along the NS surface  towards the poles.  Rotating at the NS surface, the matter is decelerated due to a turbulent friction between the rapidly rotating matter and a slowly rotating NS surface. The kinetic energy of the accreting gas is mostly deposited in two bright belts at some latitude above and below the NS equator. The width and the latitude of the belts depend on the mass accretion rate. The larger the accretion rate, the wider the belts are  and the closer they are to the NS poles.  At the accretion rate close to the Eddington limit ($L_{\\rm BL} \\sim \\Ledd$) the bright belts  expand all over the NS surface.  The observational difference between two BL models is not significant. At low accretion rates ($\\Lbl\\sim 0.01 \\Ledd$) the latitude of bright belts of the SL is small ($\\sim$ few degrees) and the vertical extension of the SL is comparable to the classical BL thickness. At high accretion rates ($\\Lbl\\sim  \\Ledd$) the classical BL thickness is comparable to the NS radius   \\citep[see][]{PS01}. Therefore, the effective temperatures of these BL models are of the same order.  In the approach by IS99, the NS radius is assumed to be larger that $3\\Rs$ (where $\\Rs=2GM/c^2$ is the Schwarzschild radius of a NS of mass $M$), but the accretion disc structure is not significantly changed up to the NS surface, and the radial velocity is always subsonic. If the radial velocity of accreting gas is supersonic at the surface (see e.g.  \\citealt*{PN92}, for a possibility of the supersonic radial velocity in classical BL, and \\citealt*{KW91}, for the ``gap accretion'' when the NS is within the innermost stable circular orbit), some  fraction of the kinetic energy (associated with a small radial velocity component) should be dissipated in an oblique shock, but most of it  still remains stored in the kinetic energy of the gas rotating  at the surface to be dissipated later  in the SL. The gap accretion  model of \\citet{KW91} is rather similar to the SL, but they did not consider  the fate of the accreting material and its spread over the surface. At  very low accretion rate, both models should produce hard Comptonization spectra extending to $\\sim100$ keV.  The aim of this work is the calculation of the radiation spectra of the SLs  and their comparison to the observed X-ray spectra of the BLs in the luminous LMXRBs. ",
        "conclusions": "We have derived the one-dimensional   equations   describing the SL model on a spherical NS surface from the usual hydrodynamic equations. The obtained equations are similar to those in IS99, except for the  energy conservation law  where we neglected the surface density of the  gravitational potential energy which is of the second order in $H/R$. This difference, however,  leads only to small quantative changes. We have also implemented a pseudo-Newtonian potential to account for the main general relativity corrections  and considered various chemical compositions of the accreting matter.  We have studied the vertical (radial) structure of the SL  with different assumptions about the vertical distributions of the radiation flux and azimuthal velocity. The temperature structure and the emergent radiation spectra of the SL are computed accounting for the effect of Compton scattering. We showed that the local (at a given latitude) emergent spectra depend very little on details of the SL vertical structure in optically thick cases with $\\Sigmas\\gtrsim 100\\  \\mbox{g\\ cm}^{-2}$ ($L \\gtrsim 0.1 \\Ledd$). These spectra can be described by the diluted Planck spectrum and are similar to the spectra  of X-ray bursts  with the same effective temperature and the effective surface gravity.  The integral SL spectra were computed accounting for relativistic effects such as the gravitational redshift and light bending, the relativistic Doppler effect and aberration. These spectra slightly depend on the inclination angle to the line of sight and on the SL luminosity. The local effective temperature increases with latitude, while the hardness factor $\\fc$  decreases. This leads to only slight variation of the color temperature on latitude. As a result, the integral spectra can also be well described by a single-temperature diluted Planck spectrum.  We compared our theoretical integral SL spectra with the observed  spectra of the LMXRBs BLs. The observed color temperature of 2.4 $\\pm$ 0.1 keV \\citep{GRM03,RG06} can be reproduced for hard equations of state of NS material. Our model constrains radii of NSs in LMXRBs to 13--16 km for a 1.4 solar mass star. Soft equations of state (smaller NS radii) can be reconciled with the observed spectra only for very low viscosity $\\alphab\\sim10^{-5}$. Calculation of $\\alphab$ from the first principles is a challenging problem that deserves further attention."
    },
    "0601/astro-ph0601530_arXiv.txt": {
        "abstract": "We present the first phase-coherent measurement of a braking index for the young, energetic rotation-powered pulsar \\psr. This 324\\,ms pulsar is located at the center of the supernova remnant Kes 75 and has a characteristic age of $\\tau_c = 723$\\,years, a spin-down energy of $\\dot{E}=8.3 \\times 10^{36}$\\,erg\\,s$^{-1}$, and inferred magnetic field of $4.9\\times 10^{13}$\\,G. Two independent phase-coherent timing solutions are derived which together span 5.5\\,yr of data obtained with the {\\textit{Rossi X-ray Timing Explorer}}. In addition, a partially phase-coherent timing analysis confirms the fully phase-coherent result. The measured value of the braking index, $n=2.65\\pm0.01$, is significantly less than 3, the value expected from magnetic dipole radiation, implying another physical process must contribute to the pulsar's rotational evolution. Assuming the braking index has been constant since birth, we place an upper limit on the spin-down age of \\psr\\ of 884\\,yr, the smallest age estimate of any rotation-powered pulsar. ",
        "introduction": "\\label{sec:intro} The measurement of pulsar braking indices ($n$) is crucial to the understanding of the physics underlying pulsar spin down, assumed to be of the form \\begin{equation} \\dot{\\nu} = -K \\nu^n, \\end{equation} where $\\nu$ is the pulse frequency, $\\dot{\\nu}$ is the frequency derivative and $K$, assumed to be constant, is related to the pulsar's magnetic field and moment of inertia. By taking a time derivative of the above spin-down equation, we find that $n = {\\nu \\ddot{\\nu}}/{\\dot{\\nu}^2}$, where $\\ddot{\\nu}$ is the second frequency derivative. It is typically assumed that pulsars radiate as perfect magnetic dipoles, implying $n=3$. This assumption is used implicitly for the estimation of pulsar magnetic fields as well as to calculate characteristic ages (defined as $\\tau_c=P/2{\\dot{P}}$ corresponding to $n=3$). However, of the five unambiguous measurements of pulsar braking indices obtained so far, all yield a value of $n<3$ \\citep{lps93,lpgc96,ckl+00,lkg05,lkgm05}. Explanations for this discrepancy include the possibility that: the relativistic pulsar wind affects the spin-down \\citep{mt69}; the pulsar may suffer a propeller torque from a putative supernova fallback disk \\citep{mph01}; the pulsar has a time-varying magnetic moment \\citep{br88}; or, the pulsar cannot be modelled as a point dipole, rather, a dipole of some finite size which leads to $n<3$ \\citep{mel97}.  There are few pulsars that are potential candidates for a significant measurement of $n$; 4 of the 5 measured values are from sources with characteristic ages under 2000\\,yr. The first requirement for a significant measurement of $n$ is that the pulsar spins down sufficiently quickly for a measurement of $\\ddot{\\nu}$. The second requirement is that the position of the pulsar is accurately known at the $\\sim$1 arcsecond level. The third requirement is that the spin-down must not be seriously affected by glitches,  sudden spin-ups of the pulsar, or timing noise, a long-term, low-frequency stochastic wandering of the rotation about the overall trend. Typically, glitches begin to seriously affect smooth spin down at characteristic ages of $\\sim 5-10$\\,kyr \\citep{ml90,mgm+04}. Thus many of the pulsars that may spin down fast enough for a measurement of $n$ are irretrievably contaminated by glitches \\citep[e.g.][]{mgm+04}. Timing noise varies from object to object, though is roughly correlated with spin-down rate \\citep[e.g.][]{antt94} and can prevent a measurement of $n$ in a finite data set in an unpredictable way.    The very young, energetic pulsar \\psr\\ was discovered at the center of the supernova remnant Kes 75 \\citep{gvb+00}. Neutral hydrogen absorption measurements indicate that the pulsar resides well across the Galaxy, roughly at a distance of 19\\,kpc \\citep{bh84}. Radiation from this pulsar is detected exclusively in the X-ray band, with a 0.5-10\\,keV luminosity of $L_x=4.1 \\times 10^{35}$\\,erg\\,s$^{-1}$ for a distance of 19 kpc \\citep{hcg03}. If correct, the spin-down energy to X-ray conversion efficiency is 1.6\\%, 6 times greater than that found in the Crab pulsar. The X-ray spectrum is typical of rotation-powered pulsars, with non-thermal power-law emission with $\\Gamma =1.4$ \\citep{hcg03}. However, its large inferred magnetic field of $B = 4.9 \\times 10^{13}$\\,G and spectrum place it in an emerging class of rotation-powered pulsars with magnetar-strength fields \\citep[e.g.][]{km05}. With a characteristic age of $\\tau_c = 723$\\,yr, \\psr\\ is likely the youngest of all known rotation-powered pulsars.  In this paper, we present the first phase-coherent timing solution for the Kes 75 pulsar based on 5.5\\,yr of X-ray timing data from a long-term monitoring campaign with the {\\textit{Rossi X-ray Timing Explorer}} (\\rxte). ",
        "conclusions": "\\label{sec:discussion} Establishing true ages of pulsars is important for several aspects of neutron-star studies, including neutron star cooling, population synthesis, spatial velocity estimates, and to consider their associations with supernova remnants. Based on spin-down properties alone, the true pulsar age is impossible to determine. However, we can estimate the age based on the standard spin-down model ($\\dot{\\nu} \\propto - \\nu^n$) given our measurement of $n$. Integrating this model yields, \\begin{equation} \\tau= -\\frac{1}{n-1} \\frac{\\nu}{\\dot{\\nu}} \\left(1 - \\left(  \\frac{\\nu}{\\nu_0}\\right)^{n-1}   \\right). \\end{equation} Assuming that the initial spin frequency was much larger than the current value, $\\nu_0 \\gg \\nu$, the age estimate for \\psr\\ becomes, \\begin{equation} \\tau \\le - \\frac{1}{n-1} \\frac{\\nu}{\\dot{\\nu}} \\la 884 \\, \\rm{yr}. \\label{eqn:age} \\end{equation} Since the initial spin frequency is unknown, the above estimate provides an upper limit on the pulsar age. Therefore, assuming that $n$ and $K$ have remained constant over the lifetime of the pulsar, the upper limit on the age of \\psr\\ is $\\sim$884\\,yr. This value is larger than the characteristic age of 723\\,yr, but still smaller than the unambiguously known age of the next youngest pulsar, that in the Crab Nebula.  \\citet{mbb+02} did an analysis of frequency measurements of the \\asca\\ and \\beppo\\ data, as well as a small subset of the \\rxte\\ data reported here. Although data gaps prevented them from a firm conclusion regarding $n$ because of the possibility of glitches, their tentative estimate (made assuming no glitches nor substantial timing noise) was $n=1.86 \\pm 0.08$, inconsistent with our result. There are two plausible explanations for this discrepancy. The first is that significant timing noise and/or glitches cannot be definitely excluded in the 6-yr gap between the successive \\asca\\ observations. The second is that the coordinate used to barycenter these data was ultimately found to be off by $\\sim 7''$. The earlier study used coordinates derived from \\asca\\ data with a $20''$ position uncertainty, prior to the availability of the new measurement provided by the \\textit{Chandra X-ray Observatory} used herein.  Our measured $n=2.65\\pm0.01$ is significantly less than 3, as is the case for all established values of $n$, which are shown in Table~\\ref{table:indices}. There are several ideas for the nature of the deviation from the prediction of simple magnetic dipole braking. A time-varying magnetic moment could produce an observed $n$ less than the true $n_0=3$ \\citep{bla94}. A varying \\textit{B}-field can in principle be verified with a measurement of the third frequency derivative, $\\nudotdotdot$, which may or may not ultimately be possible for this source, depending on the number and strength of the glitches it experiences, as well as on the strength of its timing noise. Presently, $\\nudotdotdot$ has been measured for only two pulsars, both of which are consistent with a constant $B$ \\citep{lps93,lkgm05}. Another suggested explanation for $n<3$ is that a fallback disk forms from supernova material and modulates the spin-down of young pulsars via a propeller torque. Spin-down via a combination of magnetic dipole radiation and propeller torque results in $2<n<3$ \\citep[e.g.][]{aay01}. However, this requires that the disk material does not suppress the pulsed emission during the propeller phase, which is difficult to achieve \\citep{mph01}. Angular momentum loss due to a stellar wind would result in $n=1$ \\citep{mt69}. Thus spin-down due to some combination of relativistic pulsar winds associated with young neutron stars and observed indirectly as pulsar wind nebulae \\citep[e.g.][]{rtk+03}, and magnetic dipole spin down may result in measured values $n<3$.  \\citet{mel97} presented an intriguing solution to the $n<3$ problem. He postulated that the radius pertinent to dipole radiation is not the physical neutron-star radius, but a ``vacuum radius'', associated with the location closest to the neutron star where field-aligned flow breaks down. Since this radius can be significantly larger than the neutron-star radius, the system can no longer be treated as a point dipole and $n=3$ is not necessarily true. His model provides a prediction for $n$, given three observables: $\\nu$, $\\dot{\\nu}$ and $\\alpha$, the angle between the spin and magnetic axis. This model predicts that $2<n<3$ for all pulsars and that $n$ approaches 3 as a pulsar ages. Currently, there is no measurement of $\\alpha$ for \\psr. However, for the model to be consistent with our measured $n$ within $3\\sigma$, $\\alpha$ must lie between $8.1 \\degrees$ and $9.6 \\degrees$. This is small but not unreasonable given the very broad pulse profile for this source. Future radio polarimetric observations could in principle constrain $\\alpha$, though at present, no radio detection of this source has been reported \\citep{kmj+96}.  The measurement of $n$ for \\psr\\ brings the total number of measured braking indices to six, shown in Table~\\ref{table:indices}, along with $\\nu$, $\\dot{\\nu}$, $\\tau$, $\\tau_c$, $B_{\\rm{dipole}}$ and $\\dot{E}$ for each object. Comparing this new value to the other four measurements obtained via typical timing methods (i.e. excluding the Vela pulsar whose $n$ was measured using a different method, due to the large glitches experienced by this pulsar), we find no correlations between $n$ and any of the other parameters. If the Vela pulsar is included in the analysis, there is a slight correlation between $n$ and characteristic age. However, there is a relatively large scatter among the five younger pulsars, suggesting the Vela pulsar's value could also be just an extremum of the scatter.   The large scatter in the observed values of $n$ could be intrinsic to the pulsars, for instance, there may exist a relationship between $n$ and $\\alpha$ or $B$, as suggested by \\citet{mel97}. One way to test this possibility would be to obtain more precise measurements of $\\alpha$ for the youngest pulsars, e.g. via polarimetric observations. This, however, may prove difficult given that young radio pulsars tend to show very flat position angle swings \\citep[e.g.][]{cmk01,jhv+05}. Alternatively, the scatter could be the result of a physical process outside the neutron star that is affecting pulsar spin down, such as a supernova fallback disk as suggested by \\citet{mph01}. A supernova fallback disk was recently discovered by \\citet{wck06} around a young neutron star, though there is no indication that the disk and the neutron star are interacting.  \\psr\\ is in an emerging class of high magnetic field rotation-powered pulsars \\citep[e.g.][]{ckl+00,pkc00,km05,msk+03}. The existence of these sources raises the question of why they appear to be rotation powered instead of powered by the decay of their large magnetic fields, as are the magnetars. The magnetar model proposed by \\citet{tlk02} provides a prediction for the X-ray luminosity of a neutron star with a magnetic field of $\\sim 10^{14}$\\,G given a measurement of $n$. In the model, the neutron-star magnetosphere suffers a large scale `twist' of the north magnetic hemisphere with respect to the southern hemisphere, and resulting magnetospheric currents both scatter thermal surface photons and themselves heat the surface upon impact. Both effects result in the observed X-ray emission. The model predicts a one-to-one relationship between twist angle and $n$, as well as $L_x$. In the case of \\psr, if the above model were applicable, the measurement of $n=2.65+/-0.01$ would imply a twist angle of $\\phi \\simeq 1.2$\\,rad, which would predict an X-ray luminosity of $L_x \\simeq 3.7 \\times 10^{35}$\\,erg\\,s$^{-1}$. The observed luminosity of \\psr\\ in the 0.5-10\\,keV energy range is $L_x=4.1 \\times 10^{35}$\\,ergs\\,s$^{-1}$ for a distance of 19\\,kpc \\citep{hcg03}. This value is consistent with the magnetar model prediction. This suggests that the pulsar may actually be a magnetar. However, the estimated distance is is uncertain by a factor of two and likely overestimated, since the remnant flux and size are much larger than expected for the pulsar's age. Thus the true X-ray luminosity of \\psr\\ could be significantly less than the magnetar model prediction, though this will likely remain uncertain. In any event, since the observed X-ray flux can be easily accounted for by the spin-down luminosity of the pulsar, \\psr\\ is clearly not an `anomalous' X-ray pulsar. Moreover, the pulsar's X-ray spectrum is much harder than that of any anomalous X-ray pulsar, and inconsistent with what is predicted in the magnetar model. Indeed the spectrum is much more in line with those seen from rotation-powered pulsars \\citep{got03}. More theoretical work needs to be done to explain the difference between the magnetars and the rotation-powered pulsars having inferred magnetar-strength fields."
    },
    "0601/astro-ph0601706_arXiv.txt": {
        "abstract": "We review methods for measuring the sizes, line widths, and luminosities of giant molecular clouds (GMCs) in molecular-line data cubes with low resolution and sensitivity.  We find that moment methods are robust and sensitive --- making full use of both position and intensity information --- and we recommend a standard method to measure the position angle, major and minor axis sizes, line width, and luminosity using moment methods. Without corrections for the effects of beam convolution and sensitivity to GMC properties, the resulting properties may be severely biased. This is particularly true for extragalactic observations, where resolution and sensitivity effects often bias measured values by 40\\% or more. We correct for finite spatial and spectral resolutions with a simple deconvolution and we correct for sensitivity biases by extrapolating properties of a GMC to those we would expect to measure with perfect sensitivity (i.e.~the 0~K isosurface).  The resulting method recovers the properties of a GMC to within 10\\% over a large range of resolutions and sensitivities, provided the clouds are marginally resolved with a peak signal-to-noise ratio greater than 10.  We note that interferometers systematically underestimate cloud properties, particularly the flux from a cloud.  The degree of bias depends on the sensitivity of the observations and the $(u,v)$ coverage of the observations.  In the Appendix to the paper we present a conservative, new decomposition algorithm for identifying GMCs in molecular-line observations. This algorithm treats the data in physical rather than observational units (i.e.~parsecs rather than beams or arcseconds), does not produce spurious clouds in the presence of noise, and is sensitive to a range of morphologies. As a result, the output of this decomposition should be directly comparable among disparate data sets. ",
        "introduction": "Over the last 15 years, it has become possible to observe molecular emission in nearby galaxies with sufficient resolution and sensitivity to distinguish individual giant molecular clouds (GMCs). The immediate goal of such studies is to determine whether (and how) the GMCs in other galaxies differ from those seen in the Solar neighborhood. The most common method used to address this question has been to use molecular-line tracers of H$_2$, in particular $^{12}$CO($1\\to 0$), to compare the macroscopic properties (size, line width, and luminosity) of GMCs in other galaxies to those of Milky Way GMCs. Unfortunately, a wide variety of methods have been used to reduce data from spectral line data cubes into macroscopic GMC properties. As a result, many of the differences between GMC populations found in the literature can be attributed, at least partially, to observational artifacts or methodological differences. It is therefore difficult to assess what the real differences between GMC populations are based on the reported data in the literature.  For GMCs that are either marginally resolved or marginally detected, observational biases can be severe.  Figure \\ref{RESBIAS} shows the variation of the measured spatial size and line width with the resolution for a model cloud.  Typical Galactic GMCs have sizes of a few 10s of parsecs, comparable to the spatial resolution of many data sets used to study extragalactic GMCs \\citep[e.g.~][]{vog87,ws90,wr91,is93,wil93,fu99,she00,ros03}.  Figure \\ref{RESBIAS} shows that when the size of the beam is comparable to the size of the object, the measured size is much higher than the true size of the object. Millimeter spectrometers and correlators often have excellent frequency resolution, so the spectral resolution bias is usually less important for GMC studies, but it can become substantial when data are binned to increase signal-to-noise.  Figure \\ref{EXTRAPMOMS} shows that the measured spatial size, line width, and flux of a real GMC in M~33 \\citep[EPRB1 from][]{ros03} are all strong functions of the sensitivity of the data.  We discuss another major source of bias, the method by which emission is decomposed into GMCs, in the Appendix.  To place these biases in the the context of real molecular cloud studies, Figure \\ref{DISTCOMP} shows the peak flux and angular size of typical GMCs \\citep[those of][]{srby87} as a function of distance. The sensitivities and resolutions of a representative sample of molecular cloud studies have been indicated as horizontal lines in these plots. Distances to commonly observed objects have also been labeled.  Figure \\ref{DISTCOMP} demonstrates that most studies of extragalactic GMCs are conducted where the clouds of interest are only marginally resolved and are found at low sensitivity.  Even future observations of GMCs in the Virgo cluster using the Atacama Large Millimeter Array (ALMA) will be affected by resolution and sensitivity biases.  \\begin{figure} \\begin{center} \\plotone{fg1.eps} \\figcaption{\\label{RESBIAS} The resolution bias for a model cloud. These two panels show the variations in the measured spatial size (left) and line width (right) as a function of the resolution of a data set. These plots show measurements of Gaussian clouds convolved with a Gaussian beam and integrated across square velocity channels. Marginally resolved data yields measurements affected by a significant {\\em resolution bias}. The spatial size bias is encountered frequently in extragalactic GMC measurements.} \\end{center} \\end{figure}  \\begin{figure*} \\begin{center} \\plotone{fg2.eps} \\figcaption{\\label{EXTRAPMOMS} The sensitivity bias for a bright GMC in M~33 \\citep[EPRB1 from][]{ros03}. These four panels show the variations in the measured spatial size, line width, and flux of the GMC for a range of sensitivities (signal-to-noise ratios). We simulate the effect of changing sensitivity by varying the boundary isosurface, $T_{edge}$, and measuring properties only from data within this isosurface. All four properties are a strong function of $T_{edge}$, so the data display a substantial {\\em sensitivity bias}. The gray, dashed lines show the extrapolation to the 0 Kelvin isosurface using a weighted, least-squares fit to each property as a function of $T_{edge}$.  The extrapolated value is shown as an $\\times$.} \\end{center} \\end{figure*}  \\begin{figure*} \\begin{center} \\plottwo{fg3a.eps}{fg3b.eps} \\figcaption{\\label{DISTCOMP} Peak flux density (left) and angular size (right) of molecular clouds as a function of distance.  Diagonal lines indicate the peak flux density and size of clouds with $10^3,10^4,10^5$ and $10^6~M_{\\odot}$ (bottom to top) from the relationships observed in the inner Milky Way \\citep{srby87}.  For reference, the diameters of these clouds are roughly 3, 10, 30, and 100 pc respectively.  Horizontal black lines indicate the typical 5$\\sigma$ sensitivity (left) and effective resolution (right) of various telescopes used to observe the cloud. Observational parameters from the following studies: CfA 1.2 m: \\citet{mwco}, Massachusetts-Stony Brook (MSB) survey: \\citet{san86}, FCRAO Outer Galaxy Survey: \\citet{ogs}, NANTEN studies of the Magellanic Clouds: \\citet{fu99}, BIMA studies of M33: \\citet{ros03}, ALMA: web-based sensitivity calculator.} \\end{center} \\end{figure*}  In this paper, we examine the effects of biases stemming from finite spatial resolution, spectral resolution, and sensitivity in molecular-line observations of GMCs. We recommend data analysis methods to produce a standardized set of observed cloud properties that account for these biases.  Most of the methods used in this paper have been adopted piecemeal and {\\it ad hoc} in previous studies. Here, we endeavor to justify our choice of methodologies and to synthesize various author's techniques for approaching the problems of molecular cloud data analysis.  In Section \\ref{MOMENTS}, we describe a standardized method to measure three basic properties of an emission distribution --- size, line width, and flux --- while accounting for the sensitivity and resolution of a data set. In Section \\ref{PHYSICAL}, we discuss how these measurements can be transformed into physical quantities --- radius, line width, luminosity, and implied mass.  Finally, we consider the effects of using interferometers to derive cloud properties in Section \\ref{INTERF}. The results in these three sections are applicable to all observations of molecular clouds.  In contrast, the methods used for decomposing emission into molecular clouds vary widely, and there is little basis for favoring one method over another in all cases.  Hence, we defer a presentation of our decomposition algorithm to the Appendix of the paper, leaving only a brief discussion of the decomposition problem in Section \\ref{DECOMPSECT}.  We conclude the paper by exhibiting several examples of the application of our standardized methods to previously observed data and making recommendations for future observations (Section \\ref{APPLICATIONS}).  The methods described in this paper require a computer program to apply. A documented software version of the decomposition and measurement algorithms is available from the authors as an IDL package. ",
        "conclusions": "\\label{APPLICATIONS} We conclude the paper by applying these methods to molecular line data sets that have been previously published. In future studies, the algorithm will be used to evaluate the differences between GMC populations between galaxies. Here, we simply present an analysis designed to demonstrate the method's utility. We present the median corrections found for a large set of extragalactic (Local Group) observations and a test application of our methods to Galactic data. We use all the methods discussed in the previous sections and the decomposition algorithm discussed in the Appendix.  \\subsection{The Effects of Extrapolation and Deconvolution}  We apply the methods outlined here to an array of data from across the Local Group and measure the properties of 110 spatially resolved clouds to estimate the typical magnitude of the sensitivity and resolution corrections for extragalactic data. We use BIMA data on M~33 and M~31 \\citep{ros03,ros05}; NANTEN observations of the LMC \\citep[][]{fu99}; OVRO observations of IC~10 \\citep[][]{wa05}; and a SEST map of N83 in the SMC \\citep[][]{bol03}. Table \\ref{corrections} shows the number of clouds measured in each galaxy along with the median sensitivity and resolution corrections applied to the radius, line width, and flux. For comparison, we also measure the properties of a number of clouds in the outer Galaxy (Quadrant 2, see below) from the \\citet[][]{mwco} CO survey of the Milky Way. Table \\ref{corrections} includes all spatially resolved clouds with masses (derived from the CO luminosity) of $5 \\times 10^4$ M$_{\\odot}$ or more ($92$ of the $110$ extragalactic clouds are above this mass).  The numbers quoted in Table \\ref{corrections} are ``correction factors,''  \\begin{equation} \\frac{R_{corrected}}{R_{uncorrected}}, ~\\frac{\\sigma_{v,corrected}}{\\sigma_{v,uncorrected}},\\mbox{ and } \\frac{L_{\\mathrm{CO},corrected}}{L_{\\mathrm{CO},uncorrected}} \\mbox{ .} \\end{equation} Table \\ref{corrections} shows that throughout the Local Group data the corrections suggested in this paper have magnitudes of a few tens of percent. We draw several conclusions based on these data:  \\begin{enumerate} \\item Resolution effects on the size of clouds tend to be significant --- we would overestimate cloud sizes by $\\sim 40$\\%\\ if we did not apply a deconvolution. In the Milky Way data, this effect is much less severe. The sizes of Milky Way clouds are measured to within $5\\%$ before the resolution correction. Unresolved clouds do not contribute to Table \\ref{corrections}, so if the effects of the resolution bias were completely neglected this would be much larger (a naive approach would measure these clouds to have the size of the spatial beam).  \\item Resolution effects on the line width are negligible throughout the Local Group data.  \\item Sensitivity effects are also significant. Without a correction for the sensitivity bias, the size, line width, and luminosity of clouds would all be significantly underestimated. This sensitivity bias is least severe --- only about 20 -- 30\\%\\ --- for the line width, and most significant (and variable) for the luminosity. Sensitivity corrections to the luminosity vary from 20\\% to more than 100\\%.  \\item The resolution and sensitivity biases for the size measurement tend to cancel out, so that the completely uncorrected radius measurement is often within 10 -- 20\\%\\ of the corrected value. This is a happy coincidence of resolution and sensitivity within the Local Group, not evidence that sensitivity and resolution corrections are unimportant.  \\item The magnitude of corrections across the Local Group data are fairly uniform. This is because GMCs near the resolution limit tend to outnumber higher mass GMCs. Unresolved GMCs are not included in the analysis, so the median cloud through all the data sets appears marginally resolved.  \\item In order to compare extragalactic data to Galactic data (with very good sensitivity and resolution and therefore small corrections) it is crucial to correct for the sensitivity and resolution biases. \\end{enumerate}  \\begin{center} \\begin{deluxetable*}{l c c c c c c}  \\tablecaption{\\label{corrections} Typical Corrections for Local Group Data}  \\tablehead{ \\multicolumn{1}{l}{Galaxy} & \\multicolumn{1}{c}{$N_{Clouds}$} & \\multicolumn{3}{c}{Sensitivity Correction} & \\multicolumn{2}{c}{Resolution Correction} \\\\ \\multicolumn{1}{l}{} & \\multicolumn{1}{c}{} & \\multicolumn{3}{c}{$\\overbrace{\\phm{SpanningSpann}}$} & \\multicolumn{2}{c}{$\\overbrace{\\phm{SpanningSpann}}$} \\\\ \\multicolumn{1}{l}{} & \\multicolumn{1}{c}{} & \\multicolumn{1}{c}{$\\frac{R_{corrected}}{R_{uncorrected}}$} & \\multicolumn{1}{c}{$\\frac{\\sigma_{v,corrected}}{\\sigma_{v,uncorrected}}$} & \\multicolumn{1}{c}{$\\frac{L_{CO,corrected}}{L_{CO,uncorrected}}$} & \\multicolumn{1}{c}{$\\frac{R_{corrected}}{R_{uncorrected}}$} & \\multicolumn{1}{c}{$\\frac{\\sigma_{v,corrected}}{\\sigma_{v,uncorrected}}$} \\\\}  \\startdata LMC & 46 & $1.4$ & $1.2$ & $1.7$ & $0.8$ & $1.0$ \\\\ M~31 & 28 & $1.5$ & $1.3$ & $1.6$ & $0.7$ & $1.0$ \\\\ IC~10 & 17 & $1.7$ & $1.3$ & $2.3$ & $0.7$ & $1.0$ \\\\ M~33 & 15 & $1.4$ & $1.2$ & $1.5$ & $0.7$ & $1.0$ \\\\ SMC & 4 & $1.1$ & $1.2$ & $1.2$ & $0.7$ & $1.0$ \\\\ MW\\tablenotemark{a} & 107 & $1.1$ & $1.1$ & $1.4$ & $1.0$ & $1.0$ \\\\ \\enddata \\tablenotetext{a}{Quadrant 2 clouds with $M > 5 \\times 10^4$ M$_{\\odot}$.} \\end{deluxetable*} \\end{center}  \\subsection{Analysis of Second Quadrant CO Data}  The method described in this paper has been designed with extragalactic data in mind. However, a crucial step in interpreting extragalactic measurements is to make a fair comparison with Galactic data. In this section we report some results of applying our decomposition and measurement algorithms to the survey of the second Galactic quadrant by \\citet{mwco}. We compare the results of this analysis to the results by \\citet[][]{hc01} and show that our analysis recovers results that are consistent with theirs.  We decompose and analyze $^{12}$CO($1\\to 0$) from the second quadrant \\citep[Survey 17 in Table 1 of][]{mwco}. The data set covers the Galactic plane from $\\ell = 70^{\\circ}$ to $\\ell = 210^{\\circ}$ with a noise level of $0.3$ K. We measure the distance to the molecular emission using the kinematic distances by adopting a flat rotation curve with $\\Theta_{\\mathrm{LSR}} = 220$~km~s$^{-1}$ and $R_\\odot=8.5$~kpc. We omit local emission by discarding all elements of the data cube with a kinematic distance less than 2 kpc as well as all elements in the data cube that are connected by significant emission in position or velocity space to such pixels. We apply the decomposition algorithm described in the Appendix and measured sizes, line widths, and luminosities of GMCs using the methods of \\S \\ref{EXTRAPSECT} and \\S \\ref{BEAMDCSECT}. The analysis recovers 431 clouds with resolved angular sizes and line widths located within 10 kpc of the Sun. We include the median sensitivity and resolution corrections for massive ($>5 \\times 10^4$ M$_{\\odot}$) clouds in Table \\ref{corrections} above.  Do the results from our algorithm agree with previous studies of Galactic GMCs? The data set covers the region studied by \\citet{hc01} using the $45''$ resolution of the FCRAO 14m. That data set has a lower sensitivity than the \\citet[][]{mwco} data, so we apply a rudimentary sensitivity correction (assuming that the GMCs are Gaussian and using their peak temperature and boundary isosurface) to their results and scale to $X_{\\mathrm{CO}}=2\\times 10^{20}\\mbox{ cm}^{-2} (\\mbox{K km s}^{-1})^{-1}$ before comparing GMC properties. We focus on a comparison of the virial parameter between the two studies --- a full treatment of the Galactic ``Larson's Laws'' is beyond the scope of this paper. We draw several conclusions from the comparison:  \\begin{enumerate}  \\item Figure \\ref{virparams} shows that the virial parameters derived in our analysis largely agree with those found by \\citet[][]{hc01}. Below masses of $\\approx 3 \\times 10^4$ M$_{\\odot}$, both surveys find the same virial parameter in a given mass bin.  \\item Above $\\approx 3 \\times 10^4$ M$_{\\odot}$, our analysis of the \\citet[][]{mwco} data set may yield a slightly higher virial parameter, on average. This may be evidence for the diffuse emission mentioned in \\S\\ref{EXTRAPSECT} --- the higher sensitivity \\citet[][]{mwco} data may include diffuse emission surrounding the GMCs while the \\citet[][]{hc01} data may miss this effect. It may also reflect inadequacies in the simple sensitivity correction we apply to the \\citet[][]{hc01} data. Resolution effects may also play a role --- the \\citet[][]{hc01} data set has $\\sim 5$ times better resolution than the \\citet[][]{mwco} data and so the lower resolution data may tend to lump unbound clouds together. The number of clouds in the high mass bins is relatively low, so the discrepancy may not be particularly significant.  \\item We apply our algorithm to a small portion of the OGS data and find our corrections for resolution and sensitivity bias increase the mean virial parameter to be consistent with the results from the \\citet{mwco} data. This suggests that the differences in the virial parameters for the high mass clouds in Figure \\ref{virparams} may simply be methodological.  \\item Both catalogs of outer Galaxy clouds find an inverse relationship between luminous mass and the virial parameter, approximately $\\alpha \\propto M_{\\mathrm{Lum}}^{-0.2}$ --- a relationship that was also observed by \\citet{srby87} for inner Galaxy molecular clouds. \\end{enumerate}  Thus we find agreement with the results of previous studies of Galactic molecular clouds. Our methods applied to Galactic data find the same behavior observed in earlier work and we find agreement among the method applied to several data sets.  \\begin{figure} \\begin{center} \\plotone{fg8.eps} \\caption{\\label{virparams} Virial Parameter as a Function of Luminous Mass for Clouds in the Perseus Arm. Gray dots indicate the 431 clouds for which the virial and luminous mass can both be determined.  The solid black line is the average of the data in bins of 0.25 dex in mass.  For comparison, the average virial parameter is plotted as a function of mass for the clouds in the OGS catalog (dashed line).  The two data sets show similar virial parameters at all masses, recovering nearly the same virial parameter for a given mass.} \\end{center} \\end{figure}  \\subsection{Conclusions}  We have presented a method for measuring macroscopic GMC properties --- spatial size, line width, and luminosity --- from a region of emission in a spectral line data cube. This method corrects for biases from limited sensitivity and resolution and produces reliable results that are directly comparable among a wide variety of data sets. We correct for limited sensitivity via an extrapolation to a theoretical 0 Kelvin isosurface. We apply a simple quadratic deconvolution to the extrapolated values to account for resolution biases. We find that bootstrapping methods yield believable estimates of the uncertainties in the derived parameters. We present a set of suggestions for transforming the derived properties into physical quantities of interest. In the Appendix to this paper we present a new method for decomposing emission into individual GMCs. This method is conservative and robust, designed to produce robust results from low resolution/sensitivity data. In this section we have applied all of these methods to an array of extragalactic (Local Group) and Galactic data. We find that the algorithm reproduces established results for Galactic GMCs and that the sensitivity and resolution biases are potentially significant --- often $\\sim 40\\%$ --- for even the most recent Local Group GMC measurements.  Based on this investigation of the observational biases in measuring molecular cloud properties, we note several important points that should be considered in planning observations of GMCs and interpreting the results.  First, resolution and sensitivity biases can be corrected to $<10\\%$ error provided the cloud has a modest peak signal-to-noise ($T_{peak}/\\sigma_{RMS}\\gtrsim 10$) and is marginally resolved $R_{cld} > 0.8\\theta_{\\mathrm{FWHM}}$, where $\\theta_{\\mathrm{FWHM}}$ is the full width at half maximum extent of the beam.  Given current and future telescope capabilities, even the properties of extragalactic GMCs can be accurately measured.  From Figure \\ref{DISTCOMP}, we can see that single dish surveys can accurately study clouds more massive than $10^4~M_{\\odot}$ in the Magellanic clouds and careful interferometer observations can recover cloud properties for clouds with mass $M\\gtrsim 10^5~M_{\\odot}$ in M31 and M33.  However, interferometer observations systematically underestimate molecular cloud properties.  All else being equal, interferometers can underestimate fluxes by 40\\% and cloud radii by 10\\% relative to single-dish observations.  Line widths are largely unaffected by interferometer observations. The magnitude of the systematic bias depends on the array configuration and the sensitivity of the observations.  In general, wider coverage of the $(u,v)$ plane produces better property recovery.  To minimize bias, observers should favor observations from several array configurations or from arrays with many elements.  If possible, the interferometer data should be supplemented with single-dish observations.  Finally, the decomposition algorithm used to separate emission into physically relevant structures will systematically affect molecular cloud properties.  To date, there is no algorithm that should be favored in all circumstances, but any comparative study of GMC properties should be consistent in the choice of algorithm.  For example, the results of a CLUMPFIND algorithm applied to an extragalactic data set are not directly comparable to a catalog produced by a simple contouring method applied to Milky Way data.  Provided the same algorithm is used across multiple data sets, referencing algorithm parameters to a common physical scale will minimize systematic differences.  We will make a software version of the decomposition and measurement algorithm available as an IDL package. Further, since methodology can affect the results of GMC studies so strongly, we encourage authors working in the field to make their data available to the community after publication in order to facilitate future rigorous comparisons."
    },
    "0601/astro-ph0601060_arXiv.txt": {
        "abstract": "The Carina Nebula (NGC~3372) is our richest nearby laboratory in which to study feedback through UV radiation and stellar winds from very massive stars during the formation of an OB association, at an early phase in the evolution of the surrounding proto-superbubble before supernova explosions have influenced the environment.  This feedback is triggering successive generations of new star formation around the periphery of the nebula, while simultaneously evaporating the gas and dust reservoirs out of which young stars are trying to accrete material.  This paper takes inventory of the combined effect from all the known massive stars that power the Carina Nebula through their total ionizing flux and integrated mechanical energy from their stellar winds. Carina is close enough and accessible enough that spectral types for individual stars are available, and many close binary and multiple systems have recently been spatially resolved, so that one can simply add them.  Adopting values from the literature for corresponding spectral types, the present-day total ionizing photon luminosity produced by the 65 O stars and 3 WNL stars in Carina is Q$_H\\simeq$10$^{51}$ s$^{-1}$, the total bolometric luminosity of all stars earlier than B2 is 2.5$\\times$10$^7$ $L_{\\odot}$, and the total mechanical luminosity of stellar winds is $L_{SW}\\simeq$10$^5$ $L_{\\odot}$.  The total Q$_H$ was about 25\\% higher when $\\eta$~Carinae was on the main sequence, before it and its companion were surrounded by its obscuring dust shell; for the first 3 Myr, the net ionizing flux of the 70 O stars in Carina was about 150 times greater than in the Orion Nebula.  About 400-500 $M_{\\odot}$ has been contributed to the H~{\\sc ii} region by stellar wind mass loss during the past 3 Myr.  Values for Q$_H$ and $L_{SW}$ are also given for the individual clusters Tr14, 15, and 16, and Bo10 and 11, which are more relevant on smaller spatial scales than the total values for the whole nebula. ",
        "introduction": "Feedback from young massive stars may play an integral role in star and planet formation.  Most stars are born in the vicinity of the hottest and most massive stars spawned only from giant molecular clouds, and the effects of feedback from these massive stars cannot be studied in nearby quiescent regions of star formation.  Even the nearest H~{\\sc ii} region, the Orion Nebula, may not be representative of the extreme environments where most stars are born, since it is dominated by just a single O6 dwarf.  A much more extreme collection of stars can be found in the Carina Nebula (NGC~3372).  Its star clusters are almost as spectacular as those of 30~Dor in the LMC (Massey \\& Hunter 1998), or objects in our own Galaxy like the Arches cluster near the Galactic center (Najarro et al.\\ 2004; Figer et al.\\ 1999, 2002) and NGC~3603 (Moffat et al.\\ 2002).  However, these other regions are too distant for detailed studies of small-scale phenomena like irradiated protoplanetary disks and jets, and their study is hampered by considerably more extinction. The low extinction toward Carina combined with its proximity and rich nebular content provide a worthwhile trade-off for its less concentrated star clusters.  While this may be a liability for investigating the upper end of the initial mass function, the fact that Carina's stars are more of a loose aggregate may be an advantage for studying feedback and triggered star formation.  Carina provides a snapshot of an OB association in the making, which may be more representative of the environments in which most stars form than are bound super star clusters.  It may also provide an early analog of the local Galactic environment that evolved into Gould's Belt.  Recent studies suggest that at the present epoch there exists an upper mass limit to the most massive stars of $\\sim$150 M$_{\\odot}$ (Figer 2005, Weidner \\& Kroupa 2004, Kroupa \\& Weidner 2005; Oey \\& Clarke 2005).  While the star clusters that power Carina may not be as dense as the Arches cluster or NGC~3603, Carina does contain several examples of stars that probably began their lives with 100--150 M$_{\\odot}$.  Among these are $\\eta$~Carinae, the prototypical O2 supergiant HD~93129A, three late-type hydrogen-rich Wolf-Rayet stars (WNL stars), and the remaining original members of the O3 spectral class for which this spectral type was first introduced (Walborn 1973; although see Walborn et al.\\ 2002).  For these colossal stars, lifetimes of $\\sim$3 Myr before they explode as supernovae (SNe) are shorter than the time it takes to clear away their natal molecular material.  Consequently, we already have a situation in Carina where the most massive members like $\\eta$~Car are approaching their imminent demise while new stars are being born from dense molecular gas only 5--20 pc away.  In the next 1--2 Myr, there will be several very energetic SNe in the Carina Nebula, which will carve out an even larger cavity in the ISM and form a giant superbubble in the Galactic plane, and may pollute protoplanetary disks with nuclear-processed ejecta.  In the mean time, studying this rich region provides a snapshot of a young proto-superbubble (Smith et al.\\ 2000) energized only by UV radiation and stellar winds, immediately before the disruptive SN-dominated phase.  In just the past decade, the Carina Nebula has been recognized as a hotbed of ongoing active star formation.  It provides a laboratory to study several star formation phenomena in great detail, all of which have recently been identified in this region: (1) evaporating protoplanetary disks (the so-called ``proplyds''), small cometary clouds, or globules (Smith et al.\\ 2003$a$; Smith 2002$b$; Brooks et al.\\ 2000, 2006; Cox \\& Bronfman 1995), (2) the erosion of large dust pillars and the triggering of a second generation of embedded star formation within them (Smith et al.\\ 2000, 2005; Rathborne et al.\\ 2004; Megeath et al.\\ 1996), (3) irradiated Herbig-Haro (HH) jets that are a signpost of active accretion (Smith et al.\\ 2004), (4) photodissociation regions (PDRs) on the surfaces of molecular clouds across the region (Brooks et al.\\ 2003, 2006, 1998; Rathborne et al.\\ 2002; Smith et al.\\ 2000; Mizutani et al.\\ 2004), and on the largest scales, (5) the early formation of a fledgeling superbubble (Smith et al.\\ 2000).  These phenomena trace a second generation of stars, perhaps triggered by feedback from the first generation clusters. Carina is near enough that an upcoming program with the {\\it Hubble Space Telescope} will undertake the same types of detailed studies of feedback that have been done in Orion at visual and near-IR wavelengths, but sampling a more extreme environment.  All of these phenomena respond to the cumulative UV radiation and/or stellar winds from star clusters that power the H~{\\sc ii} region, and a census of this energy input is essential in order to understand the relationship between them.  Estimating the total stellar energy input in Carina is a more difficult problem than in smaller H~{\\sc ii} regions like Orion, where a single star dominates the photoevaporation of proplyds and the irradiation of HH jets.  Instead, Carina has dozens of massive O stars over many parsecs.  Although the remarkable stellar content of the Carina Nebula has been discussed many times in the literature (e.g., Walborn 1995, 2002, 2005; Feinstein 1995; Massey \\& Johnson 1993), there is no complete and up-to-date census of the massive stars, and especially there has been no estimate of their {\\it collective} UV radiation and mechanical (wind) luminosity that powers the region.  This is the main purpose of the present investigation.  A second paper in this series will compare the total energy input from stars with the apparent energy budget inferred from observations of the surrounding Carina Nebula. ",
        "conclusions": ""
    },
    "0601/astro-ph0601583_arXiv.txt": {
        "abstract": "We present \\SPEARFIMS\\ far-ultraviolet observations near the North Ecliptic Pole. This area, at $b \\sim$ 30\\degrees\\ and with intermediate \\HI\\ column, seems to be a fairly typical line of sight that is representative of general processes in the diffuse ISM. We detect a surprising number of emission lines of many elements at various ionization states representing gas phases from the warm neutral medium (WNM) to the hot ionized medium (HIM). We also detect fluorescence bands of \\Htwo, which may be due to the ubiquitous diffuse \\Htwo\\ previously observed in absorption. ",
        "introduction": "The North Ecliptic Pole (NEP) is a region of the sky with no obviously unusual features. Many prior space-missions (\\Einstein, \\ROSAT, \\COBE, \\IRAS, etc.) have conducted deep surveys of this region. Several ground based surveys have also concentrated on this region as a complement to the space-based surveys \\citep{labov89,elvis94}. Its location at moderate galactic latitude (b=+29\\degrees) places it clear of the bright stars of the galactic plane. In this region, the Galactic N(H{\\sc I}) varies from about 2\\tten{20}\\persqrcm\\ to 8\\tten{20}\\persqrcm which corresponds to an dust opacity of \\taudust=0.6-2.3 at 1032\\AA \\citep{sasseen02} or \\taudust=0.3-1.4 at 1550\\AA \\citep{savage79}.  The Spectroscopy of Plasma Evolution from Astrophysical Radiation (\\SPEAR) instrument, designed for observing emission lines from the diffuse ISM was launched in late 2003. \\SPEAR\\, is a dual-channel FUV imaging spectrograph (short-$\\lambda$ (S) channel 900 - 1150 \\AA, long-$\\lambda$ (L) channel 1350 - 1750 \\AA) with $\\lambda/\\Delta\\lambda\\sim$550, with a large field of view (S: 4.0$^{\\circ} \\times$ 4.6', L: 7.5$^{\\circ} \\times$ 4.3') imaged at 10$\\arcmin$ resolution. See \\cite{edelstein05a, edelstein05b} for an overview of the instrument, mission, and data analyses. \\SPEAR\\ sky-survey observations consist of sweeps at constant ecliptic longitude from the NEP to the south ecliptic pole through the anti-solar point.  The duration of each sweep varies between 900 and 1500 s. Each survey scan includes the region near the north and south ecliptic poles resulting in large exposure times in these regions. We report on deep \\SPEAR\\ observations of FUV emission spectra from a 15\\degree\\ radius region centered on the NEP. ",
        "conclusions": "We have presented the FUV spectra of a 30\\degree\\ region around the NEP.  We detect a variety of emission lines from high and low ionization gas.  The high ionization lines, taken on a per species basis, are consistent with EM from 0.001 to 0.005 \\emunits\\ throughout $T=10^{4.5}$ to 10$^{5.5}$ K range.   For those few species which have multiple lines in the band, the calculated temperatures tend to fall near the CIE peak abundance for the species.  Despite this, a linear combination of CIE models with a restricted abundance vs $T$ gradient does not provide a good fit to the spectrum.  This is likely due to non-CIE effects and abundance variations are not directly $T$ related.  It is also likely that many lines include some resonantly scattered stellar radiation.  Modeling this scattering will be a significant portion of our continued research.  We have discerned line emission from species that are expected to exist in the WNM and WIM.  We interpret these lines as resonantly scattered stellar radiation.  In the future, we will investigate if spatial variation of the \\AlII\\ and \\SiII\\ features can indicate the effect of grain formation and destruction on the gas phase abundances of these elements.  Despite relatively low galactic $N(HI)$ in this direction, we see a substantial amount of \\Htwo\\ fluorescence.  Further modeling of the \\Htwo\\ fluorescence is necessary to fully understand faint spectral features which could be related to \\Htwo\\ or due to unrelated atomic lines."
    },
    "0601/astro-ph0601256_arXiv.txt": {
        "abstract": "We use high quality VLT/UVES published data of the permitted  \\ion{O}{1} triplet and \\ion{Fe}{2} lines to determine oxygen and iron abundances in unevolved (dwarfs, turn-off, subgiants) metal-poor halo stars. The calculations have been performed both in LTE and NLTE, employing effective temperatures obtained with the new infrared flux method (IRFM) temperature scale by Ram\\'{\\i}rez \\& Mel\\'endez, and surface gravities from {\\it Hipparcos} parallaxes and theoretical isochrones. A new list of accurate transition probabilities for \\ion{Fe}{2} lines, tied to the absolute scale defined by laboratory measurements, has been used. Interstellar absorption has been carefully taken into account by employing reddening maps, stellar energy distributions and Str\\\"omgren photometry.  We find a plateau in the oxygen-to-iron ratio over more than two orders of magnitude in iron abundance ($-3.2 <$ [Fe/H] $< -0.7$), with a mean [O/Fe] = 0.5 dex ($\\sigma$ = 0.1 dex), independent of metallicity, temperature and surface gravity. The flat [O/Fe] ratio is mainly due to the use of adequate NLTE corrections and the new IRFM temperature scale, which, for metal-poor F/early G dwarfs is hotter than most \\teff scales used in previous studies of the \\ion{O}{1} triplet.  According to the new IRFM \\teff scale, the temperatures of turn-off halo stars strongly depend on metallicity, a result that is in excellent qualitative and quantitative agreement with stellar evolution calculations, which predict that the \\teff of the turn-off at [Fe/H] = $-3$ is about 600-700 K higher than that at [Fe/H] = $-1$. Recent determinations of H$\\alpha$ temperatures in turn-off stars are in excellent relative agreeement with the new IRFM \\teff scale in the metallicity range $-2.7 <$ [Fe/H] $< -1$, with a zero point difference of only 61 K. ",
        "introduction": "The importance of the oxygen abundance in metal-poor stars has been strongly emphasized in the literature, as well as the problems related to its determination. There is currently no consensus as to whether the [O/Fe] ratio in halo stars is approximately constant ([O/Fe] $\\approx$ 0.4 - 0.5) or steeply increases with decreasing [Fe/H] ([O/Fe]$\\approx$ -0.35 [Fe/H]). The discrepancy is due to problems in the modeling of the different spectral features used to estimate the oxygen abundance, which are very hard to detect in the same star.  Due to its high excitation potential, the permitted \\ion{O}{1} triplet at 0.77$\\mu$m is mainly observed in FG dwarfs and subgiants, being very faint (yet detectable) in metal-poor early K stars. The low excitation potential forbidden lines [\\ion{O}{1}] at 0.63 $\\mu$m are detected in giants and cool subgiants, and barely detectable in dwarfs. Molecular OH lines are observed in the ultraviolet (electronic transitions at 0.3 $\\mu$m) and in the infrared (vibrational-rotational lines at 1.5 and 3 $\\mu$m). Both observations are difficult: the UV OH lines are close to the atmospheric cutoff and most CCDs have low sensitivity in the UV, while the weak IR OH lines require extremely high S/N. The UV OH lines are relatively strong and observed in FGK metal-poor stars, while the very weak IR OH lines are only observed in cool stars with \\teff $<$ 5000 K.  The observed spectra by themselves do not allow a direct measurement of the oxygen abundance, so a careful modeling must be performed. Since the spectral features are sensitive to both effective temperature and surface gravity, reliable atmospheric parameters must be used. The analysis is complicated by the presence of both NLTE (mainly affecting the \\ion{O}{1} triplet) and granulation (crucial for the molecular OH lines) effects.  Analyses of the forbidden [\\ion{O}{1}] (e.g. Barbuy 1988; Sneden et al. 1991) and infrared OH lines (Mel\\'endez, Barbuy \\& Spite 2001) seem to show a flat [O/Fe], or a small increase of about 0.1-0.2 dex per a factor of 10 (1 dex) in metallicity, from [O/Fe]$\\approx$ 0.3 at [Fe/H] = $-$1.5 to [O/Fe] $\\approx$ 0.4-0.5 at [Fe/H] = $-$2.5 (Sneden \\& Primas 2001; Mel\\'endez \\& Barbuy 2002). Yet, the 1D analysis of the forbidden line by Nissen et al. (2002) shows a continuous increase of [O/Fe] for lower metallicities, reaching [O/Fe] $\\approx$ +0.7 at [Fe/H] = $-2.5$, but after correcting for 3D effects they found a plateau at [O/Fe] $\\approx$ +0.3 in the range $-$2.0 $<$ [Fe/H] $<$ $-$0.7, and an increase in [O/Fe] for lower [Fe/H]. The recent 1D analysis of [\\ion{O}{1}] in subgiants by Garc\\'{\\i}a P\\'erez et al. (2005) shows a mean [O/Fe] $\\approx$ +0.5, with a very small increase of 0.09 dex in [O/Fe] for a decrease of 1 dex in [Fe/H].  On the other hand, studies of the permitted \\ion{O}{1} triplet (e.g. Abia \\& Rebolo 1989; Cavallo, Pilachowski \\& Rebolo 1997; Israelian et al. 1998, 2001; Mishenina et al. 2000) and the ultraviolet OH lines (Israelian et al. 1998, 2001; Boesgaard et al. 1999) show a steep monotonic increase of [O/Fe] for lower metallicities, reaching [O/Fe] = +1.1 at [Fe/H] = $-3$. However, other studies of the triplet show a lower (or zero) slope. For example, the analysis of the triplet in turn-off stars by Akerman et al. (2004, hereafter Ake04) shows only a mild increase of [O/Fe] with metallicity, with [O/Fe] $\\approx$ +0.4 at [Fe/H] = $-$1 and [O/Fe]$\\approx$ +0.7 at [Fe/H] = $-$3.  Asplund \\& Garc\\'{\\i}a P\\'erez (2001) have shown that 1D model atmospheres overestimate the abundances obtained from UV OH lines by as much as 0.9 dex at [Fe/H] = $-3$ (see also $\\S$8.2 and Asplund 2005), and when the UV OH lines are analyzed with 3D model atmospheres, the [O/Fe] ratio in metal-poor dwarfs is in good agreement with the results obtained in giants from the forbidden lines. Note also that the 1D analysis of the UV OH lines by Bessell et al. (1991) resulted in low [O/Fe] ratios, as is also the case of the work on UV OH lines in subgiant stars by Garc\\'{\\i}a P\\'erez et al. (2005), which shows 1D oxygen-to-iron ratios of [O/Fe] $\\approx$ +0.5 dex. Hence, analyses of three oxygen abundance features (UV OH lines, IR OH lines, [\\ion{O}{1}] lines) roughly agree in a somewhat flat (and low) [O/Fe] ratio in the metallicity range $-3 <$ [Fe/H] $< -$1. However, the high oxygen abundances obtained from the \\ion{O}{1} triplet have been difficult to reconcile with the low oxygen abundance derived from the forbidden lines (Spite \\& Spite 1991). Nissen et al. (2002) found that 1D analysis results in agreement between the [\\ion{O}{1}] line and the triplet, but when the lines are analyzed employing 3D model atmospheres the oxygen abundances obtained from the triplet are about 0.2 dex larger than those from the [\\ion{O}{1}] line (Nissen et al. 2002).  The disagreement between the \\ion{O}{1} triplet and the forbidden [\\ion{O}{1}] lines also occurs in metal-poor giants and subgiants (Cavallo et al. 1997; Fulbright \\& Johnson 2003; Garc\\'{\\i}a P\\'erez et al. 2005). Israelian et al. (2004) undertook the first NLTE analysis in very metal-poor CN-rich giants, showing that there are serious problems with standard 1D model atmospheres of those stars, even when NLTE effects are taken into account, producing a serious conflict between the oxygen abundances obtained from the forbidden line and the triplet (Israelian et al. 2004). The use of CN-enhanced model atmospheres have an important impact on the thermal structure of model atmospheres of metal-poor CN-enhanced ([C/Fe] = +2, [N/Fe] = +2) giants (Masseron et al. 2005), which affect the abundances from a few tenths of a dex up to 1.5 dex in extreme cases, with respect to standard (solar-scaled abundances) 1D models (Masseron et al. 2005). Hence, for metal-poor cool giants with extreme compositions, model atmospheres computed with the corresponding CNO overabundances may partly relieve the problems found by Israelian et al. (2004).  The use of 3D + NLTE seems to improve the situation for the oxygen abundances derived from the \\ion{O}{1} triplet and the [\\ion{O}{1}] line. Employing 1.5D NLTE calculations in a 3D model atmosphere of the metal-poor star HD 140283, Shchukina, Trujillo Bueno \\& Asplund (2005) have found that the oxygen abundances obtained from the \\ion{O}{1} and [\\ion{O}{1}] lines agree within $\\approx$ 0.1 dex, which is an important achievement considering the observational uncertainties for the [\\ion{O}{1}] line.  Previous studies of the \\ion{O}{1} triplet employing a hot temperature scale resulted in a flat and low [O/Fe] (King 1993; Carretta, Gratton \\& Sneden 2000). Note that another reason why Carretta et al. (2000) obtained low [O/Fe] ratios is the use of large NLTE corrections of the \\ion{O}{1} triplet by Gratton et al. (1999), as pointed out by Takeda (2003).  King (2000) studied the influence of stellar parameters and NLTE iron abundances (Th\\'evenin \\& Idiart 1999) on the [O/Fe] ratio obtained from UV OH, [\\ion{O}{1}] and \\ion{O}{1} lines. The [O/Fe] values obtained by King (2000) using the \\ion{O}{1} triplet and the UV OH lines are considerably reduced with respect to the high [O/Fe] ratios of the original studies (Tomkin et al. 1992; Boesgaard et al. 1999), but they are still about 0.1-0.2 dex larger than those obtained from the [\\ion{O}{1}] lines (King 2000). Partly the low [O/Fe] ratios obtained by King (2000) can be explained by the large NLTE \\ion{Fe}{1} corrections by Th\\'evenin \\& Idiart (1999), which increase the iron abundance by +0.2-0.3 dex, hence lowering [O/Fe] by 0.2-0.3 dex.  The preliminary NLTE analysis of the \\ion{O}{1} triplet by Primas et al. (2001) also resulted in a flat [O/Fe] ratio for halo dwarfs down to [Fe/H] = $-$2.4, independent of the \\teff scale employed. Primas et al. found [O/Fe] $\\approx$ +0.4 using the \\teff scale by Alonso et al. (1996), and [O/Fe] $\\approx$ +0.5 with both the \\teff scale by Carney (1983) and temperatures from Th\\'evenin \\& Idiart (1999).  The observed [O/Fe] ratio in halo stars provides tight constraints on models of Galaxy formation (e.g. Tinsley 1979; Wheeler, Sneden \\& Truran 1989; McWilliam 1997; Chiappini, Romano \\& Matteucci 2003). If the formation of the halo was fast, then the [O/Fe] ratio of halo stars should be roughly constant, because Type Ia supernovae, which originate from long lived low mass stars, would not have had enough time to lower the [O/Fe] ratio of the ISM from which the halo stars we observe today were formed. Furthermore, the scatter of the [O/Fe] ratio tells us the efficiency of the ISM in homogenizing (mixing) the ejecta of Type II supernovae (Scalo \\& Elmegreen 2004).  The oxygen abundance is extremely important in studies of the production of LiBeB by Galactic cosmic ray (GCR) spallation (e.g. Prantzos, Cass\\'e \\& Vangioni-Flam 1993; Fields \\& Olive 1999; Ramaty et al. 2000; Vangioni-Flam, Cass\\'e \\& Audouze 2000), especially for stars with [Fe/H] $< -2$, where current models struggle to explain the non-zero isotopic $^6$Li/$^7$Li ratios recently found in metal-poor dwarfs down to [Fe/H] = $-2.7$ (Asplund et al. 2005). The GCR production of $^6$Li depends on the adopted [O/Fe] ratio, and only models with extremely high oxygen abundances can account for the $^6$Li detection at [Fe/H] $= -2.7$, but introduce the problem of large overproduction of $^6$Li at higher metallicities (Rollinde et al. 2005). Given the problems faced by GCR models to explain the $^6$Li observed in halo stars, several authors have suggested that the origin of $^6$Li may be pre-Galactic (Jedamzik 2000, 2004; Asplund et al. 2005; Rollinde et al. 2005). A low [O/Fe] ratio in metal-poor stars brings support for pre-Galactic $^6$Li, since in this case GCR can not produce enough $^6$Li (see also Prantzos 2005).  Here, we analyze recent high quality VLT/UVES data of the \\ion{O}{1} triplet and \\ion{Fe}{2} lines from Ake04 and  Nissen et al. (2004, hereafter N04), respectively. We show that the use of NLTE corrections and the new infrared flux method (IRFM) temperature scale by Ram\\'{\\i}rez \\& Mel\\'endez (2005a,b; hereafter RM05a,b), which for metal-poor turn-off stars is hotter than previous temperature scales, results in a low and flat [O/Fe] ratio in metal-poor stars.  We also compare the IRFM temperature of turn-off stars (employing the new IRFM \\teff scale) with those expected from stellar evolution, showing that the strong metallicity dependence of the turn-off temperatures are very well reproduced by stellar evolution models. ",
        "conclusions": "We have determined oxygen and iron abundances in 31 metal-poor ($-3.2 <$ [Fe/H] $< -0.7$) stars close to the turn-off, employing high resolution high S/N UVES data taken from Ake04 and N04.  We find a flat [O/Fe] = +0.5, independent of metallicity, temperature and surface gravity in the ranges $-3.2 <$ [Fe/H] $< -0.7$, 5700 K $<$ \\teff $<$ 6700 K, and $3.7 <$ log $g$ $< 4.5$, respectively. Our work confirms previous studies (e.g. Carretta et al. 2000; Primas et al. 2001) which have already found a flat and low [O/Fe] ratio in halo dwarfs from the \\ion{O}{1} triplet, extending the constancy of the [O/Fe] ratios down to [Fe/H] = $-$3.2.  The flat [O/Fe] ratio is mainly due to the use of adequate NLTE calculations and the new IRFM \\teff scale by RM05a,b, which for metal-poor turn-off stars is hotter than most previous \\teff scales available in the literature. We find a low star-to-star scatter of 0.136 dex for the whole sample, or $\\sigma$ = 0.10 dex for the sample with low reddening uncertainty.  The observed star-to-star scatter ($\\sigma_{obs}$ = 0.136 dex) is almost completely explained by errors in the analysis ($\\approx$ 0.121 dex), leaving little room for intrinsic scatter. Hence, the Galaxy was extremely efficient in mixing the chemical elements ejected by supernovae. Other recent works in the literature have also found very small star-to-star scatter for other $\\alpha$-elements (e.g. Cayrel et al. 2004; Cohen et al. 2004; Arnone et al. 2005). Furthermore, the low scatter implies a small contribution to the halo from metal-poor stars that originated in dSph galaxies, since much lower [$\\alpha$/Fe] ratios are commonly seen in such galaxies (e.g. Venn et al. 2004; see also discussion and references given in Catelan 2005).  The constancy of the [O/Fe] ratio over more than two orders of magnitude in [Fe/H], from [Fe/H] = $-$0.7 to [Fe/H] = $-$3.2, is telling us that the formation of the halo was extremely fast, with a timescale shorter than the bulk of Type Ia SNe. Our data provides tight constraints for Galactic chemical evolution models (e.g. Matteucci \\& Recchi 2001; Alib\\'es, Labay \\& Canal 2001; Goswami \\& Prantzos 2000; Portinari, Chiosi \\& Bressan 1998; Samland 1998; Chiappini, Matteuci \\& Gratton 1997).  Our low [O/Fe] = +0.5 constrains the $^6$Li production by GCR models, which are not able to explain the detection of $^6$Li in a star with [Fe/H] = $-$2.7 (Asplund et al. 2005), hence reinforcing the case for a pre-Galactic origin of the recent $^6$Li detections in very metal-poor stars.  Recent determinations of H$\\alpha$ temperatures in turn-off stars by Asplund et al. (2005) are in very good relative agreeement with the RM05b \\teff scale in the metallicity range $-2.7 <$ [Fe/H] $< -1$, with a zero point difference of only 61$\\pm$62 K.  As shown in Figs. 11 and 12, the strong metallicity dependence of the temperature of turn-off stars is not unphysical, but a natural consequence of stellar evolution. Ram\\'{\\i}rez et al. (2006), has recently confirmed the hot \\teff scale of RM05b for one star, BD +17 4708, by fitting its observed absolute flux distribution from Hubble Space Telescope observations with Kurucz models. Indeed, the \\teff obtained by Ram\\'{\\i}rez is about 100 K higher than the \\teff obtained in this work with the \\teff scale of RM05b. This is not in conflict with stellar evolution calculations, since there is still room for an increase of about 100 K in the RM05b \\teff scale of metal-poor turn-off stars (Fig. 12).  In the future, it would be important to take into account also granulation effects, performing full 3D+NLTE analyses. This is the way to go for future abundance studies. Also, it is very important to perform future verifications of the \\teff scale of RM05b."
    },
    "0601/astro-ph0601126_arXiv.txt": {
        "abstract": "We present three dimensional smoothed particle hydrodynamics (SPH) calculations of warped accretion discs in X-ray binary systems. Geometrically thin, optically thick accretion discs are illuminated by a central radiation source. This illumination exerts a non-axisymmetric radiation pressure on the surface of the disc resulting in a torque that acts on the disc to induce a twist or warp. Initially planar discs are unstable to warping driven by the radiation torque and in general the warps also precess in a retrograde direction relative to the orbital flow. We simulate a number of X-ray binary systems which have different mass ratios using a number of different luminosities for each. Radiation-driven warping occurs for all systems simulated. For mass ratios q $\\sim 0.1$ a moderate warp occurs in the inner disc while the outer disc remains in the orbital plane (c.f. X$\\thinspace 1916-053$). For less extreme mass ratios the entire disc tilts out of the orbital plane (c.f. Her X-1). For discs that are tilted out of the orbital plane in which the outer edge material of the disc is precessing in a prograde direction we obtain both positive and negative superhumps simultaneously in the dissipation light curve (c.f. V603 Aql). ",
        "introduction": "There are a growing number of identified X-ray binaries that have periodic variabilities in their light curves that are significantly longer than their binary orbital periods. These `super-orbital' periods have been found in Her X-1, SS 433 and LMC X-4 for example \\cite{ClarksonEt:2003}. A number of different explanations have been proposed for these super-orbital periodicities, including the precession of a tilted or warped accretion disc (Katz 1973; Petterson 1975; Wijers \\& Pringle 1999 and references therein). X-rays from close to the accreting object are reprocessed and re-emitted at the surface of the accretion disc. As the disc presents a changing aspect toward the observer, so light variations are seen. Variations in the amount of the donor star eclipsed by the warped disc and the different X-ray illumination of the donor (Gerend \\& Boynton 1976) also generate modulations in the light curves. A tilted disc will also obscure the central X-ray source from the observer and generate `dips' in the X-ray light curves as seen for example in X1916-053 \\cite{HomerEt:2001}.  Tananbaum et al. (1972) interpreted the 35-day period in the X-ray flux from Her X-1 as resulting from an accretion disc that is precessing and tilted with respect to the orbital plane. Katz (1973) suggested that the precession was forced by a torque from the donor star, although the mechanism for the misalignment of the disc out of the orbital plane was unclear. Other mechanisms have been suggested for driving the warps such as wind torques \\cite{SchandlMeyer:1994}.  The discovery of the precessing relativistic jets found in SS 433 (see Margon 1984 for a comprehensive review) indicated that an accretion disc was present which was precessing with a period of 164 days. Moreover, the morphology of the radio jets and optical photometry of the system, in combination with the jet precession, all indicate that the entire disc precesses at the same rate. Global precession of the disc is also evident for Her X-1 (Tr\\\"umper et al. 1986; Petterson et al. 1991).  Radiation pressure warping of an initially coplanar accretion disc was first detailed by Petterson (1977). Iping \\& Petterson (1990) performed detailed numerical simulations of warped discs, suggesting that radiation pressure determines the shape of the disc and its precession rate. If radiation from the central source is absorbed at the surface of the accretion disc and then re-emitted normal to this surface, a torque will be induced on the disc (Petterson 1977 ; Pringle 1996). For values of the central radiation luminosity sufficient to overcome the viscous forces within the disc, a warp will develop and the disc will tilt out of the orbital plane (Pringle 1996).  Pringle (1996, 1997) investigated analytically and numerically the effects of radiation-warping on an initially flat disc using a linear analysis for isothermal $\\alpha$-discs. The disc was modelled using a set of concentric rings centred on the primary object. Neighbouring rings exchanged angular momentum by means of viscous torques. Mass and angular momentum were conserved, but there was no detailed consideration of the internal fluid dynamics of the disc.  Pringle concluded that an accretion disc is unstable to tilting and warping due to irradiation reaction forces when the luminosity of the central source exceeds a critical value. Wijers \\& Pringle (1999) extended Pringle's work and addressed the non-linear evolution of the instability and found very different and complex warp development.  Larwood et al. (1996) investigated tidally induced warping and precession of accretion discs in binary systems. They used a smoothed particle hydrodynamics (SPH) code to investigate how a thin accretion disc that is warped and inclined to the orbital plane can survive in close binary systems. They were able to demonstrate disc warping and precession in the tidal field of the companion. The disc was only moderately warped and precessed approximately like a rigid body. A similar study of warped circumbinary discs was performed by Larwood \\& Papaloizou (1997). Nelson \\& Papaloizou (1999) studied the dynamics of warped discs orbiting a rotating black hole where the disc and black hole angular momentum vectors where misaligned. In this case the inner regions of the disc became warped.  In these various investigations a number of different types of disc warp have been found, including stable, unstable and time-independent modes. In most cases the warp precesses relative to the inertial frame. The precession rate is generally constant and the warp precesses as a rigid body by twisting the disc such that the radiation pressure causes each radial annulus to precess at the same rate. The rate of growth of the warp is dependent on the outer disc boundary conditions. The warp generally starts in the outer parts of the disc and propagates inward.  Murray et al. (2002) modelled an initially planar disc subject to an inclined magnetic dipole field centred on the donor star. They used a SPH code to model the disc response to a magnetic field that was fixed in the binary frame. They found a warp developed that precessed uniformly in a retrograde direction relative to the direction of the disc flow. The amplitude and structure of the warp was dependent on the relative phase with respect to the magnetic field. The induced warp was at maximum amplitude when it was reinforced by the magnetic force and a minimum when the warp and magnetic force were anti-phased, giving an almost flat disc.  In this paper we present a numerical study of radiation-driven warping for a number of X-ray binary systems. We investigate the problem using a non-linear SPH code originally developed by Murray (1996, 1998). In particular we aim to establish to what extent a thin, initially coplanar disc will warp and precess subject to irradiation from a central source.  In section \\ref{sec:Basic Equations} and \\ref{sec:Numerical method} we present our formulation of the problem using SPH. Section \\ref{sec:Simulations} details the binary systems modelled, and in Section \\ref{NUMERICAL RESULTS} we present our results. Section \\ref{sec:Conclusions} is devoted to conclusions. ",
        "conclusions": "\\label{sec:Conclusions}  \\subsection{Algorithms} We have studied the radiation-driven warping of accretion discs in X-ray binaries using a non-linear smoothed particle hydrodynamics code. A convex hull algorithm was developed to find the surface particles of the disc and a ray-tracing algorithm used to evaluate regions of self-shadowing. Initially planar accretion discs were illuminated with radiation from the centre of the discs. The surface of each disc was subject to a force from the absorption of the X-rays on the disc surface and back reaction from the remitted radiation.  \\subsection{Warped disc structure} Our results indicate that twists and warps develop which undergo near rigid body precession in a retrograde direction relative to the inertial reference frame. These warps survive for many orbital periods in close binary systems, with the integrity of their local vertical structure maintained. The entire disc may also become inclined relative to the orbital binary plane.  We found that for the extreme mass ratio systems, a strong warp develops close to the centre of the disc to the extent that some of the outer parts are completely shadowed by the inner warp. This warp shadow and the viscous forces combine to keep the disc in the orbital plane and the inclination of the disc to the orbital plane is small. The warp precesses in a wave-like manner in a retrograde direction relative to the inertial frame. The disc as a whole still precesses in a prograde direction similar to a non-warped disc and superhumps are still apparent in the simulated dissipation light curve (c.f. Foulkes et al.).  For the less extreme mass ratio systems a small warp is induced in the inner regions of the disc. This warp does not shadow the outer regions of the disc and these regions also warp. The entire disc then becomes tilted to the binary plane and the whole disc precesses in a retrograde direction. The warp behaviour we find appears therefore to be consistent with the differing observational characteristics of the extreme mass ratio system X\\thinspace 1916-053 and the non-extreme q system Her \\thinspace X-1.  \\subsection{Positive and negative superhumps} In the less extreme mass ratio systems, as the tilt of the disc becomes larger the gas stream no longer impacts on the edge of the disc. Instead the stream now impacts deeper in the disc near the circularization radius and the simulated dissipation light curve shows two peaks per orbital period occuring when the stream-disc impact point is deepest in the gravitational well.  Systems that have eccentric accretion discs in which the edge of the disc precesses in a prograde direction and the entire disc is inclined to the orbital plane, positive and negative superhumps are visible in the simulated disc dissipation light curves. The positive superhumps are generated by the interaction of the edge of the disc and the material from the gas stream. Since the edge material in the disc precesses in a prograde direction the period of the positive superhumps is slightly longer than the orbital period. When the whole of the disc is also marginally inclined out of the orbital plane and the entire disc warp is precessing in a retrograde direction, then negative superhumps are generated by some of the gas stream material over or under flowing the edge of the disc. The period of the negative superhumps is slightly shorter than the orbital period.  \\subsection{Radiation field strength and behaviour at $L \\sim L_{Edd}$}  The radiation field strength throughout the disc was measured using the dimensionless parameter $F*$ as defined by Wijers \\& Pringle (1999) (their equation (8))  \\begin{equation}\\label{eq:f*} F* = \\frac{L_{*}}{6 \\pi c r \\Sigma \\Omega \\nu}. \\end{equation}  \\noindent where $L_{*}$ is central source luminosity, $c$ is the speed of light, $r$ is the distance from the compact object,  $\\Sigma$ is the disc surface density at radius $r$, $\\Omega$ is the angular velocity at radius $r$ and $\\nu$ is the shear viscosity. This dimensionless number gives the viscous time scale at a point in the disc divided by the radiation time scale. Wijers \\& Pringle (1999) found that $F*$ had to be greater than a critical value for warping of the disc.  Fig. \\ref{figure:fstar} shows $F*$ as a function of distance from the compact object for each system considered and for a luminosity equal to $2.5\\%L_{Edd}$.  We used the ray-tracing algorithm to calculate the extent of the disc illuminated by the radiation source and the figure plots $F*$ out to this distance for each disc. The figure indicates that $F*$ is largest at the edge of the disc and that the disc warp starts at the edge of the disc and propagates towards the central object.  We conducted a number of simulations in order to investigate the warp behaviour as the central luminosity approached $L_{Edd}$. In our simulations whenever the central luminosity was just slightly larger than the Eddington limit the accretion disc material would start moving away from the central object and accretion at the central object would cease. The warp angles in our simulations never exceeded \\degree{90} and were about \\degree{30} - \\degree{40} at most, even when the central luminosity was close to the Eddington limit.  \\begin{figure} \\psfig{file=figure12.eps,width=0.5\\textwidth,angle=0} \\caption{ Radiation field strength evaluated using equation (\\ref{eq:f*}). The plots are for a constant luminosity of approximately $2.5\\%$ $L_{Edd}$. The vertical axis is the dimensionless radiation field strenght and the horizontal axis is distance from the primary object normalised so that the binary separation is 1. Plot (a), (b), (c), (d) and (e) are correspond to models 1, 4, 7, 10 and 13 respectively. The extent of each plot was calculated using a ray-tracing algorithm and shows the maximum disc radius that was illuminated by the central radiation source. } \\label{figure:fstar} \\end{figure}  \\subsection{Warp wave transportation process} Fig. \\ref{figure:fstar} indicates that the outer regions of the accretion disc that are illuminated by the central radiation source are most unstable to warping. The warp starts in these outer regions of the disc and propagates toward the central regions of the disc in the form of a prograde spiral. Wijers and Pringle (1999) indicate two processes for the transportation of a warp through a disc. The first process is a wave-like effect \\cite{LubowPringle:1993,KorycanskyPringle:1995,PapaloizonLin:1995,OgilvieLubow:1999} and this regime was used by Larwood $\\&$ Papaloizon (1997) for modelling warped circumbinary discs using a SPH code. The second process involves the exchange of angular momentum in a radial direction via viscous processes (Pringle 1992). Wijers and Pringle (1999) also indicate the condition for warp waves to propagate over a radial distance $R$ (their equation 30)  \\begin{equation}\\label{eq:wave} \\alpha \\ll  \\frac{H}{R}, \\end{equation}  \\noindent where $\\alpha$ is the normal Shakura-Sunyaev (1973) viscosity parameter and $H$ is the local disc thickness. In our current stimulations this condition was not satisfied and the warp transport process was dominated by viscous torques.  \\subsection{Effects of donor star magnetic field} Murray et al. (2002) modelled an initially planar disc subject to an inclined magnetic dipole field centred on the donor star. We conducted a limited number of simulations to investigate the warping effect reported and found that the warping was transient in nature. The warp developed strongly when the magnetic dipole field was first applied to the accretion disc but then dissipated after a period of time. Our number of particles was about a tenth of that used by Murray et al. (2002).  \\subsection{Future work} In a future paper, already in preparation, we shall apply our SPH code to a number of systems that have photometric periodicities that are significantly longer than their orbital periods. In particular we will investigate Her X1/HZ X-1 and SS 433. Both of these systems have long periods in their observed optical and X-ray light curves which have been attributed to an inclined accretion disc that is precessing in a retrograde direction."
    },
    "0601/astro-ph0601304_arXiv.txt": {
        "abstract": "We present a deep observation with the {\\em Chandra X-ray Observatory}\\ of the neutron star bow shock \\pwn\\ in the supernova remnant (SNR) \\snr.  Our data confirm the cometary morphology and central point source seen previously, but also reveal considerable new structure. Specifically, we find that the X-ray nebula consists of two distinct components: a ``tongue'' of bright emission close to the neutron star, enveloped by a larger, fainter ``tail''. We interpret the tongue and tail as delineating the termination shock and the post-shock flow, respectively, as previously identified also in the pulsar bow shock \\mouse\\ (``the Mouse''). However, for \\pwn\\ the tongue is much less elongated than for the Mouse, while the tail is much broader. These differences are consistent with the low Mach number, $\\mathcal{M} \\la 2$, expected for a neutron star moving through the hot gas in a SNR's interior, supporting the case for a physical association between \\pwn\\ and \\snr. We resolve the stand-off distance between the star and the head of the bow shock, which allows us to estimate a space velocity for the neutron star of $\\sim230$~\\kms, independent of distance.  We detect thermal emission from the neutron star surface at a temperature of $102\\pm22$~eV, which is consistent with the age of SNR~\\snr\\ for standard neutron star cooling models. We also identify two compact knots of hard emission located $1''-2''$ north and south of the neutron star. ",
        "introduction": "\\label{sec_intro}   The relativistic winds from energetic pulsars generate broadband synchrotron emission when they interact with their surroundings. New high resolution X-ray observations have revealed amazing details in the structure of the resulting pulsar wind nebulae (PWNe) \\citep{wh06,gs06}. Most of these systems resemble the Crab Nebula, and typify early phases of PWN evolution. However, after a few thousand years a pulsar's motion becomes supersonic with respect to its surroundings, and a bow shock is formed. While only a small group of pulsar bow shocks have been as yet identified, these sources provide an additional opportunity to probe pulsar winds, using very different boundary conditions from those experienced by ``Crab-like'' PWNe.  \\cite{gvc+04}, hereafter G2004, presented a detailed study of \\mouse\\ (``the Mouse''), a spectacular radio and X-ray bow shock powered by the energetic pulsar \\mousepsr. This analysis led to the identification of distinctive features which should also be observable in other bow shocks. A comparable system is \\pwn\\ (also known as 1SAX~J0617.1+2221), an apparent X-ray and radio bow shock coincident with the supernova remnant (SNR) IC~443, as shown in Figure~\\ref{fig_wide} (see also Olbert \\etal\\ 2001\\nocite{ocw+01}, hereafter O2001; Bocchino \\& Bykov 2001\\nocite{bb01b}).  Pulsations from a pulsar in \\pwn\\ have not yet been detected, but the nebula's cometary morphology, with a point source near the apex, make it virtually certain that this is a pulsar-driven bow shock. Here we present a deep \\cxo\\ observation of \\pwn, with which we study the structure of this source and compare it to that seen for the Mouse.  Before discussing this system further, we note that it is unclear if \\pwn\\ and \\snr\\ are genuinely associated because, as can be seen in Figure~\\ref{fig_wide}, the cometary trail of \\pwn\\ points $\\sim80^\\circ$ away from the direction expected if the embedded pulsar is moving away from the SNR's center.  Despite this difficulty, for now we assume that the pulsar, PWN and SNR are physically associated, and consequently adopt a common distance to them of $1.5d_{1.5}$~kpc \\citep{ws03}.  In \\S\\ref{sec_assoc} we will discuss the validity of this assumption in the light of the new data presented here. ",
        "conclusions": "\\label{sec_conc}  Our new \\cxo\\ observations of \\pwn\\ reveal a thermally-emitting neutron star with a typical space velocity of $\\sim230$~\\kms, whose relativistic wind drives a bow shock through the shocked gas inside its associated SNR~\\snr. We place this system midway between PSR~B1853+01 / SNR~W44 (for which the pulsar bow shock is deep in the interior of the SNR) and PSR~B1951+32 / SNR~CTB~80 (for which the pulsar is now crossing the SNR shell) in the evolutionary sequence described by \\cite{vdk04}. While pulsations from the central source in \\pwn\\ have not been detected in the radio band, pulsed magnetospheric emission may be detectable in future deep X-ray observations.  \\pwn\\ shows a close correspondence with the Mouse, in that the X-ray emission from both sources consists of two distinct components, a ``tongue'' and a ``tail''. The specific differences between these features in \\pwn\\ and in the Mouse can be understood in terms of the much lower Mach number of the pulsar for \\pwn. The presence of these features in two systems, and our ability to interpret them consistently in terms of the expected wind shock structures, argue that these are ubiquitous features in pulsar bow shocks.  Deeper observations of other pulsar bow shocks \\cite[especially those with detected pulsars and known proper motions, see][]{gjs02,sgk+03,cbd+03}, in hand with new relativistic magnetohydrodynamic simulations of these systems \\cite[e.g.,][]{bad05}, can further add to our understanding of these systems."
    },
    "0601/astro-ph0601132_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\\be{\\begin{equation}} \\def\\ee{\\end{equation}} \\newcommand{\\bea}{\\begin{eqnarray}} \\newcommand{\\eea}{\\end{eqnarray}} \\newcommand{\\vx}{\\ensuremath{\\vec{x}}} \\newcommand{\\vk}{\\ensuremath{\\vec{k}}} \\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{% Clarifying Slow Roll Inflation and the Quantum Corrections to the Observable Power Spectra } \\author{% D. Boyanovsky$^{c,a,}$\\footnote{E-mail:boyan@pitt.edu}, \\underline{H. J. de Vega}$^{b,a,}$\\footnote{E-mail:devega@lpthe.jussieu.fr}, N. G. Sanchez $^{a,}$\\footnote{E-mail:Norma.Sanchez@obspm.fr} } \\address{% $^{a}$ Observatoire de Paris, LERMA, Laboratoire Associ\\'e au CNRS UMR 8112, \\\\ 61, Avenue de l'Observatoire, 75014 Paris, France. % \\\\$^{b}$ LPTHE, Laboratoire Associ\\'e au CNRS UMR 7589, Universit\\'e Pierre et Marie Curie (Paris VI) \\\\ et Denis Diderot (Paris VII), Tour 24, 5 \\`eme. \\'etage, 4, Place Jussieu, 75252 Paris, Cedex 05, France. \\\\ $^{c}$Department of Physics and Astronomy, University of Pittsburgh, Pittsburgh, Pennsylvania 15260, USA } \\abstract{Slow-roll inflation can be studied as an effective field theory. The form of the inflaton potential consistent with the data is $ V(\\phi) = N \\; M^4 \\; w\\left(\\frac{\\phi}{\\sqrt{N} \\; M_{Pl}}\\right) $ where $ \\phi $ is the inflaton field, $ M $ is the inflation energy scale, and $ N \\sim 50 $ is the number of efolds since the cosmologically relevant modes exited the Hubble radius until the end of inflation. The dimensionless function $ w(\\chi) $ and field $ \\chi $ are generically $ \\mathcal{O}(1) $. The WMAP value for the amplitude of scalar adiabatic fluctuations yields $ M \\sim 0.77\\times 10^{16}$GeV. This form of the potential encodes the slow-roll expansion as an expansion in $1/N$.  A Ginzburg-Landau (polynomial) realization of $ w(\\chi) $ reveals that the Hubble parameter, inflaton mass and non-linear couplings are of the see-saw form in terms of the small ratio $ M/M_{Pl} $. The quartic coupling is $ \\lambda \\sim \\frac{1}{N} \\left(\\frac{M}{M_{Pl}}\\right)^4 $. The smallness of the non-linear couplings is {\\bf not} a result of fine tuning but a {\\bf natural} consequence of the validity of the effective field theory and slow roll approximation. Our observations suggest that slow-roll inflation may well be described by an almost critical theory, near an infrared stable gaussian fixed point. Quantum corrections to slow roll inflation are computed and turn to be an expansion in  powers $ \\left(H/M_{Pl}\\right)^2 $. The corrections to the inflaton effective potential and its equation of motion are computed, as well as the quantum corrections to the observable power spectra. The near scale invariance of the fluctuations introduces a strong infrared behavior naturally regularized by the slow roll parameter $ \\Delta \\equiv \\eta_V-\\epsilon_V = \\frac12(n_s -1)+ r/8 $. We find the \\emph{effective} inflaton potential during slow roll inflation including the contributions from scalar curvature and tensor perturbations as well as from light scalars and Dirac fermions coupled to the inflaton. The scalar and tensor superhorizon contributions feature infrared enhancements regulated by slow roll parameters. Fermions and gravitons do not exhibit infrared enhancement. The subhorizon part is completely specified by the trace anomaly of the fields with different spins and is solely determined by the space-time geometry. This inflationary effective potential is strikingly {\\bf different} from the usual Minkowski space-time result. Quantum corrections to the power spectra are expressed in terms of the CMB observables: $ n_s, \\; r $ and $ dn_s/d \\ln k $. Trace anomalies (especially the graviton part) dominate these quantum corrections in a definite direction: they {\\bf enhance} the scalar curvature fluctuations and {\\bf reduce} the tensor fluctuations. } ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601242_arXiv.txt": {
        "abstract": "We report here on variable propagation effects in over twenty years of multi-frequency timing analysis of pulsar PSR B1937+21 that determine small-scale properties of the intervening plasma as it drifts through the sight line. The phase structure function derived from the dispersion measure variations is in remarkable agreement with that expected from the Kolmogorov spectrum, with a power law index of $3.66\\pm 0.04$, valid over an inferred scale range of 0.2---50 A.U. The observed flux variation time scale and the modulation index, along with their frequency dependence, are discrepant with the values expected from a Kolmogorov spectrum with infinitismally small inner scale cutoff, suggesting a caustic-dominated regime of interstellar optics. This implies an inner scale cutoff to the spectrum of $\\sim 1.3\\times 10^9$ meters. Our timing solutions indicate a transverse velocity of 9 km sec$^{-1}$ with respect to the solar system barycenter, and 80 km sec$^{-1}$ with respect to the pulsar's LSR. We interpret the frequency dependent variations of DM as a result of the apparent angular broadening of the source, which is a sensitive function of frequency ($\\propto\\nu^{-2.2}$). The error introduced by this in timing this pulsar is $\\sim$2.2 $\\mu$s at 1 GHz. The timing error introduced by ``image wandering'' from the slow, nominally refractive scintillation effects is about 125 nanosec at 1 GHz. The error accumulated due to positional error (due to image wandering) in solar system barycentric corrections is about 85 nanosec at 1 GHz. ",
        "introduction": "The dispersion measure (DM) of a pulsar probes the column density of free electrons along the line of sight (LOS). Observed DM variations over time scales of several weeks to years sample structures in the electron plasma over length scales of $10^{10}$ m to $10^{12}$ m. Diffraction of pulsar signals is the result of scattering by structures on scales below the Fresnel radius, $10^{8}$ m or so. The DM as well as the scattering measure (SM) variability along the LOS to the Crab pulsar was first reported by Rankin \\& Isaacman (1977), who reported that the DM variability poorly correlated with the SM variability. Helfand et al. (1980) inferred an upper limit for DM variations of a few parts in a thousand for several pulsars.  In an earlier study of PSR B1937+21 Cordes et al. (1990) measured a DM change of $\\Delta DM\\sim 0.003$ pc cm$^{-3}$ over a period of a thousand days. The work of Phillips \\& Wolszczan (1991) presented the variations of DM observed along the LOS to a few pulsars. They connected these variations to those on diffractive scales, and derived an electron density fluctuation spectrum slope of $3.85\\pm0.04$ over a scale range of $10^7-10^{13}$ meters. Backer et al. (1993) report on further DM variability and show that the amplitude of the variations known at that time are consistent with a scaling by the square root of DM.  Another important investigation by Kaspi et al. (1994) studied DM variations of the millisecond pulsars PSR B1937+21 and B1855+09 over a time interval of calendar years 1984 to 1993. In addition to establishing a secular variation in DM over this time interval, they also show that the underlying density power spectrum has an index of 3.874$\\pm$0.011, which is close to what we would expect if the density fluctuations are described by Kolmogorov spectrum. An anomalous dispersion event towards the Crab pulsar was reported by Backer et al. (2000), where they report a DM ``jump'' as large as 0.1 pc cm$^{-3}$.  \\begin{figure} \\begin{center} \\epsfig{file=fig1.ps,height=8cm,angle=-90} \\caption[]{Summary of our data sample. See text for details.} \\label{fig:obs_summary} \\end{center} \\end{figure}  In this work, we present results of various long term monitoring programs on PSR B1937+21. Our data, which includes that of Kaspi, Taylor \\& Ryba (1994), spans calendar years from 1983 to 2004. These data sets have been taken with five different telescopes, the NRAO\\footnote{The National Radio Astronomy Observatory (NRAO) is owned and operated by Associated Universities, Inc under contract with the US National Science Foundation.} Green Bank 42-m (140-ft) and 26-m (85-ft) telescopes, the NAIC\\footnote{The National Astronomy and Ionosphere Center is operated by Cornell University under contract with the US National Science Foundation.} Arecibo telescope and the \\nancay ~telescope, at frequency bands of 327, 610, 800, 1400 and 2200 MHz. After giving the details of our observations in \\S\\ref{sec-obsvn}, we describe our analysis methods in \\S\\ref{sec-anal}. This is followed in \\S\\ref{sec-sightline} by a discussion on the distribution of scattering material along the LOS. As we describe, the knowledge of temporal and angular broadening of the source, proper motion, and scintillation based velocity estimates enables us to at least qualitatively study the distribution of scattering matter as well as properties of its wave-number spectrum.  We have measured some of the basic refractive scintillation parameters from our observations, and these are discussed in \\S\\ref{sec-diffref}. The frequency dependence of the refractive scintillation time scale and the modulation index indicate a caustic-dominated regime that results from a large inner scale in the spectrum.  We have detected DM variations as a function of time and frequency. We determine the phase structure function of the medium with the knowledge of the time dependent DM variations, which is consistent with a Kolmogorov distribution of density fluctuations between scale sizes of about 1 to 100 A.U. These are summarized in \\S\\ref{sec-dmvar} and \\S\\ref{sec-dmfreq}.  \\begin{table} \\begin{center} \\caption[]{Template fit parameters at various frequencies.} \\begin{tabular}{lcccccccc} \\hline {\\bf $\\nu$ (MHz)} & backend & ~~$w_1$~~ & ~~$l_2$~~ & ~~$w_2$~~ & ~~$h_2$~~ & ~~$l_3$~~ & ~~$w_3$~~ & ~~$h_3$~~ \\\\ \\hline\\hline 327\t & GBPP\t& 8.7 & 0\t& 0 & 0 & 186.9  & 12.0 & 0.51 \\\\ 430\t & ABPP\t& 10.3\t& 7.0\t& 2.5\t& 0.19\t& 186.8\t& 12.4\t& 0.53 \\\\ 610\t & GBPP\t& 9.4\t& 7.1\t& 2.9\t& 0.34\t& 187.3\t& 10.8\t& 0.60 \\\\ 863\t & EBPP\t& 8.9\t& 7.7\t& 5.2\t& 0.38\t& 187.7\t& 10.9\t& 0.55 \\\\ 1000 & GBPP & 9.2   & 8.6   & 3.9   & 0.56  & 187.9 & 11.3  & 0.54 \\\\ 1419 & ABPP\t& 8.5\t& 8.5\t& 3.7\t& 0.37\t& 187.5\t& 10.1\t& 0.45 \\\\ 1689 & EBPP\t& 9.1\t& 9.2\t& 3.9\t& 0.37\t& 187.4\t& 10.5\t& 0.40 \\\\ 2200 & GBPP\t& 9.4\t& 9.8\t& 3.2\t& 0.29\t& 187.6\t& 10.4\t& 0.36 \\\\ 2379 & ABPP\t& 10.0\t& 9.8\t& 3.0\t& 0.29\t& 187.1\t& 10.8\t& 0.33 \\\\ \\hline \\end{tabular} \\label{tab:fitpars} \\end{center} \\end{table}  PSR B1937+21 is known for its short term timing stability. However, the achievable long term timing accuracy is suspected to be seriously limited by the interstellar scattering properties. With our sensitive measurements, we are in a position to quantify these errors. In \\S\\ref{sec-timingerror}, we describe in detail various sources of these errors and quantify them. ",
        "conclusions": "We have presented in this paper a summary of over twenty years of timing of PSR B1937+21. These observations have been done over frequencies ranging from 327 MHz to 2.2 GHz with three different telescopes.  Given the agreement between the measured apparent angular broadening and that estimated by the temporal broadening, and the measured proper motion velocity and that estimated by the knowledge of scintillation parameters, we conclude that the scattering material is uniformly distributed along the sightline.  There are three significant discrepancies between the expected values and the measured refractive parameters. These are, \\begin{enumerate} \\item The measured flux variation time scale is about an order of magnitude shorter than what is expected from the knowledge of the observed apparent angular broadening. \\item The flux variation time scale is observed to be directly proportional to the wavelength, whereas it is expected to vary as proportional to $\\lambda^{2.2}$ (for a Kolmogorov spectrum). \\item The flux modulation index is observed to have a wavelength dependence that is much ``shallower\" than the expected value. \\end{enumerate} These three discrepancies consistently imply that the optics is ``caustic-dominated''. This would mean that the density irregularity spectrum has a large inner scale cutoff, $1.3\\times 10^9$ m.  Our extrapolation of the phase structure function from the regime sampled by DM variations to the diffractive regime seems to indicate that the expected $\\tdiff$ value is considerably shorter than the measured value. This is in favor of the above conclusion. Accurate measurements of frequency dependence of diffractive parameters is much needed.  In general, Millisecond pulsars are known for their timing stability. Potentially, we may achieve adequate accuracy in timing some of these pulsars to understand some of the most important questions related to the gravitational background radiation, or the internal structure of these neutron stars. However, our analysis here shows that interstellar scattering could be an important and significant source of timing error.  As we have shown, although PSR B1937+21 is known to produce short term TOA errors as low as 10--20 nanosec with sensitive observations, the long term error is far larger than this. After fitting for $\\ddot{f}$ (which absorbes most of the achromatic timing noise), the best accuracy that we can achieve for this pulsar is $0.9\\;\\mu$sec at 1.4 GHz, and about 0.5 $\\mu$sec at 2.2 GHz (by one of the authors, Andrea Lommen). It appears that almost all of this error can be accounted for by various effects that we have discussed in \\S\\ref{sec-timingerror}. In general, for millisecond pulsars with substantial DM, even if achromatic timing noise is small, interstellar medium may be a major source of timing noise."
    },
    "0601/astro-ph0601597_arXiv.txt": {
        "abstract": "{We reanalyse Cosmic Microwave Background data from experiments probing both large and small scales. We assume that measured anisotropies are due not only to primary fluctuations but also, especially at small scales, to secondary effects (namely the Sunyaev-Zel'dovich effect) and possible point source contaminations. We first consider primary and secondary anisotropies only.  For the first time in such analyses, the cosmological dependence of secondary fluctuations is fully taken into account. We show in that case that a higher value of the normalisation $\\sigma_8$ is preferred, as found by previous studies, but also higher values of the optical depth $\\tau$ and power spectrum index $n_s$ are needed. In the second part of our analysis, we further include possible contaminations from unresolved and unremoved point sources. Under these considerations, we discuss the effects on the cosmological parameters. We further obtain the best combination of relative contributions of the three kinds of sources to the measured microwave power on small scales at each frequency. Our method allows us to simultaneously obtain cosmological parameters and explain the so-called small scale power excess in a consistent way. ",
        "introduction": "Measurements of the Cosmic Microwave Background (CMB) angular power spectrum are now becoming invaluable observables for cosmology. The detailed shape of this spectrum allows one to determine cosmological parameters with high precision. The CMB anisotropy spectrum has been recently measured from large to intermediate scales with WMAP \\citep[][]{bennett2003}, Archeops \\citep[][]{benoit2003} and BOOMERanG \\citep[][]{B03} experiments. At small scales, different experiments, ACBAR \\citep[][]{acbar}, CBI \\citep[][]{cbi, cbi2}, VSA \\citep[][]{rebolo2004} and BIMA \\citep[][]{bima}, have detected signal in excess of  that expected from purely primordial CMB anisotropies. This excess is not fully understood yet but several explanations have been proposed. Some relate to the early Universe \\citep[e.g.][]{cooray2002,griffiths2003}, some are associated with reionising sources \\citep[e.g.][]{oh2003} and others are related to the astrophysical processes contributing to the signal at the CMB frequencies \\citep[e.g.][]{bond2005,toffolatti2005}.  As a matter of fact, when observing at microwave wavelengths one expects to detect not only primordial CMB fluctuations but also secondary anisotropies as well as galactic and extra-galactic foregrounds. Astrophysical contributions are, as much as possible, removed during data analysis to produce clean estimates of the CMB angular power spectrum. Nevertheless, residuals may remain and contribute to the measured power.  In order to minimise galactic contamination, most of the small sky coverage experiments (probing small scales) observe at high galactic latitudes in regions where the sky is as much as possible free of galactic emissions. Moreover, the galactic signal can be removed through multi-frequency observations, or even by purely masking the contaminated regions. Since most of the residual signal is expected on large scales, galactic emissions only very moderately affect small scale CMB observations.  At small angular scales, additional contribution to the primary CMB signal arises from the interaction of CMB photons with matter along their propagation path. These are the so-called secondary anisotropies \\citep[e.g.][ and references therein]{hu2002}. The latter are dominated by the thermal Sunyaev-Zel'dovich (SZ) effect \\citep[][]{sz72,sz80}, i.e. CMB photon inverse Compton scattering off free electrons in the hot intra-cluster medium. When clusters are known and resolved by an instrument, they can easily be cleaned out of CMB maps. However, not all clusters are known and we expect a contribution to the microwave signal from unresolved/unknown clusters.  The SZ effect being frequency dependent (contrary to primary CMB fluctuations) and its spectral signature being known, one can use this information to detect and remove SZ signal from unknown clusters in a map if multi-frequency observations are conducted. On the contrary, with single frequency experiments, we are consequently left with an SZ contamination at small scales.   Contamination by extra-galactic sources has long been studied for CMB and SZ observations. The effects of radio sources were investigated for low frequency experiments \\citep[e.g.][]{ledlow96,cooray98,toffolatti98,sokasian2001,dezotti2005,gonzalez2005}. At high frequencies dusty galaxies emitting in the infra-red are the dominant sources of contamination \\citep[eg][]{blain98,white2003,lagache2005}.  Contamination by point sources can be monitored by interferometric arrays, and multi-frequency observations should help us in reducing the additional signal through component separation.  Once again, however, a residual contamination is likely to remain, especially in view of the uncertainties in the modelling of the sources themselves.  The excess power reported by small scale CMB experiments may thus be due to several astrophysical signals expected to contribute at those scales. In order to correctly estimate the relative and possible contribution of the secondary effects or extra-galactic signal on top of the primary CMB anisotropy, one needs a coherent analysis. Cosmological dependences are expected for CMB and SZ effect, and SZ and point sources are contributing on rather similar scales. Thus, unlike previous studies \\citep[e.g.][]{goldstein2003,bond2005}, we do not decorrelate small scales from those on large scales in our analysis of the CMB data. On the contrary, we look for the best model that fits data simultaneously on large and small scales. Hence, assuming that primary, secondary and point sources signals are all contributing to the model, we present a coherent analysis (i.e. including the full cosmological dependence of secondary effects) of the CMB angular power spectrum as measured from large ($\\ell=2$) to small ($\\ell=10000$) scales.   In the next two sections we describe the CMB data we used and we present our modelling of the power spectra of the different contributions that enter the analysis. We then present our results in terms of cosmological parameters inferred from an analysis accounting coherently for primary, secondary and extra-galactic contributions. We show in each case the relative contributions of all considered signals, and discuss our results. ",
        "conclusions": "We presented the first coherent analysis of CMB data available from $\\ell =2$ to $\\ell = 10000$, taking simultaneously into account contributions from the primary fluctuations and secondary anisotropies as well as point source contamination. More specifically, we have explicitely taken into account in a consistent way the cosmological dependence of secondary anisotropies. In the first part of our work, this allowed us to confirm that SZ effect contribution to the CMB power spectrum could explain the so-called small scale power excess, as found in previous studies. We showed however that this requires not only high values of $\\sigma_8$, but also high values of the spectral index $n_s$ and of the Thomson scattering optical depth $\\tau$, leaving other cosmological parameters basically unaffected.  In the second part of our work, we included also unresolved/unremoved point source contamination. In that case, we recover the cosmological parameter values as determined from large scale CMB data only.  Our consistent analysis introduces a parametrisation of the point source signal. Consequently, we obtain estimates of the contribution of each CMB source to the total power spectrum. Our method hence allows us to test for cosmological consistency of physical models of point source populations."
    },
    "0601/astro-ph0601074_arXiv.txt": {
        "abstract": "{We present 8--13 $\\mu$m spectra of eight red giants in the globular cluster 47\\,Tucanae (NGC\\,104), obtained at the European Southern Observatory 3.6m telescope. These are the first mid-infrared spectra of metal-poor, low-mass stars. The spectrum of at least one of these, namely the extremely red, large-amplitude variable V1, shows direct evidence of circumstellar grains made of amorphous silicate. ",
        "introduction": "Galactic globular clusters offer us an opportunity to study the late stages of evolution of stars with a main sequence mass $\\sim0.8$ M$_\\odot$ at an accurately known distance. These clusters span a range in metallicity from [Fe/H]$<-2$ to $\\sim0$, allowing studies of the dependence on metallicity over more than two orders of magnitude. As the most populous star clusters in the Galaxy, they also harbour rare examples of short-lived phases in stellar evolution.  Of particular interest for the evolution of low-mass stars and their r\\^{o}le in galactic evolution is the loss of $\\sim30$ per cent of their mass during post-Main Sequence evolution. Some of this occurs prior to the core helium burning stage, near the tip of the first ascent Red Giant Branch (RGB), and some on the Asymptotic Giant Branch (AGB). Alfv\\'{e}n waves generated in chromospherically active giants may drive mass loss at rates of $\\dot{M}\\sim10^{-7}$ M$_\\odot$ yr$^{-1}$ (Pijpers \\& Habing 1989; Judge \\& Stencel 1991; Schr\\\"{o}der \\& Cuntz 2005), but in the coolest giants a combination of strong radial pulsations with circumstellar dust formation may facilitate a radiation-driven wind with $\\dot{M}\\sim10^{-6}$ M$_\\odot$ yr$^{-1}$ or more (Gehrz \\& Woolf 1971; Bowen \\& Willson 1991; Winters et al.\\ 2000). Very little is known about the dependence of dust-driven winds on metallicity (van Loon 2006).  Infrared (IR) emission from circumstellar dust has been detected in several globular clusters (Ramdani \\& Jorissen 2001; Origlia et al.\\ 2002), most notably in 47\\,Tucanae --- a massive cluster at 5 kpc (Gratton et al.\\ 2003), with [Fe/H]=$-0.66$ (Carretta \\& Gratton 1997). We here present the first mid-IR spectra of globular cluster red giants, obtained to identify the nature of the dust grains. ",
        "conclusions": "We presented the first mid-infrared spectra of red giants in a galactic globular cluster, 47\\,Tucanae. The most evolved object, 47\\,Tuc\\,V1 is found to be surrounded by dust grains made of amorphous silicate, and to experience mass loss at a rate of $\\dot{M}=1.0\\times10^{-6}$ M$_\\odot$ yr$^{-1}$."
    },
    "0601/astro-ph0601238_arXiv.txt": {
        "abstract": "{Among unidentified gamma-ray sources in the galactic plane, there are some that present significant variability and have been proposed to be high-mass microquasars. To deepen the study of the possible association between variable low galactic latitude gamma-ray sources and microquasars, we have applied a leptonic jet model based on the microquasar scenario that reproduces the gamma-ray spectrum of three unidentified gamma-ray sources, \\7, \\8\\ and \\cc, and is consistent with the observational constraints at lower energies. We conclude that if these sources were generated by microquasars, the particle acceleration processes could not be as efficient as in other objects of this type that present harder gamma-ray spectra. Moreover, the dominant mechanism of high-energy emission should be synchrotron self-Compton (SSC) scattering, and the radio jets may only be observed at low frequencies. For each particular case, further predictions of jet physical conditions and variability generation mechanisms have been made in the context of the model. Although there might be other candidates able to explain the emission coming from these sources, microquasars cannot be excluded as counterparts. Observations performed by the next generation of gamma-ray instruments, like GLAST, are required to test the proposed model. ",
        "introduction": "\\label{intro}  The instruments EGRET\\footnote{http//cossc.gsfc.gov/egret/} and COMPTEL\\footnote{http//cossc.gsfc.gov/comptel/}, onboard the Compton Gamma Ray Observatory (CGRO),  detected about two hundred gamma-ray sources that still remain unidentified.  Among these sources, there is a subgroup that appears to be concentrated towards the galactic plane and presents significant variability (Torres et~al. \\cite{Torres01}, Nolan et~al. \\cite{Nolan03}). The discovery of the microquasar LS~5039, a high-mass X-ray binary (XRB) with  relativistic jets, and its association with the high-energy gamma-ray source 3EG~J1824$-$1514 (Paredes et~al. \\cite{Paredes00}), opened the possibility that some other unidentified EGRET sources (Hartman et~al. \\cite{Hartman99}) could also be microquasars. That microquasars can be high-energy  gamma-ray emitters has been confirmed by the ground-based Cherenkov telescope HESS, that detected a TeV source whose  very small 3-$\\sigma$ error box contains LS~5039 (Aharonian et~al. \\cite{Aharonian05}). In addition, high-mass microquasars have been proposed to be counterparts of at least a significant fraction of the low galactic latitude unidentified variable EGRET sources (e.g. Kaufman Bernad\\'o et~al. \\cite{Kaufman02}, Romero et~al. \\cite{Romero04a}). Recent statistical and theoretical studies on this group of sources have provided additional support to this association (Bosch-Ramon et~al. \\cite{Bosch05a}). Therefore, it seems at least plausible  that microquasars could represent a significant fraction of the variable gamma-ray sources in the galactic plane, generating not only the emission detected by EGRET but also that of variable sources detected  by other gamma-ray instruments like COMPTEL. This paper deepens the study  of the gamma-ray source/microquasar connection by applying a detailed microquasar model to three unidentified gamma-ray sources: \\7\\ and \\8, two likely variable unidentified EGRET sources in the galactic plane\\footnote{ \\7\\ and \\8, at galactic latitudes 9$^{\\circ}$ and 6$^{\\circ}$ respectively and assuming galactic distances, are at few hundreds of parsecs above the galactic plane. It does not preclude that they are relatively young objects provided that the EGRET microquasar LS~5039 is a runaway object that could during its lifetime reach galactic latitudes up to 10$^{\\circ}$ or vertical distances of 500~pc from the galactic plane (Rib\\'o et~al. \\cite{Ribo02}).} (Torres et~al. \\cite{Torres01}, Nolan et~al. \\cite{Nolan03}),  and \\cc, recently discovered by Zhang et~al. (\\cite{Zhang02}) in a re-analysis of the COMPTEL data, which is also both variable and located in the galactic plane. Our aim is to check whether a microquasar model \"under reasonable assumptions\" can be compatible with the observational constraints at different frequencies.  The contents of this paper are arranged as follows: in Sect.~\\ref{model}, the microquasar  model is described; in Sect.~\\ref{gamma}, the application of the model to each source as well as a brief discussion of its results and predictions are presented; the work is summarized in Sect.~\\ref{con}. ",
        "conclusions": "\\label{con}  A microquasar model is applied to model the emission at different wavelengths coming from the direction of \\7, \\8 and \\cc. In the context of this model, the gamma-ray emitting jets would radiate mainly via SSC, and would present a lower electron maximum energy than the microquasars LS~5039 and LS~I~+61~303. Due to the low electron maximum energy, the radio emission is low, and only detectable for the two EGRET sources at low frequencies if the electrons are effectively re-accelerated in the radio jet, which is expected to be quite extended due to the low radiative efficiency for the electron energy and magnetic field values in there. For the COMPTEL source, detectable radio emission is not expected. We have estimated under what conditions the variability could be produced in the context of both precession and/or eccentric orbit effects, although a scenario where both effects are present seems likely."
    },
    "0601/astro-ph0601524_arXiv.txt": {
        "abstract": "Discussion of  the designation of  multiple-star components leads  to a conclusion that,  apart from components, we need  to designate systems and centers-of-mass.  The  hierarchy is coded then by  simple links to parent. This  system is  adopted in the  multiple star  catalogue, now available  on-line.  A  short  review of  multiple-star statistics  is given: the frequency of different multiplicities in the field, periods of  spectroscopic sub-systems,  relative orbit  orientation, empirical stability  criterion,  and  period-period  diagram with  its  possible connection to formation of multiple stars. ",
        "introduction": "\\label{sec:1}  Actual  designations  of binary  and  multiple  stars are  historical, non-systematic and confusing. Future space missions like GAIA, ongoing searches for  planets, and other large projects  will greatly increase the number of known multiple stars and components, making it urgent to develop  a coherent and  unambiguous designation  scheme.  Recognizing this need, IAU  formed a special group that held  its meetings in 2000 and 2003.  As  a result, IAU adopted the  designation system developed for  the  Washington  Multiple  Star Catalogue  (WMC)  \\cite{WMC},  an extension of  the current designations  in the Washington  Double Star Catalogue (WDS) \\cite{WDS}.  Old  WDS designations will not change (to avoid  further confusion),  new  component names  will be  constructed hierarchically as  sequences of letters  and numbers, e.g.   Ab2.  The designations reflect  hierarchical structure of  multiple systems, but only to a  certain extent.  Future discoveries of  components at outer or  intermediate  levels of  hierarchy  and  the  constraint of  fixed designations  will inevitably lead  to situations  where names  do not reflect the true hierarchy, which has then to be coded separately.   \\begin{figure}[ht] \\centering \\includegraphics[width=10cm]{tokovF1.eps} \\caption{Hierarchical structure  of a quintuple system  and its coding by  reference  to  parent.  Given the  components  designations  and references to  parent (right), the  hierarchical tree (left)  can be constructed automatically. } \\label{fig:1} \\end{figure}  In the  future, many components  of multiple systems will  be resolved into  sub-systems, so  that  A will  become  Aa and  Ab, for  example. Nevertheless, current designations of these components (or rather {\\em super-components}, as they contain more than one star) will not become obsolete, because past and future  measurements of these entities as a whole will still  refer to the super-component names  such as ``A'' in our  example. It  turns out  that the  concept of  super-components is crucial to the  whole problem.  Once each super-component  has its own designation,  unique   within  each  multiple  system,   we  can  code hierarchies by simple reference to parent (Fig.~\\ref{fig:1}).  The  WMC   designation  system  extended  to   normal  components  and super-components   is  thus  logical   and  flexible,   permitting  to accommodate new discoveries. Its  application is not free of ambiguity, however, so  that a  common center (or  clearing house) will  still be required. The WMC itself will hopefully play this role. ",
        "conclusions": "Cataloguing of multiple systems,  however boring it might seem, offers interesting  insights  into formation  and  evolution  of stars.   New powerful   observing  techniques  (adaptive   optics,  interferometry, precise  radial velocities)  should now  be applied  to  large stellar samples in order  to ``fix'' the multiplicity statistics  in the solar neighborhood and beyond."
    },
    "0601/astro-ph0601654_arXiv.txt": {
        "abstract": "\\begin{list}{ } {\\rightmargin 0.4in} \\baselineskip = 11pt \\parindent=1pc {\\small The Edgeworth-Kuiper belt encodes the dynamical history of the outer solar system. Kuiper belt objects (KBOs) bear witness to coagulation physics, the evolution of planetary orbits, and external perturbations from the solar neighborhood. We critically review the present-day belt's observed properties and the theories designed to explain them. Theories are organized according to a possible time-line of events. In chronological order, epochs described include (1) coagulation of KBOs in a dynamically cold disk, (2) formation of binary KBOs by fragmentary collisions and gravitational captures, (3) stirring of KBOs by Neptune-mass planets (``oligarchs''), (4) eviction of excess oligarchs, (5) continued stirring of KBOs by remaining planets whose orbits circularize by dynamical friction, (6) planetary migration and capture of Resonant KBOs, (7) creation of the inner Oort cloud by passing stars in an open stellar cluster, and (8) collisional comminution of the smallest KBOs. Recent work underscores how small, collisional, primordial planetesimals having low velocity dispersion permit the rapid assembly of $\\sim$5 Neptune-mass oligarchs at distances of 15--25 AU. We explore the consequences of such a picture. We propose that Neptune-mass planets whose orbits cross into the Kuiper belt for up to $\\sim$20 Myr help generate the high-perihelion members of the hot Classical disk and Scattered belt. By contrast, raising perihelia by sweeping secular resonances during Neptune's migration might fill these reservoirs too inefficiently when account is made of how little primordial mass might reside in bodies having sizes on the order of 100 km. These and other frontier issues in trans-Neptunian space are discussed quantitatively. \\\\~\\\\~\\\\~\\\\}%   \\end{list} ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601181_arXiv.txt": {
        "abstract": "Gravitational lensing by massive galaxy clusters is a powerful tool for the discovery and study of high redshift galaxies, including those at $z{\\ge}6$ likely responsible for cosmic re--ionization.  Pell\\'o et al.\\ recently used this technique to discover a candidate gravitationally magnified galaxy at $z{=}10$ behind the massive cluster lens Abell~1835 ($z{=}0.25$).  We present new Keck (LRIS) and \\emph{Spitzer Space Telescope} (IRAC) observations of the $z{=}10$ candidate (hereafter \\#1916) together with a re-analysis of archival optical and near-infrared imaging from the \\emph{Hubble Space Telescope} and VLT respectively.  Our analysis therefore extends from the atmospheric cut-off at ${\\lambda}_{\\rm obs}{\\simeq}0.35{\\mu}{\\rm m}$ out to ${\\lambda}_{\\rm obs}{\\simeq}5{\\mu}{\\rm m}$ with \\emph{Spitzer}/IRAC.  The $z{=}10$ galaxy is not detected in any of these data, including an independent reduction of Pell\\'o et al.'s discovery $H$- and $K$-band imaging.  We conclude that there is no statistically reliable evidence for the existence of \\#1916.  We also assess the implications of our results for ground-based near-infrared searches for gravitationally magnified galaxies at $z{\\gs}7$.  The broad conclusion is that such experiments remain feasible, assuming that space-based optical and mid-infrared imaging are available to break the degeneracy with low redshift interlopers (e.g.\\ $z{\\sim}2{-}3$) when fitting spectral templates to the photometric data. ",
        "introduction": "Observations of distant QSOs (Becker et al.\\ 2001; Fan et al.\\ 2002) and the cosmic microwave background (Kogut et al.\\ 2003) together suggest that the universe was re-ionized somewhere between $z{\\simeq}6$ and $z{\\simeq}20$.  Searching for the sources of re-ionizing photons is currently an intense observational effort. Most searches naturally concentrate on luminous systems, i.e.\\ QSOs and luminous galaxies, at $z{\\sim}6{-}8$ as these should be easier to detect than less luminous and more distant objects.  However QSOs likely produced insufficient photons to accomplish re-ionization alone (Fan et al.\\ 2001; Barger et al.\\ 2003), and the same may be true of luminous ($L{\\gs}0.3L^{\\star}_{z{=}3.8}$) galaxies based on small samples from the \\emph{Hubble Space Telescope} Ultra Deep Field (hereafter \\emph{HST} UDF; Bouwens et al.\\ 2004; Bunker et al.\\ 2004; Yan et al.\\ 2004).  This raises the important possibilities that re-ionization either occurred much earlier, or that the bulk of the re-ionizing photons were emitted by sub-luminous galaxies, i.e.\\ $L{\\ls}0.1L^{\\star}$.  The UDF studies operate close to the detection threshold of the deepest optical/near-infrared imaging available.  It is therefore difficult to envisage substantial progress in the detection of more remote and/or less luminous galaxies via deep imaging of ``blank fields'' with the current generation of telescopes.  With the advent of the \\emph{James Webb Space Telescope} still some years ahead, the magnifying power of massive galaxy cluster lenses is therefore a much needed boost for the discovery power of \\emph{HST} and large ground-based telescopes.  Indeed, the galaxy redshift record has been broken on several occasions with the help of the gravitational magnification of distant galaxies by foreground galaxy clusters (Mellier et al.\\ 1991; Franx et al.\\ 1997; Hu et al.\\ 2002; Kneib et al.\\ 2004).  The faint end of the luminosity function of Lyman-${\\alpha}$ emitters at $z{=}5$ has also been constrained with the help of gravitational lensing (Santos et al.\\ 2004; Ellis et al.\\ 2001).  Extension of these techniques to $z{\\gs}7$ is therefore an important element of observational studies of cosmic re-ionization .  Pell\\'o et al.\\ (2004 -- hereafter P04) reported a gravitationally magnified (${\\mu}{\\sim}25{-}100$) galaxy at $z{=}10$ (hereafter \\#1916, following P04's nomenclature) behind the foreground galaxy cluster A\\,1835 ($z{=}0.25$).  This interpretation is based on non-detection in optical imaging from the ground ($3{\\sigma}$ limits in a $0.6''$ diameter aperture: $V{\\ge}27.4$, $R{\\ge}27.5$, $I{\\ge}26.9$) and space ($3{\\sigma}$ limit in a $0.2''$ diameter aperture: $R_{702}{\\ge}27.2$), and the shape of the continuum at ${\\lambda}_{\\rm obs}{\\ge}1{\\mu}{\\rm m}$ (using a $1.5''$ aperture: $(J{-}H){\\ge}0.6$, $(H{-}K){=}{-}0.5{\\pm}0.4$) which is reminiscent of the Lyman-break selection technique (Steidel et al.\\ 1996).  P04 corroborated the putative Lyman-break redshifted to ${\\lambda}_{\\rm obs}{\\simeq}1.3{\\mu}{\\rm m}$ with an emission line at ${\\lambda}_{\\rm obs}{=}1.3375{\\mu}{\\rm m}$ with integrated flux of $(4.1{\\pm}0.5){\\times}10^{-18}{\\rm erg\\,cm^{-2}\\,s^{-1}}$, which they interpret as Lyman-${\\alpha}$.  Lower redshift interpretations of the line ([OII] at $z{=}2.59$; [OIII] at $z{=}1.68$; H${\\alpha}$ at $z{=}1.04$) were discarded by P04 largely on the basis of the low probability of solutions at $z{\\ls}7$ when fitting synthetic spectral energy distributions (SEDs) to their photometric data.  The most likely of these lower-redshift solutions ($z{=}2.59$) was further excluded on the basis of the dust extinction required to fit the photometric data ($A_V{\\ge}2$), and the absence of doublet structure in the observed emission line.  Based solely on the photometry (optical non-detection, red $(J{-}H)$ and blue $(H{-}K)$ colors) the $z{=}10$ interpretation of \\#1916 is plausible.  However P04's preference for this solution over the lower redshift alternatives was controversial from the outset.  For example, the emission line does not have the characteristic P-Cygni profile of Lyman-${\\alpha}$, and the different photometric apertures adopted in the optical ($0.6''$ -- smaller than the ground-based seeing disc) and near-infrared ($1.5''$ -- $3{\\times}$ the seeing disk) may suppress the likelihood of lower-redshift solutions when fitting synthetic SEDs.  Bremer et al.\\ have also suggested that \\#1916 may either not exist, or be intrinsically variable, based on their non-detection in the $H$-band with NIRI on Gemini-North, $H(3{\\sigma}){>}26$, in contrast to P04's $4{\\sigma}$ detection of $H{=}25.00{\\pm}0.25$ with ISAAC on VLT.  The spectroscopic identification of \\#1916 is also in doubt.  Weatherley et al.\\ (2004) re-analyzed P04's spectroscopic data and failed to detect the emission line at ${\\lambda}_{\\rm obs}{=}1.3375{\\mu}{\\rm m}$, citing spurious positive flux arising from variable hot pixels in the ISAAC array as the likely source of the discrepancy.  A\\,1835 has been used previously as a gravitational telescope, for example targeting sub-millimeter galaxies (hereafter SMGs; e.g.\\ Smail et al.\\ 2002) and galaxies with extremely red optical/near-infrared colors (Smith et al.\\ 2002).  One of the galaxies detected in these surveys, SMMJ\\,14011${+}$0252, lies at $z{=}2.56$ and suffers an estimated extinction of $1.8{\\ls}A_V{\\ls}6.5$ (Ivison et al.\\ 2000).  Bearing in mind recent discovery of galaxy groups at $z{\\sim}2{-}3$ associated with gravitationally lensed SMGs (Kneib et al.\\ 2004; Borys et al.\\ 2004), and the strong clustering of SMGs (Blain et al.\\ 2004), \\#1916 is plausibly at a similar redshift to SMMJ\\,14011${+}$0252, and may also be obscured by dust.  Further circumstantial evidence for a lower redshift interpretation of \\#1916 comes from Richard et al.\\ (2003) who used the same spectroscopic data as presented by P04 to discover a strongly reddened star-forming galaxy at $z{=}1.68$.  This redshift coincides with the [OIII] interpretation of PO4's putative emission line at ${\\lambda}_{\\rm obs}{=}1.3375{\\mu}{\\rm m}$.  In this paper, we address three questions: (i) is \\#1916 at $z{\\sim}2{-}3$?; (ii) is \\#1916 intrinsically variable?; (iii) does \\#1916 exist?  These tests exploit new optical and mid-infrared observations using the Keck-I 10-m telescope and the \\emph{Spitzer Space Telescope} respectively, plus an independent reduction of P04's $H$- and $K$-band imaging data from VLT.  Throughout this article we assume that the emission line at ${\\lambda}_{\\rm obs}{=}1.3375{\\mu}{\\rm m}$ is a false detection (Weatherley et al.\\ 2004).  In \\S\\ref{data} we present the new Keck and \\emph{Spitzer} data, explain in detail the re-reduction of the archival VLT/ISAAC data, and summarize the archival \\emph{Hubble Space Telescope (HST)} data.  Then in \\S\\ref{analysis} we describe the analysis and key results, focusing on the three questions posed above; this section closes with a summary of the current observational status of \\#1916.  Finally, we discuss the implications of our results for future ground-based near-infrared searches for galaxies at $z{\\gs}7$ (\\S\\ref{discuss}) and summarize our conclusions in (\\S\\ref{conclusions}).  We assume $H_0{=}65{\\rm km\\,s^{-1}Mpc^{-1}}$, ${\\Omega}_{\\rm M}{=}0.3$ and ${\\Omega}_{\\rm \\Lambda}{=}0.7$ throughout.  Unless otherwise stated all error bars are at $1{\\sigma}$ significance; photometric detection limits are at $3{\\sigma}$ significance; upper and lower limits on colors are based on $3{\\sigma}$ detection thresholds in the non-detection filter.  Magnitudes are stated in the AB system; conversion between the AB and Vega systems for the specific filters used in this paper are as follows:- ${\\Delta}_B{=}B_{\\rm AB}{-}B_{\\rm Vega}{=}{-}0.1$, ${\\Delta}_V{=}0.1$, ${\\Delta}_R{=}0.2$, ${\\Delta}_{F702W}{=}0.3$, ${\\Delta}_I{=}0.5$, ${\\Delta}_J{=}0.9$, ${\\Delta}_H{=}1.4$, ${\\Delta}_K{=}1.9$, ${\\Delta}_{3.6{\\mu}{\\rm m}}{=}2.8$, ${\\Delta}_{4.5{\\mu}{\\rm m}}{=}3.2$. ",
        "conclusions": "We have analyzed new and archival observations of the $z{=}10$ galaxy \\#1916 behind the foreground galaxy cluster lens A\\,1835 ($z{=}0.25$) spanning $0.35{\\le}{\\lambda}_{\\rm obs}{\\le}5{\\mu}{\\rm m}$. Statistically significant flux is not detected in any of these data, including our independent re-analysis of P04's discovery $H$- and $K$-band data.  The $3{\\sigma}$ detection thresholds are: $F({\\lambda}_{\\rm obs}{=}0.44{\\mu}{\\rm m}){\\gs}4.5{\\times}10^{-19}{\\rm erg\\,s^{-1}\\,cm^{-2}}$ per spectral resolution element, $R_{702}{\\ge}27.0$, $H{\\ge}25.0$, $K{\\ge}25.0$, $F(3.6{\\mu}{\\rm m}){\\le}0.75{\\mu}{\\rm Jy}$ and $F(4.5{\\mu}{\\rm m}){\\le}0.75{\\mu}{\\rm Jy}$, where the photometric limits are calculated in apertures with a diameter $3{\\times}$ that of the seeing disk.  Combining these results with those of Bremer et al.\\ (2004) and Weatherley et al.\\ (2004), we are therefore forced by the data to conclude that there is no statistically sound evidence for the existence of \\#1916.  We also show that inefficient bad-pixel rejection and issues relating to the calculation of the background noise can broadly account for the differences between P04's near-infrared photometry and our own using the same data.  The balance of probability is therefore that \\#1916 was a false detection in P04's analysis.  From a broader perspective, the demise of \\#1916 warns of the hazards of operating close to the detection threshold of deep ground-based near-infrared data.  The need for gravitational magnification to boost the observed flux of faint ($L{\\ls}0.1\\,L^\\star$) galaxies at $z{\\sim}6$--8 and populations of galaxies at still higher redshifts (see \\S\\ref{intro}) is undiminished by our results on \\#1916.  However an important issue is whether it is feasible to find such galaxies using ground-based facilities.  We explore this issue using {\\sc hyperz} to fit synthetic spectral templates to representative data.  Our main conclusions are that deep near-infrared imaging similar to that presented by P04 in combination with a single sensitive optical non-detection from \\emph{HST} imaging and moderate depth \\emph{Spitzer}/IRAC imaging at $3.6$ and $4.5{\\mu}{\\rm m}$ can discriminate between low (e.g.\\ $z{\\sim}2{-}3$) and high (e.g.\\ $z{\\gs}7$) redshift solutions with strong statistical significance.  In summary, ground-based near-infrared surveys of massive galaxy cluster lenses with 10-m class telescopes remain a powerful tool for the discovery of intrinsically faint galaxies ($L{\\ls}0.1L^\\star$) at $z{>}7$ that may be responsible for cosmic re-ionization.  Future surveys should combine these data with sensitive space-based optical and mid-infrared observations."
    },
    "0601/astro-ph0601462_arXiv.txt": {
        "abstract": "Although published in 1995, the Gauss-Seidel method for solving the non-LTE radiative transfer problem has deserved too little attention in the astrophysical community yet. Further tests of the performances and of the accuracy of the numerical scheme are provided. ",
        "introduction": "Fast and {\\em accurate} numerical schemes for the solution of the non--LTE radiation transfer (RT) problem still need to be pushed further, in order to be able to deal with increasingly complex models (e.g., multi-dimensional geometry, multi-level atoms, polarisation...). Hereafter, we present new numerical tests against {\\em exact} solutions of the Gauss-Seidel (GS) and SOR methods initially proposed by Trujillo Bueno \\& Fabiani Bendicho (1995), and inspired by the classical GS iterative method in numerical analysis (which can be modified into SOR methods when an overcorrection is made as compared to GS).  Very few tests, indeed, are available for the validation of any new numerical scheme dealing with the resolution of the non--LTE radiative transfer problem. The usual one comes from the computation of numerical solutions under the Eddington approximation i.e., adopting a very coarse angular quadrature such as $\\mu= \\pm 1/ \\sqrt{3}$ (also known as the ``two-stream'' approximation; see \\S 4.3.1.  in Rutten 2003). In such a case, analytical solutions -- of a ``reduced'' RT problem though -- are available for numerical solutions to be checked against. However this test may lead to an {\\em erroneous} estimation of the accuracy of the numerical scheme, as pointed out in Chevallier et al. (2003, see \\S 5. therein). ",
        "conclusions": "\\begin{figure} \\centerline{ \\begin{tabular}{cc} \\includegraphics[width=6.3cm]{paletou_fig1b.eps} & \\includegraphics[width=6.3cm]{paletou_fig1a.eps} \\end{tabular}}  \\caption{{\\em Left:} evolution of the maximum true error vs. the number of iterations (20 grid-points per decade) {\\em Right:} CPU-time for, respectively, the ALI and the GS/SOR numerical schemes vs. the number of points per decade (for a 1D finite slab of optical thickness 1000 and a collisional destruction probability $10^{-4}$).}  \\label{example} \\end{figure}  The GS/SOR numerical schemes are better implemented when one adopts the ``short characteristics'' (SC) approach for the so-called formal solution of the RT equation (see Auer \\& Paletou 1994). Chevallier et al. (2003) showed how the accuracy of an Accelerated $\\Lambda$--Iteration (with a monotonic parabolic SC formal solver) numerical code's solutions do scale with the refinement of both spatial and angular grids, by comparing the latter to very high-accuracy analytical solutions given by the ARTY code (Chevallier \\& Rutily 2004).  We computed GS/SOR solutions for a grid of 1D finite slab, two-level atom (in CRD) models that we compared to ARTY reference solutions: the rate of convergence for various numerical schemes is displayed in Fig.~1 (left). Our main conclusions are that: (a) as expected, the {\\em accuracy} of the GS/SOR code is identical to the one of the ALI-SC-based code, (b) the numerical ``overcost'' for GS iterations (due to a modified formal solver) is {\\em negligible} but, (c) that the gain in computing time with GS/SOR is {\\em very significant.} As shown in Figs.~1, for high-order quadratures, which are {\\em absolutely needed} in order to keep the accuracy of the ALI-SC method better than 1\\% (see Chevallier et al. 2003), the gain in CPU-time provided by the GS/SOR scheme can be as large as a factor of 50! And even with standard acceleration techniques for ALI or GS, GS/SOR remains the fastest algorithm (see Trujillo Bueno \\& Fabiani Bendicho 1995).  We feel that this fully justifies to consider {\\em very seriously} GS/SOR schemes for future radiative modelling codes (note also that GS/SOR have already been generalized to complex models). It is particularly important for the development of those diagnosis tools required by major projects such as GAIA or HERSCHEL, for instance."
    },
    "0601/astro-ph0601148_arXiv.txt": {
        "abstract": "We present new limits on ultra-high energy neutrino fluxes above $10^{17}$ eV based on data collected by the Radio Ice Cherenkov Experiment (RICE) at the South Pole from 1999-2005. We discuss estimation of backgrounds, calibration and data analysis algorithms (both on-line and off-line), procedures used for the dedicated neutrino search, and refinements in our Monte Carlo (MC) simulation, including recent {\\it in situ} measurements of the complex ice dielectric constant. An enlarged data set and a more detailed study of hadronic showers results in a sensitivity improvement of more than one order of magnitude compared to our previously published results. Examination of the full RICE data set yields zero acceptable neutrino candidates, resulting in 95\\% confidence-level model dependent limits on the flux $E_\\nu^2d\\phi/dE_\\nu<10^{-6} {\\rm GeV}/({\\rm cm^2s~sr})$ in the energy range $10^{17}< E_\\nu< 10^{20}$ eV. The new RICE results rule out the most intense flux model projections at 95\\% confidence level.  $^*${\\tt Contact for additional information.} ",
        "introduction": "} The motivation for producing a new map of the Universe, as viewed through neutrino `eyes' is now well-established. Ultra-high energy (``UHE''; E$>$1 PeV) neutrinos point back to sources of high energy cosmic rays, providing a direct picture of the source and the acceleration mechanism. Detection of UHE neutrino fluxes simultaneous with gamma-ray bursts (GRB's)\\cite{GRBs} could provide essential information on the nature of these extraordinarily luminous sources and help resolve the question of whether GRB's are responsible for the bulk of the UHE cosmic ray particle flux incident at Earth. In the $10^{18}-10^{20}$ eV energy regime, ``GZK'' neutrinos may distinguish between source evolution models for UHE cosmic rays\\cite{SeckelS05}. At even higher energies, ``Z-burst\" neutrino models\\cite{Weiler} have been proposed to explain the $\\ge$$10^{20}$~eV cosmic ray flux claimed by the AGASA experiment\\cite{AGASA98}; neutrinos may also identify more exotic sources such as topological defects\\cite{Berezinsky}. In the realm of particle physics, detection of UHE neutrinos from cosmological distances, if accompanied by flavor identification, may permit measurement of neutrino oscillation parameters over a wide range of $\\Delta m^2$\\cite{Francis-David-mixing-sensitivity,Pakvasa03} or observation of $\\nu_\\tau$ via ``double-bang'' signatures\\cite{doublebang}. Additionally, the angular distribution of upward-going neutrino events could be used to measure weak cross-sections at energies unreachable by man-made accelerators\\cite{sigma}. Alternately, if the high energy weak cross-sections are known, they can be used to test Earth composition models along an arbitrary cross-section, so-called `neutrino tomography'\\cite{tomography}.  Several recent projects (AMANDA\\cite{AMANDA1}, IceCube\\cite{IceCube}, NESTOR\\cite{NESTOR}, NEMO\\cite{NEMO}, Lake Baikal\\cite{BAIKAL}, ANTARES\\cite{Antares}, e.g.) have demonstrated photomultiplier-tube based detection of high energy cosmic ray muon neutrinos, by observing the optical Cherenkov cone which results from muons produced in $\\nu_\\mu$ weak current interactions. To detect neutrinos at UHE, it is more effective to exploit coherence of radio Cherenkov emissions. The amplitude of radio frequency signals emitted by neutrino-generated electromagnetic showers in the MHz-GHz regime increases nearly linearly with energy, making radio detection the most efficient scheme presently known at ultra-high energies.  RICE employs this strategy to search for neutrino interactions occurring in cold polar ice, which has exceptional transmission properties favorable for radio detection.  Here we extend our previous results \\cite{astroph-cal,astroph-results}, which were based on shorter running times and less advanced detector response modeling.  We report new limits based on additional data and further consideration of systematic errors. ",
        "conclusions": "} Using the full dataset accumulated thus far (1999-2005), we have presented upper limits on the incident neutrino flux. Despite suboptimal dense-packing of the array, RICE provides superior sensitivity in the energy regime between $10^{17}-10^{20}$ eV. Limits are considerably stronger than previously reported values, and the most intense flux model projections are ruled out at 95\\% confidence level."
    },
    "0601/astro-ph0601572_arXiv.txt": {
        "abstract": "We report extensive  VLA and ATCA observations of the two diffuse radio sources in the cluster of galaxies Abell 548b, which confirm their classification as relics.  The two relics (named A and B) show similar flux density, extent, shape, polarization and spectral index and are located at projected distances of about 430 and 500 kpc from the cluster center, on the same side of the cluster's X-ray peak. On the basis of  spectral indices of discrete radio sources embedded within the diffuse features, we have attempted to distinguish emission peaks of the diffuse sources from unrelated sources.  We have found that both relics, in particular the B-relic, show possible fine structure, when observed at high resolution. Another diffuse source (named C) is detected close in projection to the cluster center.  High-resolution images show that it contains two discrete radio sources and a diffuse component, which might be a candidate for a small relic source.  The nature and properties of the diffuse radio sources are discussed.  We conclude that they are likely related to the merger activity in the cluster. ",
        "introduction": "In recent years growing interest has been directed toward the study of diffuse radio sources in clusters of galaxies. Cluster radio halos have been found to permeate cluster centers, and diffuse radio sources classified as relics have been detected in the cluster peripheries (see e.g. review by Feretti 2003).  Extended relics, showing a steep spectrum, high polarization degree and no obvious association with any cluster galaxy, have been suggested to originate along cluster merger shocks, either by diffusive shock acceleration of electrons (En{\\ss}lin et al. 1998) or adiabatic recompression of fossil radio plasma (En{\\ss}lin \\& Br\\\"uggen 2002). Currently we know about 30 clusters showing relics, but detailed studies are available for only a few sources of this class (see e.g. Giovannini \\& Feretti 2004 for a review). In six clusters, the relics are double and located on symmetric sides with respect to the cluster center. Relics may show quite different observational structures and locations, which might indicate different physical properties. New information on relic sources is therefore of great importance.    \\begin{table*} \\caption{Observing log} \\begin{flushleft} \\begin{tabular}{llllllll} \\hline \\noalign{\\smallskip} RA ~ (2000) & DEC & Freq. &  Bandw. & R.Tel. (Config.)  &  Date  & Durat. & Figure\\\\ h ~ m ~ s & ~ $^{\\circ}$ ~ $^{\\prime}$ ~ $^{\\prime\\prime}$ &   MHz & MHz  & & & min \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} 05 45 10.0 & $-$25 50 00 & 1365/1435 & 50 & VLA (C) & 28 Apr 00 & 110 & Figs. \\ref{vlaopt}, \\ref{polar}, \\ref{radx} \\\\ 05 45 04.9 & $-$25 47 40 & 1385/1465 & 50 & VLA (BnA) & 18 May 02 & 20 & Fig. \\ref{relicS} \\\\ &             & 4835/4885 & 50 &  VLA (BnA) & 18 May 02 & 20 & Fig. \\ref{sourceA} \\\\ &             & 8435/8485 & 50 &  VLA (BnA) & 18 May 02 & 20 \\\\ 05 45 22.1 & $-$25 47 31 & 1385/1465 & 50 &  VLA (BnA) & 18 May 02 & 8 \\\\ &             & 4835/4885 & 50 &  VLA  (BnA) & 18 May 02 & 8 \\\\ 05 44 51.4 & $-$25 50 45 & 1384  & 104 &  ATCA (6A) &  17 Feb 00 & 720 &  Fig. \\ref{slee1} \\\\ &            & 1384  & 104 &  ATCA (6B) &  22 Jun 00 & 720 & Fig. \\ref{slee1} \\\\ &            & 1704  & 104 &  ATCA (6B) &  22 Jun 00 & 720 \\\\ &            & 2496  & 104 &  ATCA (6A) &  17 Feb 00 & 720 & Fig. \\ref{slee2} \\\\ \\noalign{\\smallskip} \\hline \\noalign{\\smallskip} \\end{tabular} \\end{flushleft} \\end{table*}    The presence of a relic source in A548 was suggested by Giovannini et al. (1999), using radio data of the NRAO VLA Sky Survey (NVSS, Condon et al. 1998).  The cluster A548, at an average redshift z=0.04 (Struble \\& Rood 1999), shows a rather complex structure both in the optical and in the X-ray band. Based on ROSAT data, Davis et al. (1995) pointed out the existence of three main sub-clusters. These are confirmed by optical data which indicate also significant further substructures, some of which overlap along the line of sight (Escalera et al. 1994, Den Hartog \\& Katgert 1996, Andreuzzi et al. 1998, Wegner et al. 1999, Colless et al. 2001, Nikogossian 2001, Smith et al. 2004). The combination of X-ray and optical data  indicates  that A548 is a cluster in a collapsing phase and therefore not yet dynamically relaxed.  The diffuse relic source detected by Giovannini et al. (1999) by inspection of the NVSS is located in the subcluster A548b, referred to in the literature also as A548S or A548W or A548SW, at a redshift z= 0.0424 (Den Hartog \\& Katgert 1996).  Radio data for this cluster were  published previously by Gregorini et al. (1994) and Marvel et al. (1999), but they refer to radio galaxies only and Marvel et al's image does not cover the region discussed in the present paper.  We present in this paper the results of new radio observations obtained with the Very Large Array (VLA) and the Australia Telescope Compact Array (ATCA) to study the radio emission in A548b.  The VLA observations were aimed at imaging the diffuse emission with improved sensitivity and resolution with respect to the NVSS.  The aim of the ATCA observations was to search for fine structure in the image of the radio relic similar to that seen in four relics by Slee et al. (2001), and resolve possible discrete sources. Thanks to the high surface brightness sensitivity of the present radio data, we detect for the first time in a cluster two radio relics located on the same side with respect to the cluster center. Moreover, a third possible diffuse source is detected close in projection to the cluster center. The paper is organized as follows: the radio data are presented in Sect. 2, the results from low resolution images in Sect. 3, those from high resolution images in Sect. 4. The results are discussed in Sect. 5.  For the computation of intrinsic parameters, we adopt the concordance cosmology with H$_0$=70 km s$^{-1}$ Mpc$^{-1}$, $\\Omega_m$ = 0.3, and $\\Omega_{\\Lambda}$ = 0.7. The luminosity distance D$_L$ of the subcluster under study is 188 Mpc; at this distance 1 arcsec corresponds to 0.84 kpc.   \\begin{figure*} \\centerline{ \\psfig{figure=fig1.ps,height=15cm}} \\caption{Contour radio emission in A548b at 1.4 GHz obtained with the VLA in C configuration at the resolution of  15$^{\\prime\\prime}$  $\\times$ 30$^{\\prime\\prime}$ (FWHM, RA $\\times$ DEC). Contour levels are  0.3, 0.6, 1, 2, 4, 8, 16, 32, 64 mJy/beam. The noise level is  0.09 mJy/beam. The grey-scale represents the optical R-band image taken from the DSS2. Labels A, B and C indicate diffuse radio sources. } \\label{vlaopt} \\end{figure*} ",
        "conclusions": "From the results and discussion presented in this paper we derive the following conclusions.  Extensive radio observations of the two diffuse radio sources A and B in the the cluster of galaxies Abell 548b confirm their classification as radio relics. The two relics, of similar flux density, extent, shape, polarization and spectral index, are located at projected distances of about 430 and 500 kpc from the cluster center, on the same side of the cluster's X-ray peak.  We present images at high resolution of the cluster and measure radio spectral indices of discrete radio sources embedded within the diffuse features.  We attempt to distinguish emission peaks of the diffuse sources from unrelated sources, and we find that both relics show possible fine structure, when observed at high resolution. We analyze the  high resolution images of the strong radio source located approximately midway between the two relics, and conclude that it has very likely no physical connection to the diffuse sources.  Another diffuse source C is detected close in projection to the cluster center.  High-resolution images show that it contains two discrete radio sources 7a and 7c and a diffuse component 7b. The nature of this complex is uncertain. It seems probable that the more compact components are not physically related to the more extended one, although all lie in projection within $\\sim$ 50 kpc of the X-ray centroid and the galaxy VV\\,162.  Although the spectral index of the complex is fairly normal, the amorphous morphology of the extended component, the absence of an active nucleus, and the missing flux in the high-resolution images, are in favor of a diffuse nature for this complex, which we thus consider a possible relic candidate.  We suggest that the two external relics may trace the same shock wave in the outer cluster region, whereas the possible relic close to the cluster center would indicate a much more internal shock wave. Future studies in the radio and X-ray domains will allow to detect the details of the merger process and shed more light on the connection between cluster mergers and the formation of radio relics."
    },
    "0601/astro-ph0601091_arXiv.txt": {
        "abstract": "Air fluorescence measurements of cosmic ray energy must be corrected for attenuation of the atmosphere. In this paper we show that the air-showers themselves can yield a measurement of the aerosol attenuation in terms of optical depth, time-averaged over extended periods. Although the technique lacks statistical power to make the critical hourly measurements that only specialized active instruments can achieve, we note the technique does not depend on absolute calibration of the detector hardware, and requires no additional equipment beyond the fluorescence detectors that observe the air showers. This paper describes the technique, and presents results based on analysis of 1258 air-showers observed in stereo by the High Resolution Fly's Eye over a four year span. ",
        "introduction": "Fluorescence detectors use the atmosphere calorimetrically to measure the energy deposited by extensive air-showers. Ultra-violet fluorescence emitted by particle cascades can be observed tens of kilometers away by a photosensitive detector when the primary cosmic particle is above 10$^{18}$ eV. The energy of the primary particle is measured in proportion to the total number of photons yielded by the shower.  Monitoring atmospheric clarity is required to calibrate for atmospheric propagation losses of light between the shower and the detector. Obtaining this calibration requires routine measurements by specialized equipment, generally lasers, and LIDARS. While essential, this equipment is challenging to construct, maintain, and calibrate, especially in the remote deserts where fluorescence detectors are located. Active systems are limited in their beams can not be so bright as to swamp the fluorescence detectors and cause saturation of the data acquisition systems. Additional methods to measure the aerosol optical depth and cross-check these conventional measurements can be helpful, especially when no additional equipment is needed.  The High Resolution Fly's Eye observatory (HiRes), located at Dugway, Utah, USA features two fluorescence detector stations separated by 12.6 km. (See \\cite{proto00} and \\cite{B2002}.) Each station views nearly the full azimuth. The HiRes-1 station has one ring of telescopes that view 3.5 to 16 degrees of elevation. A second ring of telescopes extends the elevation coverage of the HiRes-2 station to 30 degrees. Each telescope features a 3.75 m$^2$ mirror that focuses light onto a camera of 256 photomultiplier tubes (PMTs). Each PMT views approximately (1$^\\circ$x1$^\\circ$).  The atmosphere is modeled using molecular scattering and ozone absorption as a baseline. The remaining attenuation is attributed to aerosols. The HiRes experiment includes steerable lasers used to measure aerosol attenuation. For more details, see \\cite{A2005}.  This paper describes an independent measurement of aerosol optical depth that uses air-showers viewed in stereo. It has the advantage that it is insensitive to the absolute photometric calibration of the detector hardware and the total fluorescence yield that are two of the largest uncertainties of the fluorescence technique. In the sense that air-showers are a natural part of the primary data sample, the technique incurs no additional cost. Furthermore, the measurement is made over the band of wavelengths that air-showers produce, and for the range of distances over which HiRes measures air-showers. Nor does this analysis require a comprehensive reconstruction of the air-shower light profile, energy, or primary particle composition; it is enough to reconstruct the shower axis, identify segments of the shower viewed in common by two detectors, and apply a set of selection criteria.  We note that the technique has limitations. The relatively low flux of extensive air-showers restricts the statistical power of the technique to the measurement of one parameter, total aerosol optical depth, averaged over years. To reduce sensitivity of the result to details of the aerosol vertical distribution, the technique assumes that most of the aerosol is distributed below the shower segments used in the analysis. The assumption is supported by an analysis of laser shots (Abassi 2005), that found that the aerosol vertical distribution is consistent with an average scale height of about 1 km. For this analysis we use shower segments at least 1 km above the detectors.  \\begin{figure} \\epsscale{1} \\label{lbal-diagram} \\plotone{fig0-diagram.eps} \\caption{Diagram of a cosmic ray air-shower as viewed in stereo.} \\end{figure} ",
        "conclusions": "Stereo cosmic ray showers are used to measure $\\tau$ in a manner that is independent of absolute detector calibration. While it cannot replace the hourly and daily measurements obtained by specialized equipment, this method gives a cross check of the amount of aerosols present as averaged over extended periods. This technique may be of use to other experiments that measure air-showers with more than one fluorescence detector station. A trade-off between sensitivity and statistics is expected, depending on the distance between stations. We note, in passing, that this work is the first reported systematic use of air-showers to estimate atmospheric clarity."
    },
    "0601/astro-ph0601058_arXiv.txt": {
        "abstract": "{ The joint influence of numerical parameters such as the number of particles $N$, the gravitational softening length $\\varepsilon$ and the time-step $\\Delta t$ is investigated in the context of galaxy simulations. For isolated galaxy models we have performed a convergence study and estimated the numerical parameters ranges for which the relaxed models do not deviate significantly from its initial configuration. By fixing $N$, we calculate the range of the mean interparticle separation $\\lambda(r)$ along the disc radius. Uniformly spaced values of $\\lambda$ are used as $\\varepsilon$ in numerical tests of disc heating. We have found that in the simulations with $N=1\\,310\\,720$ particles $\\lambda$ varies by a factor of $6$, and the corresponding final Toomre's parameters $Q$ change by only about $5$ per cent. By decreasing $N$, the $\\lambda$ and $Q$ ranges broaden. Large $\\varepsilon$ and small $N$ cause an earlier bar formation. In addition, the numerical experiments indicate, that for a given set of parameters the disc heating is smaller with the Plummer softening than with the spline softening. For galaxy collision models we have studied the influence of the selected numerical parameters on the formation of tidally triggered bars in galactic discs and their properties, such as their dimensions, shape, amplitude and rotational velocity. Numerical simulations indicate that the properties of the formed bars strongly depend upon the selection of $N$ and $\\varepsilon$. Large values of the gravitational softening parameter and a small number of particles results in the rapid formation of a well defined, slowly rotating bar. On the other hand, small values of $\\varepsilon$ produce a small, rapidly rotating disc with tightly wound spiral arms, and subsequently a weak bar emerges. We have found that by increasing $N$, the bar properties converge and the effect of the softening parameter diminishes. Finally, in some cases short spiral arms are observed at the ends of the bar that change periodically from trailing to leading and vice-versa -- the wiggle. ",
        "introduction": "} One of the long-standing issues in $N$-body simulations is how a specific selection of numerical parameters influence an outcome. In astrophysical and cosmological simulations there are a number of these parameters that may play an important role, and therefore they have to be carefully selected. For instance, it is well known that collisionless simulations of galaxies are performed with a number of particles $N$, that is far less than the number of stars in a typical galaxy, hence, it induces artificial non-uniformities of the gravitational field. In order to smooth the field and to avoid the divergence of the gravitational force when the separation between particles becomes small, close encounters between particles should be avoided. This is usually done by softening the force between two particles using some modified form of the Newtonian gravitational potential. For example, the Plummer softening is widely used since the first simulations of \\citet{b2}. Other softening algorithms exist that give a better approximation to the Newtonian force such as, for example, higher powers of the Plummer softening \\citep{b4} or the spline softening \\citep{b27}, among others \\citep{Dehnen01}. The degree of the force softening is defined by the gravitational softening parameter $\\varepsilon$. Small values of $\\varepsilon$ cause high relaxation rates but give a better resolution of the internal structure of the system. On the other hand, a large softening reduces the relaxation rate, but errors in the force calculation increase because of the deviation from the Newtonian force.  The influence of $\\varepsilon$ on the force calculation error was discussed by \\citet{b31}, where he found that its optimal choice is related to the minimum error that is characterized by the so-called bias and variance. He also gives two empirical expressions to estimate $\\varepsilon$ for the \\citet{b55} and \\citet{b28} density profiles. This work was later extended to another configurations \\citep{b64, b4} and to other softening kernels \\citep{Dehnen01}. Although these criteria give $\\varepsilon$ that minimizes the force errors, they have been derived for rigid Monte-Carlo realizations, each of which, when evolved in time, will drift to its individual equilibrium state. Such criteria allow to estimate only the magnitude of $\\varepsilon$ that minimizes the initial perturbation due to the modification of gravity, but ignores the further evolution. For this reason it is important to check whether these criteria apply for dynamical evolving systems such as galaxies.  Two-body relaxation effects and force calculation errors depend on $\\varepsilon$ and $N$, and both of them should be simultaneously minimized. It is, however, difficult to measure them separately \\citep{b30}, e.g. they manifest themselves as a joint error in the energy of the particles \\citep{b60}.  Another source of error stems from an inadequate choosing of the integration time-step $\\Delta t$, which also results in poor energy conservation. This latter error is typically influenced by the maximum acceleration, which is constrained by the softening parameter. In addition, there are truncation and roundoff errors, but they are small and can be easily controlled.  The problem of an adequate selection of numerical parameters ($N$, $\\varepsilon$, $\\Delta t$) for $N$-body simulations has been widely discussed in the literature, usually separately. In particular, concerning the effect of the softening parameter in 2D galaxy simulations, it was established that in addition to reduction of relaxation, it plays a stabilizing role, and the meaning of a given value of the $Q$ parameter strongly depends on the value of $\\varepsilon$ \\citep{Miller71, Romeo94}. Indeed, relaxation can heavily affect the results of numerical studies of disc stability \\citep[e.g,][]{White88}. Criteria for choosing $\\varepsilon$ for such systems, based on dynamical requirements and type of softening, were discussed by \\citet{Romeo94, Romeo97, Romeo98}. A 3D study to estimate the numerical parameters was made by \\citet{b26, b30} in which the influence of the softening parameter and the number of particles on the force errors were investigated. Also, there is a rich discussion concerning the choosing of numerical parameters and its effects on cosmological simulations and small scale structure formation \\citep[e.g.,][]{b63}. On the other hand, analytical estimations give severe limits on the reliability of $N$-body calculations. For instance, it was shown that the orbits of particles diverge exponentially with time from their original trajectories \\citep{Goodman93, b65}.  For galaxy simulations, the above mentioned effects influence any dynamic process, such as the formation of bars in spirals that we study in the present work. This problem involves physical as well as numerical aspects. Since the 70s, several numerical simulations have analysed this problem and found that cold, rotationally supported stellar discs are globally unstable to bisymmetric distortions and quickly evolve into well defined spontaneous bars \\citep[e.g.,][]{b58}. The formation of a bar in unstable disc galaxies was extensively studied in both two- and three-dimensional simulations \\citep[e.g.,][]{b66, b50, b38, b52}. It was established that the bar formation can be suppressed by introducing high enough velocity dispersions \\citep{b57}, or by a massive spherical halo \\citep{b56}. The 2D-simulations have shown that the bar's length is not only defined by physical parameters but also strongly depends on the magnitude of the softening parameter \\citep{b66, b50}. Furthermore, recent 3D-simulations of isolated galaxies \\citep{b14, b16, b6} have shown that the bar semi-major axis is more than twice as long as the observed length, which is close to the exponential length-scale of the disc \\citep{b12}. On the other hand, it was found that bar rotation velocities were slowed down to less than twice the observed velocities in just a few Gyrs \\citep{b16, b6}. This result coincides with predictions of \\citet{b73} that bars are slowed down by dynamical friction with a massive halo. However, it was argued that this slowdown is associated with the low halo spatial resolution that diminishes the rotational velocity of the formed bar \\citep{b40}. A high resolution simulation with $20$ million particles \\citep{ONeill03} still indicates the bar slowdown, although it is not so drastic. Moreover, it was found that even in a stable galactic disc a bar may emerge due to discreteness of the halo \\citep*{b69}, and \\citet{b70, Athan03} showed that these effects are related to the disc's angular momentum redistribution and resonances of halo particles.  Numerical experiments performed by other authors have shown that bars may also form in disc galaxies which interact with a companion galaxy \\citep*{b33, b23, b36}. A systematic 2D study of tidally triggered bar formation \\citep{b36} in spiral galaxies has shown that bar formation depends simultaneously on the shape of the rotational curve, the disc-to-halo mass ratio, the halo concentration, the strength and the geometry of the perturbation. These results have been partially confirmed by self-consistent 3D simulations \\citep*{b10, b67, b74}, but the formed bars are found to be slower than in isolated models, indicating a possible explanation for the dichotomy of bar characteristics found by \\citet{b12}.  Motivated by the above arguments, in this work we investigate for an isolated galaxy in equilibrium and for the collision of two galaxies, the range of the numerical parameters ($N$, $\\varepsilon$, $\\Delta t$) that permit us to minimize numerical artifacts in simulations. In addition, we compare the spline gravitational softening with the Plummer one. For the collision of two disc galaxies we study the influence of the numerical parameters on the formation and characteristics of tidal bars.  This work is organized as follows: In Sect. \\ref{BHDic} we describe the initial conditions of the galaxy model and the way in which the numerical parameters are computed. We obtain a particular range for ($N$, $\\varepsilon$, $\\Delta t$), which in Sect. \\ref{conv-stu} is tested to observe deviations from an initial configuration. Based on a convergence study we choose the most suitable values for the numerical parameters. Then, in Sect. \\ref{coll} we analyse the influence of the parameters on the bar formation after the collision of two spirals. Finally, in Sect. \\ref{discu} we discuss the results and in Sect.\\ref{conclusions} outline our conclusions. ",
        "conclusions": "In this work we have investigated the influence of the triad of numerical parameters ($N$, $\\varepsilon$, $\\Delta t$) on the properties of an equilibrium isolated galaxy model and the dynamics of two interacting galaxies. The main results for the galaxy equilibrium models can be summarized as follows: \\begin{enumerate}  \\item The transfer of the disc angular momentum to the bulge and the halo depends on $\\varepsilon$, $N$, and type of softening.  \\item The convergence study indicates that for $N=491\\,520$ and $\\varepsilon=[0.001-0.01]$, the range of final values of $Q$ changes only by $14$ per cent, while for smaller $N$, the $\\varepsilon$ and $Q$ range broadens. Further decrease of numerical heating is observed for larger $N$ simulations, namely, for $N = 1\\,310\\,720$, $Q$ only varies $5$ per cent.  \\item The final value of the Toomre stability parameter decreases linearly with the increasing of the Plummer softening.  \\item For the same values of $N$, $\\Delta t$ and $\\varepsilon$, the spline softened simulations produce bars earlier than those performed with Plummer softening. \\end{enumerate}  For the galaxy collision models it may be concluded that: \\begin{enumerate}  \\item The susceptibility of galaxy models to tidal bar formation and its dissolution depend strongly on the chosen numerical parameters. Therefore, carefully chosen parameters have to be taken for specific simulations.  \\item For sufficiently small $N$, the size of the bar formed during an encounter, may become too long or even not appear depending on the value of $\\varepsilon$. With the increase of $N$ artificial effect of the softening is reduced.  \\item The short spiral arms that are formed at the ends of the bar show a periodic change from trailing to leading and vice-versa -- the wiggle.  \\item The spline softening gives smaller bars than in case of the Plummer softening even when they have the same potential depth.  \\item The properties of the merger remnant are also affected by the softening. The remnants of cold systems maintain their barred thick disc shape, whereas those of hot systems produce rather an ellipsoid. \\end{enumerate}"
    },
    "0601/astro-ph0601602_arXiv.txt": {
        "abstract": "The astronomical dark matter could be made of weakly interacting massive species whose mutual annihilations should produce antimatter particles and distortions in the corresponding energy spectra. The propagation of cosmic rays inside the Milky Way plays a crucial role and is briefly presented. The uncertainties in its description lead to considerable variations in the predicted primary fluxes. This point is illustrated with antiprotons. Finally, the various forthcoming projects are rapidly reviewed with their potential reach. ",
        "introduction": "\\label{sec:propagation}  Supersymmetric neutralinos or Kaluza--Klein particles could explain the dark matter that is indirectly observed around galaxies, inside their clusters and on cosmological scales. These species are hunted for by several experimental collaborations. As they annihilate inside the Milky Way halo, they produce high--energy cosmic rays and in particular antimatter particles such as antiprotons, antideuterons and positrons. These are rare elements which are already manufactured inside the galactic disk through the spallations of primary cosmic rays on the interstellar gas. That conventional process provides the natural background inside which the DM induced antimatter signature may be burried and whose spectral distortions could reveal an exotic signal. It is therefore crucial to derive as precisely as possible the energy spectra of the various antimatter species irrespective of the production mechanisms and to evaluate how well they are known. To reach that goal, a precise understanding of cosmic ray propagation is mandatory. We briefly sketch here how that propagation is understood and how well it is modeled.  Interstellar nuclei are accelerated by the passage of supernovae induced shock waves. The sources for primary cosmic ray nuclei are therefore located inside the disk of the Milky Way. Primaries propagate then inside the galactic magnetic fields and undergo collisions on its irregularities -- the Alfv\\'en waves. This motion is described in terms of a diffusion process whose coefficient $K(E) \\; = \\; K_{0} \\, \\beta \\, {\\mathcal R}^{\\delta}$ increases with the rigidity ${\\mathcal R}$ of the cosmic ray particle as a power law. Because the scattering centers move with a velocity $V_{a} \\sim$ 20 to 100 km s$^{-1}$, a second order Fermi mechanism yields some diffusive reacceleration whose coefficient $K_{EE}$ may be expressed as \\beq K_{EE} \\; = \\; {\\displaystyle \\frac{2}{9}} \\, V_{a}^{2} \\, {\\displaystyle \\frac{E^{2} \\beta^{4}}{K(E)}} \\;\\; . \\eeq Ionization, adiabatic and Coulomb energy losses must also be taken into account through the energy loss rate $b^{\\rm loss}(E)$. Finally, galactic convection could wipe cosmic rays away from the disk with a velocity $V_{C} \\sim$ 5 to 15 km s$^{-1}$. The number of particles par unit of volume and space $\\psi = dn / dE$ may be derived form the master equation \\beq V_{C} \\, \\partial_{z} \\Psi \\; - \\; K \\, \\Delta \\Psi \\; - \\; \\partial_{E} \\left\\{ b^{\\rm loss}(E) \\, \\Psi \\; + \\; K_{EE}(E) \\, \\partial_{E} \\Psi \\right\\} \\; = \\; q \\;\\; , \\label{master_equation} \\eeq that applies for any cosmic ray species -- primary or secondary -- as long as the various production processes are appropriately described in the rate $q$. That general scheme may be implemented either numerically~\\cite{Strong:1998pw,Moskalenko:2004fe} or semi--analytically through the two--zone model discussed in~\\cite{Usine1}.  The latter method allows an easier derivation of the theoretical uncertainties associated to the various parameters at stake -- namely $K_{0}$, $\\delta$, $V_{a}$, $V_{C}$ and the thickness $L$ of the confinements layers that extend above and beneath the Milky Way disk. The space of these propagation parameters has been extensively scanned~\\cite{Usine1} in order to select the allowed regions where the predictions on B/C -- a typical CR secondary to primary ratio -- match the observations. Tens of thousands of propagation models have survived that crucial test. The diffusion parameters are indeed still largely undetermined contrary to what is generally believed in our community. Recently, the same conclusion has been independently reached~\\cite{Lionetto:2005jd} with the help of a fully numerical code~\\cite{Strong:1998pw} in which the convective wind $V_{C}$ increases linearly with vertical heigh $z$. However, the B/C ratio could not be accounted for when galactic convection and diffusive reacceleration were both implemented at the same time. ",
        "conclusions": ""
    },
    "0601/astro-ph0601328_arXiv.txt": {
        "abstract": "{The question of the degree of order in the magnetic fields of relativistic jets is important to any understanding of their production. Both vector-ordered (e.g. helical) and disordered, but anisotropic fields can produce the high observed degrees of polarization. We outline our models of jets in FR\\,I radio galaxies as decelerating relativistic flows. We then present theoretical calculations of the synchrotron emission from different field configurations and compare them with observed emission from FR\\,I jets. We show that large-scale helical fields (with significant poloidal and toroidal components) are inconsistent with observations. The combination of an ordered toroidal and disordered poloidal component is consistent with our data, as is an entirely disordered field. Jets must also contain small, but significant amounts of radial field. ",
        "introduction": "This paper addresses two questions: \\begin{enumerate} \\item Is the magnetic field in extragalactic radio jets vector-ordered or does it have many reversals? \\item What is its three-dimensional structure? \\end{enumerate} Observations of polarized synchrotron emission at frequencies where Faraday rotation is negligible can be used to infer a two-dimensional projection of the field structure on the plane of the sky (we refer to the {\\em apparent magnetic field direction}, which is the perpendicular to the observed ${\\bf E}$-vector position angle). This emission is an emissivity-weighted integral through the jet, affected by relativistic aberration if the flow speed is comparable with $c$ but independent of reversals in the field. Measurement of Faraday rotation can potentially be used to determine the vector field component along the line of sight within a jet, but the thermal plasma responsible must be within the jet volume, rather than in front of it, as usually appears to be the case (e.g. Laing et al. 2006).  We have developed models of jets in low-power FR\\,I (Fanaroff \\& Riley 1974) radio sources as decelerating relativistic flows. By fitting to deep radio images, we can determine the geometry and the distributions of velocity and emissivity. Our technique also constrains the three-dimensional structure of the magnetic field. We cannot determine this structure uniquely, but we can estimate the ratios of the longitudinal, toroidal and radial components and eliminate some of the simpler proposed field geometries.  Our fundamental assumption is that jets are intrinsically symmetrical and relativistic. Aberration then acts differently on radiation from the approaching and receding jets, so their observed synchrotron images are two-dimensional projections of the field structure viewed from different directions. From these, subject to some additional assumptions, we can infer the three-dimensional structure of the field. One key assumption is axisymmetry: if this holds, we can also use the symmetry of the transverse brightness and polarization profiles to distinguish (at least in some cases) between vector-ordered and disordered fields. ",
        "conclusions": ""
    },
    "0601/astro-ph0601434_arXiv.txt": {
        "abstract": "We determine the number counts and $z=0$--5 luminosity function for a well-defined, homogeneous sample of quasars from the Sloan Digital Sky Survey (SDSS).  We conservatively define the most uniform statistical sample possible, consisting of 15,343 quasars within an effective area of 1622 deg$^2$ that was derived from a parent sample of 46,420 spectroscopically confirmed broad-line quasars in the 5282 deg$^2$ of imaging data from SDSS Data Release Three.  The sample extends from $i=15$ to $i=19.1$ at $z\\lesssim3$ and to $i=20.2$ for $z\\gtrsim3$. The number counts and luminosity function agree well with the results of the Two-Degree Field QSO Redshift Survey (2QZ) at redshifts and luminosities where the SDSS and 2QZ quasar samples overlap, but the SDSS data probe to much higher redshifts than does the 2QZ sample. The number density of luminous quasars peaks between redshifts 2 and 3, although uncertainties in the selection function in this range do not allow us to determine the peak redshift more precisely.  Our best fit model has a flatter bright end slope at high redshift than at low redshift.  For $z<2.4$ the data are best fit by a redshift-independent slope of $\\beta = -3.1$ ($\\Phi(L) \\propto L^{\\beta}$).  Above $z=2.4$ the slope flattens with redshift to $\\beta \\gtrsim -2.37$ at $z=5$.  This slope change, which is significant at the $\\gtrsim5$-sigma level, must be accounted for in models of the evolution of accretion onto supermassive black holes. ",
        "introduction": "The advent of the Two Degree Field (2dF) QSO Redshift Survey (2QZ; \\markcite{bsc+00,csb+04}{Boyle} {et~al.} 2000; {Croom} {et~al.} 2004) and Sloan Digital Sky Survey (SDSS; \\markcite{yaa+00}{York} {et~al.} 2000) has resulted in a more than ten-fold increase in the number of known quasars over the past decade.  While the evolution of quasars and active galactic nuclei (AGN) in general has been of considerable interest since the first identification of quasar redshifts \\markcite{sch63,sch68}({Schmidt} 1963, 1968), there has been a resurgence of interest in the subject as a result of recent work in understanding the role of AGN in galaxy evolution.  In particular, the formation of bulges and supermassive black holes at the centers of galaxies appear to be intimately related (the so-called ${\\rm M_{BH}}-\\sigma$ relationship; \\markcite{fm00,gbb+00,tgb+02}{Ferrarese} \\& {Merritt} 2000; {Gebhardt} {et~al.} 2000; {Tremaine} {et~al.} 2002), emphasizing the importance of understanding the role that quasar activity plays in the formation and evolution of the galaxy population as a whole.  It has also been argued that feedback mechanisms \\markcite{beg04}(e.g., {Begelman} 2004) may play an important role in determining the ${\\rm M_{BH}}-\\sigma$ relationship and the co-evolution of black holes and the spheroid component of their host galaxies \\markcite{dcs+03,wl03,gds+04,so04,dsh05}(e.g., {Di Matteo} {et~al.} 2003; {Wyithe} \\& {Loeb} 2003; {Granato} {et~al.} 2004; {Scannapieco} \\& {Oh} 2004; {Di Matteo}, {Springel}, \\&  {Hernquist} 2005). Furthermore, an accurate description of the quasar luminosity function (QLF) is needed to map the black hole accretion history of the Universe \\markcite{yt02}(e.g., {Yu} \\& {Tremaine} 2002) and determine how quasars contribute to the feedback cycle.  Until recently, the quasar population was parameterized by a broken power law in luminosity with a peak in space density at $z\\sim$2--3. The luminosity at the power-law break has most often been characterized by ``pure luminosity evolution'', whereby the rarity of luminous quasars today is a result of a fixed population of quasars becoming less luminous with time \\markcite{who94,csb+04}({Warren}, {Hewett}, \\& {Osmer} 1994; {Croom} {et~al.} 2004).  However, pure luminosity evolution fails beyond the peak (at $z\\sim2.5$) of the luminous quasar space density \\markcite{ssg95,fss+01}({Schmidt}, {Schneider}, \\& {Gunn} 1995; {Fan} {et~al.} 2001).  Furthermore, hard X-ray surveys \\markcite{Ueda03,bcm+05}(e.g., {Ueda} {et~al.} 2003; {Barger} {et~al.} 2005), which probe both optically obscured AGN and substantially fainter optical quasars, have found that AGN evolution is best fit by a model in which less luminous AGN peak in space density at smaller redshifts.  This behavior has been termed ``cosmic downsizing,'' whereby the most massive black holes did most of their accreting in the distant past, while less massive objects underwent active accretion in the more recent past \\markcite{cbb+03,mer04,hkb+04}(e.g., {Cowie} {et~al.} 2003; {Merloni} 2004; {Heckman} {et~al.} 2004).  While X-ray and infrared surveys \\markcite{Ueda03,hsl+04,tuc+04,bcm+05,hms05}(e.g., {Ueda} {et~al.} 2003; {Haas} {et~al.} 2004; {Treister} {et~al.} 2004; {Barger} {et~al.} 2005; {Hasinger}, {Miyaji}, \\& {Schmidt} 2005) provide a more complete census of non-stellar nuclear activity in galaxies than do optical surveys, and radio surveys have demonstrated that the decline of quasars at high redshift is not due to dust obscuration \\markcite{wjs+05}({Wall} {et~al.} 2005), the optical luminosity function remains a powerful diagnostic tool for our understanding of luminous AGNs.  This is, in no small part, because of the large areas covered by optical surveys such as the SDSS ($\\sim10,000\\,$deg$^2$).  Sensitive hard X-ray and IR surveys, although sampling much higher AGN densities than optical surveys, suffer from much smaller survey areas ($\\lesssim1\\,$deg$^2$) and have difficulty constraining the AGN population where it is intrinsically least dense (e.g., the most luminous and highest redshift objects).  This paper presents a long sought after result: the optical luminosity function of a large, homogeneous sample of luminous type 1 quasars covering the entire span of observed quasar redshifts.  This overall goal has already been roughly met by splicing together different surveys \\markcite{pei95}(e.g., {Pei} 1995): in particular, the combination of results from the 750 deg$^2$ 2QZ survey to $z\\sim2.2$ \\markcite{bsc+00,csb+04}({Boyle} {et~al.} 2000; {Croom} {et~al.} 2004) (also earlier work including that of \\markcite{sg83}{Schmidt} \\& {Green} 1983, \\markcite{kk88}{Koo} \\& {Kron} 1988, \\markcite{bsp88}{Boyle}, {Shanks}, \\& {Peterson} 1988, and \\markcite{hfc93}{Hewett}, {Foltz}, \\& {Chaffee} 1993, among others); the photometrically-selected COMBO-17 survey \\markcite{wwb+03}({Wolf} {et~al.} 2003) spanning $1.2<z<4.8$ but only 0.78 deg$^2$; and $z>3$ surveys such as \\markcite{who94}{Warren} {et~al.} (1994), \\markcite{ssg95}{Schmidt} {et~al.} (1995), \\markcite{kdd+95}{Kennefick}, {Djorgovski}, \\& {de  Carvalho} (1995), and \\markcite{fss+01}{Fan} {et~al.} (2001).  Herein we accomplish this goal using a single carefully constructed subset of the SDSS-DR3 data that was designed for maximal homogeneity; this subsample has $\\sim15,000$ quasars selected over $\\sim1600$ deg$^2$.  This paper describes the luminosity function of quasars with $15.0<i<19.1$ and $0.0 \\lesssim z \\lesssim 3.0$, and extending to $i<20.2$ for higher redshifts up to $z\\sim5$.  The magnitude limits of the SDSS survey only just approach the ``break'' magnitude ($b_J\\sim19.5$) seen in the 2QZ number counts, and are nearly two magnitudes brighter than that of the combined SDSS+2dF (2SLAQ) sample from \\markcite{rca+05}{Richards} {et~al.} (2005), both of which are restricted to $z\\lesssim2.2$. The SDSS data complement these smaller but deeper optical surveys both by extending to $z\\sim5$ and having superb multicolor photometry.\\footnote{Work on the QLF at $z\\sim6$ involves going beyond the main selection algorithm of SDSS \\markcite{fss+01}({Fan} {et~al.} 2001).}  In \\S~\\ref{sec:data} we describe the creation of a homogeneous statistical sample of quasars from the SDSS data.  The selection function is presented in detail in \\S~\\ref{sec:selfunct}.  The less technically-minded reader may choose to skip to Section~\\ref{sec:number_counts} where we present the number counts relationship for our sample of quasars.  In \\S\\S~\\ref{sec:kcorrect} and \\ref{sec:lumfunct} we discuss the application of $K$-corrections to the data and the luminosity function (both binned and maximum likelihood) that we derived after doing so.  Finally some discussion and conclusions are presented in \\S~\\ref{sec:conclusions}.  Throughout this paper we use a $\\Lambda$ cosmology with $\\Omega_m=0.3$, $\\Omega_\\Lambda=0.7$, a Hubble Constant of $H_0=70\\,{\\rm km\\,s^{-1}\\,Mpc^{-1}}$ \\markcite{svp+03}(e.g., {Spergel} {et~al.} 2003), and luminosity distances determined according to \\markcite{hog99}{Hogg} (1999) for this cosmology. ",
        "conclusions": "\\label{sec:conclusions}  One of the most interesting results to come out of recent AGN surveys is the evidence in favor of ``cosmic downsizing,'' wherein the peak of AGN activity occurs at higher redshifts for more luminous objects than less luminous objects \\markcite{cbb+03,Ueda03,mer04,bcm+05}({Cowie} {et~al.} 2003; {Ueda} {et~al.} 2003; {Merloni} 2004; {Barger} {et~al.} 2005).  Comparison of X-ray, infrared, and optical surveys requires careful consideration of the fact that many groups find that the ratio of obscured (type 2) to unobscured (type 1) AGN is inversely correlated with AGN luminosity (e.g., \\markcite{law91,Ueda03,hao05}{Lawrence} 1991; {Ueda} {et~al.} 2003; {Hao} {et~al.} 2005a; but see \\markcite{tu05}{Treister} \\& {Urry} 2005). Ignoring this effect and examining the most uniform luminous sample that we can form over the largest redshift range ($M_i<-27.6$), Figure~\\ref{fig:fig20n} shows that the peak in type 1 quasar activity occurs between $z=2.2$ and $z=2.8$.  Unfortunately, this redshift range is the least sensitive in the SDSS and subject to large error, see Figure~\\ref{fig:fig9n}.  A substantial observing campaign for $z\\sim2.5$ quasars that are buried in the stellar locus (i.e., a sample with close to unity selection function in this redshift region) is needed to resolve this issue.  To this end \\markcite{Chiu04}{Chiu} (2004) and \\markcite{jiang05}{Jiang et al.} (2006) describe complete (i.e., not sparsely sampled) surveys of quasars in the mid-$z$ range to address this problem.  In addition, near-IR selected samples such as can be obtained from {\\em Spitzer Space Telescope} photometry should be able to better isolate the peak redshift of luminous type 1 quasars \\markcite{Brown05}({Brown} {et~al.} 2006).  Our most interesting result is the flattening of the slope of the QLF with increasing redshift.  This flattening has been demonstrated before using small samples of high-$z$ quasars \\markcite{ssg95,fss+01}({Schmidt} {et~al.} 1995; {Fan} {et~al.} 2001), but never so robustly and over such a large redshift range as with these data.  While there is little overlap in luminosity between the lowest and highest redshift data (deeper surveys at high redshift are clearly needed), previous constraints on the QLF and the presumption that the QLF will be well-behaved outside of the regions explored (e.g., that the slope does not get {\\em steeper} for faint high-redshift quasars), suggests that the slope change is due to redshift and not luminosity.  Small area samples such as the most sensitive hard X-ray surveys \\markcite{Ueda03,bcm+05}({Ueda} {et~al.} 2003; {Barger} {et~al.} 2005) and the COMBO-17 survey \\markcite{wwb+03}({Wolf} {et~al.} 2003) primarily probe the low-luminosity end of the QLF, where the slope is flatter, thus it is not surprising that they systematically find flatter slopes (see Fig.~\\ref{fig:fig19n}).  Our confirmation of the flattening of the high-redshift slope has significant consequences in terms of our understanding of the formation and evolution of active galaxies, particularly in light of the popularity of recent models invoking kinetic and radiative AGN feedback in the evolution of galaxies \\markcite{sr98,fab99,wl03,hhc+05}(e.g., {Silk} \\& {Rees} 1998; {Fabian} 1999; {Wyithe} \\& {Loeb} 2003; {Hopkins} {et~al.} 2005a).  A particularly interesting explanation for the steepness of the low-redshift QLF comes from the model of \\markcite{wl03}{Wyithe} \\& {Loeb} (2003).  In their scenario, the predicted slope at low redshift is much flatter than the observed slope (see Fig.~1 of \\markcite{wl03}{Wyithe} \\& {Loeb} 2003) for the most luminous quasars.  They argue that the so-called ``break'' between less and more luminous quasars for $z\\lesssim2$ is the result of ``the inability of gas to cool inside massive dark matter halos'', thus preventing the formation of $v_c\\gtrsim500\\,{\\rm km\\,s^{-1}}$ galaxies and their resulting luminous quasars in the most massive dark matter halos.  Such an idea may be in conflict with the work of \\markcite{hhc+05z}{Hopkins et al.} (2005b) who find that the break in the QLF occurs naturally in their models.  In their case the break occurs at the maximum of a (roughly) log normal peak luminosity distribution with more luminous quasars accreting near Eddington and less luminous quasars perhaps accreting at lower rates.  However, in both the \\markcite{wl03}{Wyithe} \\& {Loeb} (2003) and \\markcite{hhc+05z}{Hopkins et al.} (2005b) PLE models the predicted slopes for luminous quasars are actually {\\em steeper} at high redshift than at low redshift (see Fig.~1 in \\markcite{wl03}{Wyithe} \\& {Loeb} 2003 and Fig.~11 in \\markcite{hhc+05z}{Hopkins et al.} 2005b) --- {\\em opposite of that which we observe}.  In the case of \\markcite{hhc+05z}{Hopkins et al.} (2005b), the evolution to high redshift is determined simply by adjusting the break luminosity of their $0<z<3$ model to fit the existing high redshift data.  Thus, to match our observed flattening at high redshift, \\markcite{hhc+05z}{Hopkins et al.} (2005b) will either need to change their model or (at the very least) the details of its extrapolation to higher redshift.  To explain a flatter QLF slope at high redshift in their model, one would need to invoke a broader distribution of quasar peak luminosities at high redshifts than low, with relatively more low-luminosity objects at low redshift than high.  Thus our observations and future observations of even fainter type 1 quasars at $z>3$ provide an important litmus tests for models of galaxy evolution.  While our analysis was based on over 15,000 quasars, to form this uniform sample we were forced to drop over half of the objects in the DR3 Quasar Catalog because of inhomogeneity in the selection algorithms.  Data Release 5 of the SDSS, which is planned to occur in mid-2006, will contain more than twice as many ``new\" quasars as are in our DR3 uniform sample.  We can define appreciably larger samples yet, using the photometric selection and photometric redshift techniques of \\markcite{rng+04}{Richards} {et~al.} (2004a) and \\markcite{wrs+04}{Weinstein} {et~al.} (2004); such methods will result in a sample approaching a million quasars probing appreciably further down the luminosity function, albeit at the price of less certain redshifts.  We will also connect the low-luminosity end of the quasar luminosity function to that of Seyfert galaxies measured from the SDSS galaxy sample \\markcite{hst+05}({Hao} {et~al.} 2005b) and explore the continuity of the AGN population at low redshifts and luminosities.  We will explore the dependence of the luminosity function on color, to determine, for example, whether the luminosity function of intrinsically red quasars \\markcite{rhv+03}({Richards} {et~al.} 2003) has the same slope and varies with redshift the same way that blue quasars do; differences could indicate correlations of color with a physical parameter such as accretion rate, mass, or orientation, or possible redshift- or luminosity-dependent dust obscuration.  Finally, a number of groups are carrying out deeper spectroscopic quasar surveys based on deep SDSS photometry (see \\markcite{rca+05,jiang05}{Richards} {et~al.} 2005; {Jiang et al.} 2006), and we can look forward to a yet more comprehensive view of the quasar luminosity function in a few years' time."
    },
    "0601/astro-ph0601678_arXiv.txt": {
        "abstract": "The dynamics of accretion discs around galactic and extragalactic black holes may be influenced by their magnetic field. In this paper we generalise the fully relativistic theory of stationary axisymmetric tori in Kerr metric of Abramowicz et al.\\shortcite{Abr78} by including strong toroidal magnetic field and construct analytic solutions for barotropic tori with constant angular momentum. This development is particularly important for the general relativistic computational magnetohydrodynamics that suffers from the lack of exact analytic solutions that are needed to test computer codes. ",
        "introduction": "Accretion discs of black holes have been a subject of intensive observational and theoretical studies for decades. In most of the studies the discs are treated as purely fluid flows \\cite{ShS,FiM,Abr78,Ree}. That is their magnetic fields are considered as dynamically weak and not important for the force balance determining the disc structure. On the other, it has been long recognised that MHD-turbulence may hold the key to the nature of the disc viscosity that enables the inflow of matter into the black hole \\cite{ShS}. This expectation was confirmed with the discovery of the magneto-rotational instability \\cite{BaH} that has become now one of the major topics in theoretical and numerical studies of accretion discs. In the simulations of radiatively inefficient discs they remain relatively weakly magnetized with the magnetization parameter $\\beta =(\\mbox{gas pressure})/(\\mbox{magnetic pressure}) \\simeq 10-100$ within the main body of the disc and rising up to $\\beta \\simeq 1$ only near its inner edge and in the plunging region \\cite{HKDH,MG}. However, cooling accretion flows may well evolve towards the state with $\\beta<1$ where the vertical balance against the gravitational force of the central object is achieved by means of magnetic pressure alone \\cite{MNM,PBB}.  On the other hand, there is a general consensus that dynamically strong, ordered, poloidal magnetic field is required both for  acceleration and collimation of the relativistic jets that are often produced by the astrophysical black hole-accretion disc systems. Such magnetic field is a key ingredient both in the models where the jet is powered by the black hole \\cite{BZ} and in the models where it is powered by the accretion disc \\cite{BR,BP}. The origin of this field is not very clear. The most popular idea is that it is carried into central parts of the disc-hole system by the accreting flow itself and that the net magnetic flux gradually builds up there during the long-term evolution of the system \\cite{BR,TM}.  This is more or less what is observed in recent numerical simulations of radiatively inefficient discs with initially weak poloidal field, \\cite{HKDH,MG}. However, the strength of magnetic field that can be reached in this way is still unclear, e.g. \\cite{LOP,SU,Meier,McK}. The most extreme case, where the disc pressure is dominated by ordered poloidal magnetic field is argued by Meier \\shortcite{Meier} as a model for the low/hard state of X-ray binaries. Although much weaker magnetic field is required to explain the observed power of quasar jets within the Blandford-Znajek theory,  the pressure of poloidal magnetic field still has to be of the same order as the radiative pressure at the inner edge of the radiatively supported accretion discs \\cite{BBR}. Bogovalov and Kel'ner \\shortcite{BK} constructed a model of accretion disc where its angular momentum is carried away by a magnetized wind alone and the inflow of matter and magnetic field is driven entirely by the magnetic torque applied to the disc by the wind. These and other results show that the structure of accretion discs with dynamically strong ordered magnetic is a matter of significant astrophysical interest.  General steady-state axisymmetric magnetized dicks are described by a very complex equation of the Grad-Shafranov type which does not allow general analytic solutions \\cite{Lov}. However, the equilibrium models with pure azimuthal  (toroidal) field are much simpler and can be constructed using the approach developed for unmagnetised discs \\cite{Abr78,Koz}. In fact, the predominant motion in accretion discs is a differential rotation and one would expect the azimuthal magnetic field to dominate in the interior of such discs.  Okada et al.\\shortcite{Oka} constructed a particular exact solution for the problem of equilibrium torus with purely azimuthal magnetic field.  In their approach the Paczynski \\& Wiita \\shortcite{PW} potential is utilised to introduce the black hole gravity and the flow dynamics is described by the non-relativistic equations. Although this approach seems to work fine in the case of non-rotating black holes, it is no longer satisfactory in the astrophysically important case of rapidly rotating black hole where full relativistic treatment is required. In this paper we develop the general relativistic theory of magnetized tori around rotating black holes. Our analysis closely follows the one of Kozlowski et al.\\shortcite{Koz} and Abramowicz et al.\\shortcite{Abr78}, including the notation.   Even if tori with dynamically strong magnetic fields do not exist in nature the exact solutions presented here can still be useful for numerical general relativistic magnetohydrodynamics, GRMHD, which has attracted a lot of interest recently \\cite{KSK,Kom01,DvH,GMT,Kom,Due,Ant}. One of the problems with general relativistic computational magnetohydrodynamics is the lack of exact analytic and semi-analytic solutions that could be used to verify computer codes. The current list of such solutions includes 1) the spherical accretion onto a nonrotating black hole with monopole magnetic field where the magnetic field is dynamically passive \\cite{KSK}, 2) the perturbative force-free solution of Blandford-Znajek \\shortcite{BZ}, where both the pressure and the inertia of matter is neglected and the black hole rotates slowly, and 3) the so-called Gammie flow, that is a one-dimensional solution of the Weber-Davis type that applies only to the equatorial plane of a black hole \\cite{Gam}. One may also try to utilize the axisymmetric semi-analytic self-similar solutions constructed recently by Meliani et al.\\shortcite{Mel} but those apply only along the symmetry axis. Obviously, the magnetized torus solution is very welcome an addition to this scarce list. In fact, it is the only solution that is not only a) fully multi-dimensional and b) involves dynamically important magnetic field, but c) also applies in the case of a rapidly rotating black hole. In fact, we have already used this solution to test the 2D GRMHD codes described in Komissarov \\shortcite{Kom01,Kom}.   Throughout the paper we use the relativistic units where $c=G=M=1$ and $(-+++)$ signature for the space-time geometry. Following Anile \\shortcite{Ani} the magnetic field is rescaled in such a way that the factor $4\\pi$ disappears from the equations of relativistic MHD. ",
        "conclusions": "In this paper we have generalized the relativistic theory of thick accretion discs around Kerr black holes \\cite{FiM,Abr78,Koz} by including dynamically strong toroidal (azimuthal) magnetic field. As expected, this inclusion of magnetic field leads to a whole new class of equilibrium solutions that differ in strength and spacial distribution of the field. In particular, we have described the way of constructing barotropic tori with constant angular momentum -- under such conditions the differential equations of magnetostatics are easily integrated and reduce to simple algebraic equations.  By now it has become clear that magnetic field plays many important roles in the dynamics of astrophysical black hole-accretion disc systems. However, the structure of this field may be more complex than just azimuthal loops and as the result the analytic solutions constructed in this paper may turned out not to be particularly suitable for modelling real astrophysical objects. Even so they will be certainly helpful in testing computer codes for general relativistic MHD that are becoming invaluable tools in the astrophysics of black holes/accretion discs."
    },
    "0601/astro-ph0601352_arXiv.txt": {
        "abstract": "The determination of chemical abundances from stellar spectra is considered a mature field of astrophysics. Digital spectra of stars are recorded and processed with standard techniques, much like samples in the biological sciences. Nevertheless, uncertainties typically exceed 20\\%, and are dominated by systematic errors. The first part of this paper addresses what is being done to reduce measurement errors; and what is not being done, but should be. The second part focuses in some of the most exciting applications of stellar spectroscopy in the arenas of galactic structure and evolution, the origin of the chemical elements, and cosmology. ",
        "introduction": "Stellar absorption lines, first discovered in the solar spectrum circa 1804, can be used to quantify the proportions of different chemical elements in stars. The strength of the lines depends on the abundance of absorbers, but also on the environment where the absorption takes place: the stellar atmosphere. The determination of chemical abundances requires an accurate knowledge of the physical conditions (e.g. temperature, density, radiation field) in the outer layers of the star. In other words, inferring chemical abundances from spectra is just a part of the problem of the physics of stellar atmospheres, and cannot be decoupled from it.  After Eddingon, Milne, and others established the theoretical basis, the field of quantitative stellar spectroscopy  has evolved over time to become almost an industry. Stellar abundances have shed light on the origin of the elements, the early universe,  what the interior of the stars look like, or how galaxies form and evolve. As modeling is refined, so are the quality and reach of the observations, making possible to use stellar abundances to tackle new problems.  Nowadays,  a small army of some $10^3$ researchers are {\\it professionally} dedicated to measure chemical abundances from stellar spectra in the entire world. Practitioners favor B-type stars on the warm side, and F- and G-type stars on the cool side.  B-type stars can be observed at large distances, while the cooler F- and G-types live long, allowing us to study the past.  Cooler stars have far more complex spectra, shaped by molecular absorption and dust.  In this paper, with no ambition of formally reviewing the field, I simply highlight recent accomplishments and tendencies, both in methods (Section \\ref{howto}) and applications (\\S \\ref{whatfor}). The closing section is devoted to a personal (and surely biased) reflection on where the field is going, and where one wishes it would go. ",
        "conclusions": "Recent noteworthy trends in research on stellar abundances include the development of improved model atmospheres for solar-like stars considering hydrodynamics, the availability of full-NLTE line-blanketed model atmospheres for hot stars, significant improvements in the opacities and the treatment of clouds for the coolest stars and brown dwarfs, the emergence of large spectroscopic surveys producing public data bases (e.g., SDSS, Elodie, S$^4$N, SEGUE, GALEX), significant progress toward the automation of spectroscopic analyses, and in availability of accurate atomic data for neutron-capture elements.  A desirable short-term future would include 3D model atmospheres for late-type stars of all types, 1D models in full NLTE, NLTE line formation in 3D models, solar atlases observed at different positions on the disk to test line formation and model atmospheres, stronger efforts to measure/compute the necessary atomic and molecular data, stronger efforts to use the newly available atomic and molecular data, full analysis automation, a new generation of all-in-one archives (or virtual observatories) for both multi-wavelength observations and models, and ultra-high resolution spectrographs on large telescopes. Let's make it happen!"
    },
    "0601/astro-ph0601164_arXiv.txt": {
        "abstract": "We present the first hydrodynamic, multi-dimensional simulations of He-shell flash convection. Specifically, we investigate the properties of shell convection at a time immediately before the He-luminosity peak during the $15^\\mem{th}$ thermal pulse of a stellar evolution track with initially two solar masses and metallicity $Z=0.01$. This choice is a representative example of a low-mass asymptotic giant branch thermal pulse. We construct the initial vertical stratification with a set of polytropes to resemble the stellar evolution structure. Convection is driven by a constant volume heating in a thin layer at the bottom of the unstable layer. We calculate a grid of 2D simulations with different resolutions and heating rates. Our set of simulations includes one low-resolution 3D run.  The computational domain includes $11.4$ pressure scale heights. He-shell flash convection is dominated by large convective cells that are centered in the lower half of the convection zone. Convective rolls have an almost circular appearance because focusing mechanisms exist in the form of the density stratification for downdrafts and the heating of localized eddies that generate upflows. Nevertheless, downdrafts appear to be somewhat more focused. The He-shell flash convection generates a rich spectrum of gravity waves in both stable layers above and beneath the convective shell. The magnitude of the convective velocities from our 1D mixing-length theory model and the rms-averaged vertical velocities from the hydrodynamic model are consistent within a factor of a few. However, the velocity profile in the hydrodynamic simulation is more asymmetric, and decays exponentially inside the convection zone. An analysis of the oscillation modes shows that both g-modes and convective motions cross the formal convective boundaries, which leads to mixing across the boundaries.  Our resolution study shows consistent flow structures among the higher resolution runs, and we see indications for convergence of the vertical velocity profile inside the convection zone for the highest resolution simulations. Many of the convective properties, in particular the exponential decay of the velocities, depend only weakly on the heating rate. However, the amplitudes of the gravity waves increase with both the heating rate and the resolution. ",
        "introduction": "\\label{kap:intro}  We present hydrodynamic simulations of the interior shell convection zone driven by the He-shell flash in thermal pulse Asymptotic Giant Branch (AGB) stars.  \\subsection{AGB evolution} AGB stars are the final evolution stage of low- and intermediate mass stars before formation of a white dwarf \\citep{iben:83b}.  Nuclear production in AGB stars contribute to the chemical evolution of galaxies \\citep[for example][]{travaglio:04,renda:04}. The chemical yields are support the interpretation stellar abundance from dwarf-spheroidal galaxy satellites of the Milky Way relative to stellar abundances of galactic halo stars \\citep{venn:04,geisler:05}. This may provide important clues about the cosmological origin of our galaxy. Another role of AGB stars have recently been discussed as a potential source of the abundance anomalies observed in globular cluster star members \\citep{ventura:02,denissenkov:03a}, a question that is still debated.  AGB stars produce substantial amounts of Li, C, N, \\nezw, \\nadr, and the neutron-heavy Mg isotopes.  These stars are the progenitors of planetary nebulae nuclei and white dwarfs, including those thought to orbit C-rich extremely metal-poor stars with \\spr\\ signature \\citep[CEMP-s, ][]{beers:05}. The favoured interpretation of their strongly non-solar abundance pattern is the pollution of the EMP star with material transfered from the white dwarf progenitor AGB star. This interpretation is largely based on the fact that the abundance patterns of CEMP-s stars in general resemble the C- and \\spr\\ overabundance predicted for EMP AGB stars. However, very significant discrepancies between observed and predicted abundances exist which could indicate some serious problem with current 1D stellar evolution and nucleosynthesis simulations. Most likely those problems are related to the 1D approximate treatment of mixing that originate from multi-dimensional, convective fluid motions.  Especially at extremely low metallicity, AGB stars can produce many species in a primary mode, i.e.\\ only with the H and He available from the Big Bang \\citep[for example]{siess:02,herwig:04a}.  AGB stars are also the nuclear production site of the main component of the \\sprn\\ \\citep{busso:99}, which produces roughly half of all trans-iron elements \\citep{arlandini:99}. The interpretation of high-precision laboratory measurements of isotopic ratios of heavy elements in pre-solar meteoritic SiC grains depends sensitively on the thermodynamic and the closely related mixing conditions at the bottom of the He-shell flash convection zone \\citep{zinner:98,lugaro:02b,lugaro:02a}. In this particular layer the $\\nezw(\\alpha,\\n)\\mgfu$ reaction releases neutrons very rapidly, and activates many s-process branchings. Experimental measurements of nuclear properties of the unstable branch-point nuclei are now underway \\citep{reifarth:05} and in conjunction with grain measurements will provide new constraints for mixing at convective boundaries in the future.  Currently, only 1D stellar evolution models exist to describe the processes of the AGB stellar interior quantitatively. However, while spherical symmetry is globally a very good approximation, a detailed description of mixing processes induced by the hydrodynamic fluid-flows requires multi-dimensional simulations. Only over long times, equivalent to many convective turn-overs, can they be described accurately by the average quantities provided by the commonly used local theories, like the mixing-length theory \\citep{boehm-vitense:58}. However, it is well known that these theories do not describe the quantitative properties of mixing at convective boundaries. For these reasons, we start our investigations of the hydrodynamic properties of stellar interior convection with the He-shell flash convection in thermal pulse AGB stars.    \\subsection{Hydrodynamic simulations of stellar convection and application to 1D stellar evolution calculations} Most multi-dimensional simulations of stellar convection address the outer convection zones. For example the solar convection zone including the surface granulation \\citep{nordlund:82,stein:98}, convection below the surface \\citep{stein:89} or the role of magnetic fields and rotation in the solar convection zone \\citep{brun:04b} Other simulations include work on deep envelope convection in giants \\citep{porter:94,porter:00,freytag:02,robinson:04}, shallow surface convection, for example in A stars and white dwarfs \\citep{freytag:96} and the general properties of slab convection \\citep{hurlburt:86}. Hydrodynamic simulations of stellar interiors include an investigation of semi-convection in massive stars by \\citet{merryfield:95} and core convection in rotating A-type stars \\citep{browning:04}.  The work by \\citet{freytag:96} motivated \\citet{herwig:97} to include the exponential overshooting derived from hydrodynamical simulations to all convective boundaries. The e-folding distance for convective boundaries in the stellar interior was derived semi-empirically for main-sequence stars, and then applied to AGB stars \\citep{herwig:99a}. Similar approaches were considered by \\citet{mazzitelli:99}, \\citet{ventura:00}, \\citet{cristallo:04}, and \\citet{althaus:05}.  However, convection in shells driven by nuclear energy release is different than envelope convection. Typically, the Mach numbers are smaller throughout the convection zone, and the bordering radiative layers are relatively more stable. Accordingly, the semi-empirical determination of the overshooting value $f_\\mem{ov}$\\footnote{Here, the overshoot formulation is $ D_{\\rm OV} = D_0 \\exp{\\left( \\frac{-2 z}{f_\\mem{ov} \\cdot H_{\\rm p}}\\right)}$, where $D_0$ is the mixing-length theory mixing coefficient at the base of the convection zone, $z$ is the geometric distance to the convective boundary, $H_{\\rm p}$ is the pressure scale height at the convective boundary, and $f_\\mem{ov}$ is the overshooting parameter \\citep{herwig:97}.} gives smaller values for the stellar interior ($f_\\mem{OV} < 0.016$) compared to the shallow surface convection ($0.25 < f_\\mem{OV} < 1.0$)  according to \\citet{freytag:96}. For deep envelope (or core) convection $f_\\mem{ov}$ can be expected to be considerably smaller since the ratio of the Brunt-V\\\"ais\\\"al\\\"a time-scales of the stable to the unstable layers decrease with increasing depth, i.e. adiabaticity \\citep{herwig:97}.  Relevant in this regard are the 2D simulations of interior O-burning shell convection in massive stars immediately before a supernova explosion of \\citet{bazan:98} and \\citet{asida:00}. In these simulations hydrodynamic fluid motions are not confined to the convectively unstable region. The authors describe the excitation of gravity waves in the adjacent stable layers \\citep{hurlburt:86}, as well as an exponential decay of convective velocities at and beyond the convective border. As a result highly reactive fuel in the form of C is mixed from above into the O-shell. However, the authors caution that the quantitative results may suffer some uncertainty due to numerical effects like resolution. The main difference between these simulations and our regime of the He-shell flash is the lower driving energy in our case, and we expect correspondingly lower Mach numbers, less mixing and less overshooting.  \\citet{young:05b} present initial results on updated O-shell 2D and 3D convection simulations and apply their findings on hydrodynamics mixing to 1D stellar evolution models of supernova progenitors \\citep{young:05a}. The internal structure in these models is markedly different from standard models, with important implication for the supernova explosion properties.    \\subsection{About this paper} The main purpose of our study is to obtain a characterization of the hydrodynamic properties of shell convection in low-mass stars which is determined by weaker driving due to smaller nuclear energy generation compared to massive stars. We investigate the typical morphology, the dominating structures and the time-scales of shell-flash convection, and how these properties depend on resolution and driving energy. We identify areas of parameter space that require further study in two- and three-dimensions. In our present work we do not intend to quantify mixing across the convective boundary, but we want to gain insight for future simulations that shed more light on this important question.  In the following section we describe the 1D stellar evolution code and the hydrodynamics codes used in this work. \\kap{kap:setup} describes the stellar evolution sequence and the specific model that serves as a template for the hydrodynamic simulations. Then we describe the results in \\kap{kap:results}; specifically we describe the general properties of the convection, the onset of convection, the evolved convection and the convective and oscillation modes observed in our simulations. We close the paper with a discussion (\\kap{sec:concl}) of the results, in particular the comparison of the 1D MLT velocity profile and the time-averaged vertical velocity profile from our hydrodynamic simulations. ",
        "conclusions": "\\label{sec:concl}  We presented the first set of hydrodynamic simulations specifically addressing the conditions in the convection zone close the peak luminosity of the He-shell flash. Compared to previous calculations of other stellar convection zones our simulations reveal some important differences. One of the most important is the stiffness of the convective boundaries of He-shell flash convection in contrast to the shallow surface convection of A-stars and white dwarfs as studied by \\citet{freytag:96}. Other examples include solar-type convection \\citep{dintrans:05} or much of the parameter range in relative stability between the unstable and stable layers studied by \\citet{rogers:05}.  Even in the simulations of core convection in A-type stars \\citep{browning:04} convective plumes can cross the convective boundary significantly. In shallow surface convection zones the convective downdrafts easily cross the convection boundary and penetrate deep into the stable layers beneath. No significant g-mode excitation has been detected by \\citet{freytag:96}, which is probably due to the absence of a stiff resonance floor with which convective plumes and the pressure field interact. However, it seems that because of this particular feature of the shallow surface convection zone \\citet{freytag:96} were able to study the decay of the convective velocity field in isolation. In our He-shell flash convection we do see exponential decay of the convective velocity field as well. However it already starts inside the convection zone, and at or beyond the border quickly intermingles with the effects of the gravity waves.  The comparison to the new O-shell convection simulations by \\citet[][and 2006b, in prep]{maekin:06a} confirms the importance of g-modes for mixing and to correctly account for the physics at the convective boundary. Similar to our simulations their convective boundaries are very stiff, with overshooting of convective plumes into the stable layer reduced to a minimum. More important are horizontal wave motions which eventually induce turbulent mixing. A partial mixing zone at the convective boundaries is minimal. The O-shell convection simulations include $\\mu$-gradients, contrary to our similations. $\\mu$ gradients increase the impenetrable character of the convective boundaries significantly, and we plan to include this effect in our simulations in the future.  We tentatively compare the mixing length convective velocity profile from our 1D stellar evolution template model with the averaged vertical velocities in the hydrodynamic simulation with realistic heating rate in \\abb{fig:v_hdmlt}. Obviously, there are many differences in the assumptions of the two calculations, including geometry, micro-physics and of course the treatment of convection itself. Nevertheless, we find that the absolute value of the velocities inside the convection zone are similar. We do observe a different profile in the two representations. The hydrodynamic simulation show a more pronounced slope in the velocity profile within the convection zone than the MLT profile. The peak in the velocity profile from the hydrodynamic simulation is close to the bottom of the convection zone. The most important difference is the significant decline of the velocity field already inside the convection zone. Near the lower boundary the vertical velocity is almost 3 orders of magnitude lower than the peak value. This may have implications for s-process nucleosynthesis as nuclei are on average exposed much longer to the hottest temperature at the very bottom of the convection zone. This should be considered in particular for cases that involve branchings sensitive to the convective time-scale of He-shell flash convection, like \\xeac\\ \\citep{reifarth:04}.  Consistent with previous work on stellar interior convection we observe a rich spectrum of g-modes excited in the stable layers. These g-modes interact with the decaying convective flows at the boundary of the unstable region. Our simulations show that He-shell flash convection displays the main ingredients observed in previous stellar convection simulations, but with different relative contributions of the various components. The actual overshooting of convective plumes across the convective boundaries is much smaller than in shallow surface convection and O-shell convection, while the excitation of internal gravity is probably stronger because the boundaries are stiffer.  Although there is a significant overlap of the two processes we have shown that the analysis of the oscillation modes by means of the $k$-$\\omega$ diagram allows us to disentangle the different contributions to the velocities. With this approach we observe that there is a finite mixing of convective motions accross the convective boundary. We postpone quantifying this mixing to future work. In addition, we plan to quantify the contribution from wave mixing at He-shell flash convection boundaries.  We have identified a number of numerical problems at low heating rates and low resolution that need to be dealt with in future 3D runs. These include the diffusion of entropy from above into the top zones of the unstable layer, impeding the development of the correct hydrodynamic flow-pattern. Our convergence and heating rate study shows that simulations with larger heating rate are in some ways equivalent to runs with higher resolution. Properties like the exponential decay rate of the convective velocity field seem to be only weakly dependent on the heating rate. However, the excitation of g-modes and their amplitudes depends sensitively on both the resolution and the heating rate. While we do see signs of convergence for the convective properties of our simulations, this is not the case for the g-modes.  We consider 2D vs.\\ 3D systematics, the effect of $\\mu$-gradients, a realistic treatment of nuclear energy generation, and a more detailed study of the interaction of wave mixing with convective overshooting as our most immediate tasks for future improvements. In addition we can make our simulations more realistic by adopting an apppropriate stellar equations of state, by accounting for the spherical geometry of the He-shell, and by including a realistic gravitational potential."
    },
    "0601/astro-ph0601487_arXiv.txt": {
        "abstract": "{ We present new images of the Supernova Remnant (SNR) G127.1+0.5 (R5), based on the 408 MHz and 1420 MHz continuum emission and the HI-line emission data of the Canadian Galactic Plane Survey (CGPS). The radio spectrum of the central compact source (0125+628) is analyzed in the range 178 MHz - 8.7 GHz, indicating a flat spectrum with synchrotron self-absorption below 800 MHz. The SNR's flux density at 408 MHz is 17.1$\\pm$1.7 Jy and at 1420 MHz is 10.0$\\pm$0.8 Jy, corrected for flux densities from compact sources within the SNR. The SNR's integrated flux density based spectral index (S$_{\\nu}$$\\propto$$\\nu$$^{-\\alpha}$) is 0.43$\\pm$0.10. The respective T-T plot spectral index (derived from the relative size of brightness temperature variations between two frequencies, see text for details) is 0.46$\\pm$0.01.  There is no evidence at 1$\\sigma$ for spatial variations in spectral index within G127.1+0.5. In particular, we compared the northern shell, southern shell and central diffuse region.  HI observations show structures associated with the SNR in the radial velocity range of -12 to -16 km$/$s, suggesting G127.1+0.5's distance is 1.15 kpc. The estimated Sedov age is 2 - 3 $\\times$$10^{4}$ yr. ",
        "introduction": "We report new CGPS radio observations of the  bilateral symmetry shell-type SNR G127.1+0.5 (R5). G127.1+0.5 is considered as a younger remnant (18000 yr, Milne 1988) for its appearance and brightness. Because compact sources are detected near the center of the SNR, there have been many studies at radio, optical and X-ray wavelengths of the SNR and the compact sources (Pauls 1977; Pauls et al., 1982; Fu\\\"rst et al., 1984; Joncas et al., 1989; Xilouris et al., 1993; Kaplan et al. 2004). But its basic physical features are still uncertain, such as its distance and age and radio spectrum. G127.1+0.5's radio spectral index in frequency range of 0.4 GHz - 5 GHz has been estimated as between 0.45 and 0.6 previously (F\\\"urst et al., 1984; Joncas et al., 1989). The spectrum of the southern shell is slightly steeper than for the northern shell based on observations with a competitive resolution but lower sensitivity than current observations. In this paper, we present the SNR's continuum images at higher sensitivity than previously at both 408 MHz and 1420 MHz in order to determine its flux densities and spectrum. We search the HI-line emission for detecting possible interactions of the remnant with the surrounding gas and estimating its distance and age. ",
        "conclusions": "We present new maps of G127.1+0.5 at 408 MHz and 1420 MHz and study the radio spectrum, corrected for point source flux densities, using different methods.  All results are consistent with the best determined value from the T-T plot method of 0.46$\\pm$0.01. We find no evidence for spectral index variation across G127.1+0.5. A probable association of HI features with G127.1+0.5 yields a distance of 1.15 kpc for G127.1+0.5, consistent with association with the open cluster NGC 559. The HI deficit associated with the SNR gives a local density of $\\simeq 2 cm^{-3}$, which yields a dense shell transition radius of 7.9 pc, roughly equal to the observed radius. So G127.1+0.5 should be approximately at the age where it enters the dense shell phase, and have an age of $2-3\\times 10^4$yr."
    },
    "0601/astro-ph0601214_arXiv.txt": {
        "abstract": "We have determined the proper motion (PM) of the Large Magellanic Cloud (LMC) relative to four background quasi-stellar objects, combining data from two previous studies made by our group, and new observations carried out in four epochs not included the original investigations.  The new observations provided a significant increase in the time base and in the number of frames, relative to what was available in our previous studies.  We have derived a total LMC PM of $\\mu$ = ($+$2.0$\\pm$0.1) mas yr$^{-1}$, with a position angle of $\\theta$ = (62.4$\\pm$3.1)$^\\circ$.  Our new values agree well with most results obtained  by other authors, and we believe we have clarified the large discrepancy between previous results from our group. Using published values of the radial velocity for the center of the LMC, in combination with the transverse velocity vector derived from our measured PM, we have calculated the absolute space velocity of the LMC.  This value, along with some assumptions regarding the mass distribution of the Galaxy, has in turn been used to calculate the mass of the Milky Way.  Our measured PM also indicates that the LMC is not a member of a proposed stream of galaxies with similar orbits around our galaxy. ",
        "introduction": " ",
        "conclusions": "\\subsection{The ALP-PAM Discrepancy}  Given the implications of the result obtained by ALP for the PM of the LMC, in relation to our understanding of the interactions between the Galaxy and the Magellanic Clouds (see, e.g, Momany \\& Zaggia, 2005), and the reality of streams of galaxies with similar orbits around the Galaxy (see, e.g, Piatek et al., 2005), the main objective of the present work was to clarify the discrepancy between the previous determinations of the PM of the LMC by our group: the \"ALP-PAM Discrepancy\".  In this section, we further elaborate on some of the thoughts originally proposed in PAM, in order to explain the discrepancy of ALP, originally with PAM, and now also with the new result presented in this paper.  First, the fact that the observations used here were made with essentially the same equipment and instrumental set-up as those by ALP, precludes any arguments relating the observed discrepancy to the existence of systematic errors in the observational data.  Such errors would affect our data in the same way way as those of ALP.  Second (as explained in \\S 2), in the reduction process of the ALP and PAM data incorporated in the present work we adopted the same QSO and reference stars centroid coordinates (x,y) used in those works. Furthermore, the new data included in the present calculations was processed using the same procedure used in ALP to obtain the (x,y) coordinates. Therefore, the centroid coordinates should not be a source of a systematic error either.  The subsequent procedures to obtain the PM were also basically the same, the sole exception being the inclusion of a quadratic term in the transformation equations used for the registration (also included in PAM's equations, but not in ALP's).  Tests carried out using ALP's data alone showed however that the effect of including quadratic terms is marginal (as was suspected), and does not account for the observed discrepancy.  Considering that our current result -which includes re-processed data from ALP- agrees quite well with measurements by other groups, we conclude that ALP's results might be affected by an unidentified systematic error in Decl.  Since in the present work we used ALP's unmodified (x,y) coordinates, we believe that this error could have originated in the processing of the Decl. PM instead of the coordinates themselves.  It should be pointed out that the UCAC2-based result from Momany \\& Zaggia (2005), which is consistent with that from ALP, is currently also considered to be affected by an as yet unidentified systematic error (Momany \\& Zaggia, 2005; Kallivayalil et al., 2005).  We would finally like to note that our new result for field Q0459-6427 is consistent with PAM.  \\subsection{Membership of the LMC to a Stream}  Lynden-Bell \\& Lynden-Bell (1995), have proposed that the LMC, together with the SMC, Draco and Ursa Minor, and possibly Carina and Sculptor, define a stream of galaxies with similar orbits around our galaxy. Their models predict a PM for each of member of the stream, which can be compared to their measured PM to evaluate the reality of the stream.  For the LMC they predict PM components of [$\\mu_{\\alpha}$cos$\\delta$,$\\mu_{\\delta}$] = [+1.5,0] mas yr$^{-1}$, giving a total PM of $\\mu$ = +1.5 mas yr$^{-1}$, with a position angle of $\\theta$ = 90$^\\circ$.  A comparison of this prediction with our result [$\\mu$ = ($+$2.0$\\pm$0.1) mas yr$^{-1}$, $\\theta$ = (62.4$\\pm$3.1)$^\\circ$], shows that our measured values of $\\mu$ and $\\theta$ are, respectively, 5.1$\\sigma$ and 8.9$\\sigma$ away from the predicted values.  This result indicates that the LMC does not seem to be a member of the above stream (it is worth mentioning that Piatek et al. (2005), using HST data, have concluded that Ursa Minor is not a member of this stream).\\\\    \\newpage"
    },
    "0601/astro-ph0601508_arXiv.txt": {
        "abstract": "{We report the detection during the {\\it Chandra Orion Ultradeep Project} (COUP) of two soft, constant, and faint X-ray sources associated with the Herbig-Haro object HH\\,210. HH\\,210 is located at the tip of the NNE finger of the emission line system bursting out of the BN-KL complex, northwest of the Trapezium cluster in the OMC-1 molecular cloud. Using a recent H$\\alpha$ image obtained with the {\\it ACS} imager on board {\\it HST}, and taking into account the known proper motions of HH\\,210 emission knots, we show that the position of the brightest X-ray source, \\object{COUP\\,703}, coincides with the emission knot 154-040a of HH\\,210, which is the emission knot of HH\\,210 having the highest tangential velocity (425\\,km\\,s$^{-1}$). The second X-ray source, \\object{COUP\\,704}, is located on the complicated emission tail of HH\\,210 close to an emission line filament and has no obvious optical/infrared counterpart. Spectral fitting indicates for both sources a plasma temperature of $\\sim$0.8\\,MK and absorption-corrected X-ray luminosities of about $10^{30}$\\,erg\\,s$^{-1}$ (0.5--2.0\\,keV). These X-ray sources are well explained by a model invoking a fast-moving, radiative bow shock in a neutral medium with a density of $\\sim$12000\\,cm$^{-3}$. The X-ray detection of COUP\\,704 therefore reveals, in the complicated HH\\,210 region, an energetic shock not yet identified at other wavelengths. ",
        "introduction": "\\begin{figure}[!ht] \\centering \\includegraphics[width=\\columnwidth]{Hk231_f1.ps} \\caption{{\\it HST/WFPC2} image of HH\\,210 for epoch 2000.70 in [\\ion{S}{ii}] \\citep[adapted from Fig.~2 of][]{doi02}. Each pixel is $0\\farcs1$ on a side, the color stretch is linear. HH\\,210 shows a typical bow shock structure containing numerous emission knots and filaments. The two crosses mark epoch 2003.04 positions of COUP\\,703 and COUP\\,704 \\citep{getman05b}; cross sizes indicate the total positional uncertainties of these X-ray sources ($0\\farcs14$). The white arrows indicate velocity and epoch 2003.04 position for emission knots with proper-motion measured in [\\ion{S}{ii}] by \\citet{doi02}. The white $0\\farcs08$-radius circles indicate the optical resolution of {\\it HST}. } \\label{map} \\end{figure}  Outflow activity is ubiquitous in young stellar objects (YSO), and intimately connected to the accretion process building up the mass of stars. The interaction of outflow from YSO and the interstellar medium produces bow shocks associated with optically luminous small nebulae, called Herbig-Haro (HH) objects (\\citealt{herbig50,haro52};\\citealp[ see review by][]{reipurth01}). Optical emission from HH objects includes Hydrogen Balmer lines, [\\ion{O}{i}], [\\ion{S}{ii}], [\\ion{N}{ii}], [\\ion{O}{iii}], tracing plasma with temperatures of several $10^3$ to $10^5$\\,K. During these last years a handful of HH objects have been detected in X-rays, tracing plasma with temperatures of a few $10^6$\\,K, as expected from the high velocities of these shocks \\citep[e.g.,][]{raga02}. The emission knot H of \\object{HH\\,2} was the first detected in X-rays; it also exhibits strong H$\\alpha$ and free-free radio continuum emission \\citep{pravdo01}. \\object{HH\\,80} and \\object{HH\\,81} are excited by a massive protostar, and are among the most luminous HH objects in the optical; emission knots~A and G/H of HH\\,80, and emission knot~A of HH\\,81 are now the most luminous known in X-rays \\citep{pravdo04}.  The soft X-ray emission from L1551~IRS5, associated with \\object{HH\\,154} by \\citet{favata02}, is now considered as arising from a shock at the base of the jet of this protostar binary \\citep{bally03,bonito04}. Recently, soft X-ray excesses were found in the X-ray spectra of the \\object{`Beehive' proplyd} \\citep{kastner05} and \\object{DG\\,Tau} \\citep{guedel05}, two YSOs with jets. Remarkably, in both sources, this soft and constant X-ray emission suffers less extinction than the harder and variable X-ray component likely emitted by the stellar corona, suggesting that it comes from shocks along the base of these jets.  In January 2003, the {\\it Chandra Orion Ultradeep Project} (COUP) was carried out with ACIS-I \\citep{garmire03} on board the \\cxo{} X-ray Observatory \\citep{weisskopf02}. COUP consisted of a nearly continuous exposure over 13.2\\,days centered on the Trapezium cluster in the Orion Nebula, yielding a total on-source exposure time of 838\\,ks, i.e.\\ 9.7\\,days \\citep{getman05b}. Among the 1616 COUP X-ray sources, there are only about fifty X-ray sources with the median energy of their X-ray photons lower than 1\\,keV, i.e.\\ having a very soft X-ray spectrum. The bulk of this soft X-ray source sample is associated with optical/infrared stars; two sources may be spurious; and two other sources without counterparts may be newly discovered very low mass members of the Orion nebula cluster. Of the very soft sources, only COUP\\,703 and COUP\\,704 are possibly associated with an HH object, HH\\,210 \\citep{getman05a,kastner05}. More generally there is in COUP no other X-ray source associated with HH objects \\citep{getman05a}.  HH\\,210 is located 2.7\\arcmin-north of $\\theta^1$\\,Ori\\,C in the OMC-1 molecular cloud, and is one of the brightest HH objects in this area revealed by emission of [\\ion{O}{i}] \\citep{axon84,taylor86}. It belongs to the powerful CO outflow from the BN-KL region, which produces a spectacular set of H$_2$ bow shocks and trailed wakes; it is located at the tip of the NNE finger which is one of the brightest emission systems in [\\ion{Fe}{ii}] and lacks a trailed H$_2$ wake \\citep{allen93}. HH\\,210 has a pronounced bow shape containing numerous knots and complicated filaments in [\\ion{S}{ii}] (see Fig.~\\ref{map}).  We present in Sect.~\\ref{observation} the astrometry and spectral analysis of the X-ray sources associated with HH\\,210; we discuss in Sect.~\\ref{discussion} the origin of these X-ray emissions. ",
        "conclusions": "\\label{discussion}  HH\\,210 displays the largest proper motion among the outflow fingers extending away from BN-KL, moreover the line emission is blue-shifted, suggesting a deprojected velocity of about 500\\,km\\,s$^{-1}$ \\citep{axon84,taylor86,hu96}. 154-040a\\,HH\\,210 is also the brightest feature in [\\ion{O}{iii}] and has a small bow shock attached to it \\citep{odell97,doi02}. The tip of the bow shock is seen as a knot in [\\ion{N}{ii}], H$\\alpha$, [\\ion{O}{iii}]; hence the tip is the fastest moving portion of the shock. However, the fact that the [\\ion{O}{iii}] emission requires only a shock with a speed of about 90\\,km\\,s$^{-1}$ and that the [\\ion{O}{iii}] emission is seen only at the tip of the bow, indicates that this HH object is moving in the wake of other shocks moving ahead of it \\citep{doi02}.  The observed X-ray emission can be explained with fast moving bow shocks. The postshock temperature in the adiabatic (non-radiative) portion of a shock is given by $T_\\mathrm{ps}/{\\mathrm K}=2.9 \\times 10^5/(1+X) \\times V^2_{100}$ \\citep{Ostriker88}, where $V_{100}$ is the velocity of the shock front with respect to the downstream material (the shock velocity) in units of 100\\,km\\,s$^{-1}$ and $X$ is the ionization fraction of the preshock gas ($X$=0 for a neutral medium and 1 for a fully ionized one). Thus, COUP\\,703 and COUP\\,704 plasma temperatures of $\\sim$0.8\\,MK require a shock speed of $\\sim$170\\,km\\,s$^{-1}$ for a shock moving into a mostly neutral medium, or $\\sim$240\\,km\\,s$^{-1}$ for a shock moving into a fully ionized medium. The lower speed is comparable with the bow shock velocity of 133\\,km\\,s$^{-1}$ estimated by \\citet{odell97} from the deprojected velocity and the opening angle of the bow wing.  The expected X-ray luminosity of a shock-heated source depends on the emissivity per unit volume, which depends on the plasma temperature and density, the volume of the emitting region, and the type of shock, radiative or non-radiative \\citep{raga02}. Neglecting the line emission, \\citet{raga02} find for a radiative shock $L_{\\rm r}/{\\rm erg\\,s^{-1}}=1.64\\times10^{28}\\,n_{100}\\,r^2_{\\rm b,16}\\,V^{5.5}_{100}$~; and for a non-radiative shock $L_{\\rm nr}/{\\rm erg\\,s^{-1}}=1.8\\times10^{29}\\,n_{100}^2\\,r^3_{\\rm b,16}\\,V_{100}$, where $n_{100}$ is the preshock density in units of 100\\,cm$^{-3}$ and $r_{\\rm b,16}$ is the radius of the object driving the shock in units of $10^{16}$\\,cm. The X-ray luminosity of a bow shock is the minimum of these two values. The multiplicative factor to take into account line emission is 2.5 and 3.0 at 0.1\\,MK and 1\\,MK, respectively \\citep{raga02}. The {\\it HST/ACS} observation of 154-040a~HH\\,210 shows a radius of $0\\farcs15$ (Fig.~\\ref{trichro}), or about $10^{15}$\\,cm at a distance of 450\\,pc. Using this dimension for the X-ray source, a shock speed of 170\\,km\\,s$^{-1}$, and a preshock density of about 12000\\,cm$^{-3}$, we find $L_{\\rm r}$$\\sim$$3.3 \\times 10^{29}$\\,erg\\,s$^{-1}$. Thus taking into account line emission, a bow shock flow can readily explain the X-ray luminosities observed from COUP\\,703 and COUP\\,704. citet{allen93} found a similar density of $\\sim$10$^4$\\,cm$^{-3}$ for the electronic density inferred from the [\\ion{Fe}{ii}] spectra of the emission knots. The location of HH\\,210 inside OMC-1 but close to the limit of the \\ion{H}{ii} region may explain this density and why the medium is not yet fully ionized by the radiation field of the Trapezium cluster.  We conclude that COUP\\,703 is the counterpart of the emission knot 154-040a of HH\\,210, and its X-ray emission can be explained by a radiative bow shock. The X-ray emission of COUP\\,704 can also be explained by a fast-moving, radiative shock toward the tail of HH\\,210. Optical/infrared observations are still needed at other epochs to reveal in the complicated HH\\,210 region the proper motion of the counterpart to this X-ray source."
    },
    "0601/astro-ph0601036_arXiv.txt": {
        "abstract": "We use \\textit{GALEX} (Galaxy Evolution Explorer) near-UV (NUV) photometry of a sample of early-type galaxies selected in \\textit{SDSS} (Sloan Digital Sky Survey) to study the UV color-magnitude relation (CMR). $NUV-r$ color is an excellent tracer of even small amounts ($\\sim 1$\\% mass fraction) of recent ($\\la 1$ Gyr) star formation and so the $NUV-r$ CMR allows us to study the effect of environment on the recent star formation history. We analyze a volume-limited sample of 839 visually-inspected early-type galaxies in the redshift range $0.05 < z < 0.10$ brighter than $M_{r}$ of $-21.5$ with any possible emission-line or radio-selected AGN removed to avoid contamination. We find that contamination by AGN candidates and late-type interlopers highly bias any study of recent star formation in early-type galaxies and that, after removing those, our lower limit to the fraction of massive early-type galaxies showing signs of recent star formation is  roughly $30 \\pm 3\\%$ This suggests that residual star formation is common even amongst the present day early-type galaxy population.  We find that the fraction of UV-bright early-type galaxies is 25\\% higher in low-density environments. However, the density effect is clear only in the lowest density bin. The blue galaxy fraction for the subsample of the brightest early-type galaxies however shows a very strong density dependence, in the sense that the blue galaxy fraction is lower in a higher density region. ",
        "introduction": "There is observational evidence pointing to a very simple evolutionary model for early-type galaxies. This model of \\textit{Monolithic Collapse} was first proposed by \\citet*{1962ApJ...136..748E} to explain the origin of the Milky Way halo. According to this model, the Milky Way halo formed through the rapid collapse of a cloud of gas very early on in the history of the universe, forming all of its stars in an initial burst followed by a passive evolution of the stellar population. A similar model is often invoked as the simplest explanation for the old and seemingly homogeneous stellar populations seen in early-type galaxies \\citep{1975MNRAS.173..671L}.   The apparently universal relationship between galaxy color and luminosity in early-type galaxies was first studied in detail by \\citet*{1977ApJ...216..214V}, even though the relation had been observed before \\citep{1959PASP...71..106B, 1961ApJS....5..233D, 1968AJ.....73.1008M}. This \\textit{Color-Magnitude Relation} (CMR) is often used as a tool for understanding the formation and evolution of early-type galaxies.  A seminal investigation on the optical CMR was undertaken by \\citet*{1992MNRAS.254..589B} on the elliptical galaxies in the Virgo and Coma clusters. Their study revealed a remarkably small intrinsic scatter around the mean relation. In the context of the monolithic paradigm, they interpreted the small scatter as the result of a small age dispersion amongst galaxies of the same age and the slope as a result of a mass-metallicity relation \\citep{1997A&A...320...41K}. Further, they concluded that massive early-type galaxies did not have any major episodes of star formation at redshifts z $< 2$.  More massive galaxies are likely to be in deeper potential wells and are therefore more able to retain metals ejected from supernovae from the initial generations of young stars at high redshift, leading to the observed mass-metallicity relation \\citep{1974MNRAS.169..229L}. In addition to this, the observed levels of  $\\alpha$-enhancement \\citep{1992ApJ...398...69W,1993MNRAS.265..553C, 1994MNRAS.270..743C, 1994MNRAS.270..523C, 1997A&A...320...41K, 2000AJ....120..165T} in many giant ellipticals imply that the initial formation starburst was of a very short duration of less than 1 Gyr \\citep{1993ApJ...405..538B, 1999MNRAS.302..537T}.  Later studies have found that there is no significant evolution in the optical CMR out to  z=1 and further \\citep{1997ApJ...483..582E, 1998ApJ...501..571G,1998ApJ...492..461S, 2000ApJ...541...95V, 2003ApJ...596L.143B, 2005ApJ...635..243F}. All this adds up to a picture of massive early-type galaxies forming in an initial, intense starburst at high redshift followed by a relatively-passive evolution.  However, we know since the simulations of \\citet*{1972ApJ...178..623T} that the product of a spiral-spiral merger can be an elliptical galaxy \\citep{1983MNRAS.205.1009N, 1988ApJ...331..699B, 1992ApJ...400..460H, 2003ApJ...597..893N}. An alternative approach to understanding early-type galaxies takes into account dynamical interactions and mergers. In the \\textit{Hierarchical Merger} paradigm, small galaxies form first and later assemble into larger objects \\citep{1978MNRAS.183..341W}.  The advent of semi-analytical models (SAMs) in the 1990s has greatly enhanced our understanding of galaxy evolution in such a hierarchical universe. \\citet{1996MNRAS.283L.117K} find that in the Canada-France Redshift Survey, only 1/3 of elliptical and lenticular galaxies at redshift z=1 were fully assembled and showed colors expected of old passively evolving systems. There is of course older evidence for a strong dependence of the population of early-type galaxies on density and redshift. \\citet{1980ApJ...236..351D} found that approximately 80\\% of galaxies in a sample of 55 clusters were of early-type morphology, a much higher fraction than in the field suggesting that the denser cluster environment does affect galaxy evolution. When \\citet{1984ApJ...285..426B} looked at higher redshift clusters, they found that the fraction of blue, spiral galaxies in cluster environments increased with redshift.  Later studies confirmed that this trend  was not a selection effect \\citep{1997ApJ...490..577D, 2000ApJ...541...95V}. This evolution is accompanied with an increase in merger rates \\citep{1998ApJ...497..188C, 1999ApJ...520L..95V, 2001ApJ...561..517K}.  In a purely monolithic collapse model, the star formation history of early-type galaxies is almost trivial, as they are composed of uniformly old stars. As soon as we allow for any sort of hierarchical merging, the story becomes much more complex. Rather than being uniform, the star formation histories become highly degenerate, as disparate stellar populations from progenitor galaxies are mixed together. Beyond this simple addition, the merging history of the galaxy and its progenitors adds further complication as entirely new populations are created during interactions and mergers. Thus, assigning a single age to the stellar population of an early-type galaxy is misleading - there is no single age. The typically-derived luminosity-weighted ages are in this sense nontrivial to interpret. We know now that the combined effects of age, dust, metallicity and - potentially - a multitude of progenitors are highly degenerate. \\citet*{1992MNRAS.254..589B} took the apparent uniformity and low intrinsic scatter as a very strong constraint on the evolution of the Virgo \\& Coma Early-type population. While monolithic evolution is the simplest possible explanation of these observations, however, it does not necessarily exclude other interpretations.  \\citet{2005MNRAS.360...60K} have argued using merger models that the optical early-type CMR is useful for constraining evolution models \\textit{only} if we believe \\textit{a priori} in a monolithic model. The effect of progenitor bias - the fact that a progressively larger fraction of the progenitor set of present-day ellipticals is contained in late-type star-forming galaxies at higher redshift - means that we are \\textit{not} probing the entire star formation history of early-type galaxies, but rather a progressively more biased subset. Besides, the level of star formation predicted by SAMs incorporating AGN and supernova feedback is very low; on the order of a few percent by stellar mass. Optical filters, including $U$-band, are not sufficiently sensitive to detect such a low-level star-forming activity. This is why we must turn to the UV.  The \\textit{Galaxy Evolution Explorer (GALEX)} \\citep{2005ApJ...619L...1M} near-UV filter is capable of detecting even a small ($\\sim 1\\%$ mass fraction) young stellar population and so ideal for tracing the recent star formation history of early-type galaxies. The UV color-magnitude relation allows us to identify the last important episode of star formation in galaxies. \\citet{2005ApJ...619L.111Y} have already shown using \\textit{GALEX} information that a significant fraction of massive early-type galaxies at low redshift exhibit levels of star formation undetectable in the optical but visible in the UV. Our paper presents the results of our search for the effect of environment on the recent star formation.  We assume a standard $\\Lambda$CDM cosmology with $(\\Omega_{M}, \\Omega_{\\Lambda}) =(0.3, 0.7)$ and a Hubble constant of $H_{0}=70\\, {\\rm km\\,s^{-1}\\,Mpc^{-1}}$ \\citep{2003ApJS..148..175S}. ",
        "conclusions": "We have used the UV color-magnitude relation of low redshift, massive early-type galaxies to study their recent star formation history. Our sample is volume-limited, ranging in redshift from z=0.05 to 0.1 and is limited in absolute magnitude to $M_{r} < -21.5$. Our sample is highly unlikely to be contaminated by any significant number of late-type galaxies, as all our galaxies have been visually inspected.  In order to classify galaxies by their environment, we have devised a method for measuring environment that takes the proximity, and not just the number density of neighboring galaxies, into account (see Section \\ref{env_section}). This method can easily be modified to different samples within SDSS and can take into account a larger part of the luminosity function if restricted to lower redshift limits than $z=0.1$. Our measure works very well for the brightest cluster galaxies in the C4 cluster catalog and in addition also performs for field galaxies.  In our sample of 839 early-type galaxies with $M_{r} < -21.5$, the recent star formation (RSF) galaxy fraction is $30 \\pm 2$\\%. Our ellipticals, the bulk of our sample,  have an RSF fraction of $29\\% \\pm 3$, while the lenticulars show $39\\% \\pm 5$. This implies that \\textit{residual star formation is common amongst the present day early-type galaxy population}. Our estimates are very likely lower limits on the true fractions, as our criteria for RSF are conservative in the consideration of internal extinction and the UV contribution from the old populations.  The UV color-magnitude relation differs from the optical color-magnitude relation \\citep{1992MNRAS.254..589B, 2004ApJ...601L..29H} in that it does vary more clearly with environment. The recent star formation history of early-type galaxies also varies with environment. It is well-known that more massive galaxies reside in higher-density environments (Figure \\ref{fig9}), but we show for the first time that UV-bright early-type galaxies preferentially reside in low density environments. The RSF fraction is a function of environment and drops by 25\\% from field to group but then puzzlingly remains relatively constant at higher densities, even when split into luminosity bins (Figure \\ref{fig12}). Interestingly, the most massive galaxies ($-23.82 < M_{r} \\leq -22.13$) show the strongest dependence on environment and alone exhibit a further drop in RSF fraction from medium to high density.  One possible way to understand the drop in the RSF fraction between low and medium density is in the context of ram pressure stripping. Galaxies moving fast in the deep gravitational potential of a galaxy cluster are bound to lose most of their gas during their orbital motion \\citep{1972ApJ...176....1G}. The density dependence of gas content in galaxies has long been established empirically as well \\citep{1981ApJ...247..383G}. Our RSF fraction-density relation is in the right direction. The gas goes into the ICM and so could potentially explain the star formation we do see in high density environments.  Another noteworthy observation is the fact that those early-type galaxies which have been identified as AGN - by emission lines and/or radio - are significantly bluer than those who are not (see Figure \\ref{fig3}). We have removed these AGN from our sample since we cannot disentangle the UV flux from the AGN from that of a possible young stellar population. However, the blueness of the AGN colors is intriguing - are we really just seeing the AGN itself, or is this from the star formation triggered by the jets and outflows from the AGN \\citep{1998A&A...331L...1S}? In the latter case, the RSF fraction would increase further from our estimates, and the AGN regulations might present a possible physical mechanism responsible for the star formation that we observe. In fact, \\citet{1995ApJ...452..549H}, using IUE observations of nearby Type 1 and 2 Seyferts, suggest that at most 20\\% of the UV continuum emission seen in them can originate from the nucleus itself. This means that the vast majority of our AGN candidate (removed) would qualify as RSF galaxies.  It is important to note that we only deal with \\textit{fractions} of star-forming galaxies and not actual star-formation rates. Thus, our RSF fractions simply give us an indication of how likely an early-type galaxy with certain properties and in a certain environment is to have experienced recent star formation. Neither environment, luminosity nor axis ratio seems to be the primary physical quantity that regulates recent star formation in early-type galaxies. The relative insensitivity to environment in any environment denser than the field is also surprising and warrants further study. The observational trends presented here give us new constraints for theoretical models of galaxy evolution."
    },
    "0601/astro-ph0601493_arXiv.txt": {
        "abstract": "Strong gravitational lens systems with extended sources are of special interest because they provide additional constraints on the models of the lens systems.  To use a gravitational lens system for measuring the Hubble constant, one would need to determine the lens potential and the source intensity distribution simultaneously. A linear inversion method to reconstruct a pixellated source brightness distribution of a given lens potential model was introduced by Warren \\& Dye.  In the inversion process, a regularisation on the source intensity is often needed to ensure a successful inversion with a faithful resulting source. In this paper, we use Bayesian analysis to determine the optimal regularisation constant (strength of regularisation) of a given form of regularisation and to objectively choose the optimal form of regularisation given a selection of regularisations.  We consider and compare quantitatively three different forms of regularisation previously described in the literature for source inversions in gravitational lensing: zeroth-order, gradient and curvature.  We use simulated data with the exact lens potential to demonstrate the method.  We find that the preferred form of regularisation depends on the nature of the source distribution. ",
        "introduction": "The use of strong gravitational lens systems to measure cosmological parameters and to probe matter (including dark matter) is well known \\citep*[e.g.\\ ][]{R64,KSW04}.  Lens systems with extended source brightness distributions are particularly useful since they provide additional constraints for the lens modelling due to surface brightness conservation.  In such a system, one would need to fit simultaneously the source intensity distribution and the lens potential model (or, equivalently the lens mass distribution) to the observational data.  The use of a pixellated source brightness distribution has the advantage over a parametric source brightness distribution in that the source model is not restricted to a particular parameter space.  \\citet{WD03} introduced a linear inversion method to obtain the best-fitting pixellated source distribution given a lens model and the observational data.  Several groups of people \\citep*[e.g.\\ ][]{W96, TK04, DW05, K05, BL05} have used pixellated source distributions, and some \\citep{K05, SB05} even used a pixellated potential model for the lens.  The method of source inversion described in \\citet{WD03} requires the source distribution to be ``regularised'' (i.e., smoothness conditions on the inverted source intensities to be imposed) for reasonable source resolutions.\\footnote{The source pixel sizes are fixed and are roughly a factor of the average magnification smaller than the image pixel sizes. In this case, regularisation is needed because the number of source pixels is comparable to the number of data pixels.  On the other hand, if the number of source pixels is much fewer than the effective number of data pixels (taking into account of the signal-to-noise ratio), the data alone could be sufficient to constrain the pixellated source intensity values and regularisation would play little role.  This is equivalent to imposing a uniform prior on the source intensity distribution (a prior on the source is a form of regularisation), a point to which we will return later in this article.}  For fixed pixel sizes, there are various forms of regularisation to use and the differences among them have not been addressed in detail.  In addition, associated with a given form of regularisation is a regularisation constant (signifying the strength of the regularisation), and the way to set this constant has been unclear.  These two long-standing problems were noted in \\citet{KSW04}.  Our goal in this paper is to use Bayesian analysis to address the above two issues by quantitatively comparing different values of the regularisation constant and the forms of regularisation.  \\citet{BL05} also followed a Bayesian approach for pixellated source inversions.  The main difference between \\citet{BL05} and this paper is the prior on the source intensity distribution.  Furthermore, this paper quantitatively compares the various forms of regularisation by evaluating the so-called ``evidence'' for each of the forms of regularisation in the Bayesian framework; \\citet{BL05} mentioned the concept of model comparison but did not apply it.  \\citet{DW05} use adaptive source grids to avoid the use of explicit regularisation (i.e., uniform priors are imposed since adapting the grids is an implicit form of regularisation); however, the Bayesian formalism would still be useful to set the optimal scales of the adaptive pixel sizes objectively.  Furthermore, regularised source inversions (as opposed to unregularised -- see footnote 1) permit the use of smaller pixel sizes to obtain fine structures.  The outline of the paper is as follows.  In Section \\ref{sect:BayesInf}, we introduce the theory of Bayesian inference, describing how to fit a model to a given set of data and how to rank the various models.  In Section \\ref{sect:appgl}, we apply the Bayesian analysis to source inversions in strong gravitational lensing and show a way to rank the different forms of regularisations quantitatively. ",
        "conclusions": "\\label{sect:conclusion}  We introduced and applied Bayesian analysis to the problem of regularised source inversion in strong gravitational lensing.  In the first level of Bayesian inference, we obtained the most probable inverted source of a given lens potential and PSF model $\\responseSet$, a given form of regularisation $\\regSet$ and an associated regularisation constant $\\lambda$; in the second level of inference, we used the evidence $P(\\dataVec|\\lambda, \\responseSet,\\regSet)$ to obtain the optimal $\\lambda$ and rank the different forms of regularisation, assuming flat priors in $\\lambda$ and $\\regSet$.  We considered three different types of regularisation (zeroth-order, gradient and curvature) for source inversions.  Of these three, the preferred form of regularisation depended on the intrinsic shape of the source intensity distribution: in general, the smoother the source, the higher the derivatives of the source intensity in the preferred form of regularisation.  In the demonstrated examples of first two Gaussian sources, and then a box with point sources, we optimised the evidence $P(\\dataVec|\\lambda, \\responseSet,\\regSet)$ and numerically solved for the regularisation constant for each of the three forms of regularisation.  By comparing the evidence of each regularisation evaluated at the optimal $\\lambda$, we found that the curvature regularisation was preferred with the highest value of evidence for the two Gaussian sources, and gradient regularisation was preferred for the box with point sources.  The study of the three forms of regularisation demonstrated the Bayesian technique used to compare different regularisation schemes objectively.  The method is general, and the evidence can be used to rank other forms of regularisation, including non-quadratic forms (e.g. maximum entropy methods) that lead to non-linear inversions \\citep[e.g.\\ ][]{W96, WWLH05, BL05}.  We restricted ourselves to linear inversion problems with quadratic forms of regularising function for computational efficiency.  In the demonstration of the Bayesian technique for regularised source inversion, we assumed Gaussian noise, which may not be applicable to real data.  In particular, Poisson noise may be more appropriate for real data, but the use of Poisson noise distributions would lead to non-linear inversions that we tried to avoid for computational efficiency.  Nonetheless, the Bayesian method of using the evidence to rank the different models (including noise models) is still valid, irrespective of the linearity in the inversions.  We could also use Bayesian analysis to determine the optimal size of source pixels for the reconstruction.  The caveat is to ensure that the annular region on the image plane where the source plane maps is unchanged for different pixel sizes.  Currently the smallest pixel size is limited by the region of low magnifications on the source plane.  In order to use smaller pixels in regions of high magnifications, adaptive source gridding is needed. This has been studied by \\citet{DW05}, and we are currently upgrading our codes to include this.  The Bayesian approach can also be applied to potential reconstruction on a pixellated potential grid.  \\citet*{BSK01} proposed a method to perturbatively and iteratively correct the lens potential from a starting model by solving a first order partial differential equation.  This method has been studied by \\citet{K05} and \\citet{SB05}.  The perturbation differential equation can be written in terms of matrices for a pixellated source brightness distribution and a pixellated potential, and the potential correction of each iteration can be obtained via a linear matrix inversion.  This pixellated potential reconstruction is very similar to the source inversion problem and we are currently studying it in the Bayesian framework.  The Bayesian analysis introduced in this paper is general and was so naturally applicable to both the source and potential reconstructions in strong gravitational lensing that we feel the Bayesian approach could be useful in other problems involving model comparison."
    },
    "0601/astro-ph0601170_arXiv.txt": {
        "abstract": "{ We present the first deep, optical, wide-field imaging survey of the young open cluster Collinder 359, complemented by near-infrared follow-up observations. This study is part of a large programme aimed at examining the dependence of the mass function on environment and time. We have surveyed 1.6 square degrees in the cluster, in the $I$ and $z$ filters, with the CFH12K camera on the Canada-France-Hawaii 3.6-m telescope down to completeness and detection limits in both filters of 22.0 and 24.0 mag, respectively. Based on their location in the optical ($I-z$,$I$) colour-magnitude diagram, we have extracted new cluster member candidates in Collinder 359 spanning 1.3--0.03\\,M$_{\\odot}$, assuming an age of 60 Myr and a distance of 450 pc for the cluster. We have used the 2MASS database as well as our own near-infrared photometry to examine the membership status of the optically-selected cluster candidates. Comparison of the location of the most massive members in Collinder 359 in a ($B-V$,$V$) diagram with theoretical isochrones suggests that Collinder 359 is older than \\APer{} but younger than the Pleiades. We discuss the possible relationship between Collinder 359 and IC\\,4665 as both clusters harbour similar parameters, including proper motion, distance, and age. } ",
        "introduction": "\\label{cr359:intro}  The number of known brown dwarfs (hereafter BDs) has increased dramatically over the past few years, in the field \\citep{kirkpatrick99,kirkpatrick00,burgasser02,cruz03}, as companions to low-mass stars \\citep{burgasser03a,gizis03,close03,bouy03}, in star-forming regions \\citep{lucas00,briceno02,luhman03b}, and in open clusters \\citep[][and references therein]{bouvier98,zapatero00,barrado01a,barrado02a,oliveira03}. Since the discovery of the first BDs in the Pleiades \\citep{rebolo95,rebolo96}, numerous open clusters have been targeted in the optical and in the near-infrared to uncover their low-mass and substellar populations, including the Pleiades \\citep{bouvier98,tej02,dobbie02a,moraux03}, $\\alpha$\\,Per \\citep{stauffer99,barrado02a}, M\\,35 \\citep{barrado01a}, IC\\,2391 \\citep{barrado01b}, and NGC\\,2547 \\citep{oliveira03}.  The knowledge of the Initial Mass Function (hereafter IMF) is of prime importance in understanding the formation of stars. The IMF is fairly well constrained down to about 0.5\\,M$_{\\odot}$ and well approximated by a three segment power law with $\\alpha$ = 2.7 for stars more massive than 1\\Msun{}, $\\alpha$ = 2.2 from 0.5 to 1.0\\Msun{}, and $\\alpha$ = 0.7--1.85 in the 0.08--0.5\\Msun{} mass range with a best estimate of 1.3 \\citep{kroupa02}, when expressed as the mass spectrum. However, the mass spectrum remains somewhat uncertain in the low-mass and substellar regimes. Most studies conducted in Pleiades-like open clusters suggest a power law index $\\alpha$ in the range 0.5--1.0 across the hydrogen-burning limit \\citep{martin98a,bouvier98,tej02,dobbie02a,moraux03,barrado01a,barrado02a}. A comparable result is obtained for the mass function in the solar neighbourhood \\citep{kroupa93,reid99a}.  Several theories have recently emerged to explain the formation of BDs. First, the picture of the turbulent fragmentation of molecular clouds can be extended to lower masses \\citep{klessen01,padoan02}. Second, \\citet{whitworth04} proposed that BDs could result from the erosion of pre-stellar cores in OB associations. Third, gravitational instabilities of self gravitating protostellar disks might also be responsible for the formation of BDs \\citep{watkins98a,watkins98b,lin98,boss00}. Moreover, BDs could stop accreting gas from the molecular cloud due to an early ejection from a multiple system \\citep{reipurth01,bate02,delgado_donate03,sterzik03}. Finally, as BDs straddle the realms of stars and planets, they might form within a circumstellar disk in a similar manner to giant planets \\citep{papaloizou01,armitage02}.  Observational evidence for disks around young BDs has been reported in several star-forming regions in the near-infrared \\citep{muench02,wilking99,luhman99a,oliveira02}, in the L$'$-band at 3.8\\,$\\mu$m \\citep{liu03,jayawardhana03b}, in the mid-infrared \\citep{natta02,testi02,apai02}, and at millimetre wavelengths \\citep{klein03}. These results suggest a common formation mechanism for stars and BDs. The recent lack of substellar objects in Taurus compared to the Trapezium Cluster and IC\\,348 \\citep{briceno02,luhman03b} and the distinctive binary properties of field BDs \\citep{reid01b,burgasser03a,gizis03,close03,bouy03} hint that stars and BDs may represent two independent populations \\citep{kroupa03b}. Additional studies of nearby field BDs and young BDs in open clusters are necessary to pin down on the formation mechanism(s) of substellar objects.  We have initiated a large Canada-France-Hawaii Telescope Key Programme (CFHTKP) surveying 80 square degrees in star-forming regions ($\\leq$\\,3 Myr), pre-main-sequence clusters (10--50 Myr), and in the Hyades (700 Myr) in $I$ and $z$ filters with the CFH12K wide-field camera. The goal of our project is to investigate in an homogeneous manner the dependence of the IMF with time and environment as well as the distribution of stars and BDs in clusters to provide clues on their formation. Our detection and completeness limits are ($I$,$z$)\\,$\\sim$\\,24 and 22 mag in both passbands, respectively. In this paper, we present the results of a 1.6 square degree (5 CFH12K fields-of-view) survey complemented by near-infrared photometry in one young open cluster, Collinder 359 (= Melotte\\,186).  Collinder 359 is a young open cluster located in the Ophiuchus constellation around the B5 supergiant 67 Oph (HD\\,164353). The cluster was first mentioned by \\citet{melotte15} as a ``large scattered group of bright stars'' and its presence confirmed later by \\citet{collinder31}. A handful of cluster members are known \\citep{collinder31,vantveer80,rucinski87}. It lies between 200 pc (Lyng\\aa{} 1987) and 650 pc (Kharchenko 2004; personal communication) with a mean value of 435 pc from the HIPPARCOS parallax measurement \\citep{perryman97}. The age of \\Coll{} is estimated to be about 30 Myr \\citep{wielen71,abt83}.  This paper is structured as follows: A literature review of our current knowledge of Collinder 359 is presented in \\S\\ref{cr359:literature} including a discussion regarding its proper motion, distance and age. In \\S\\ref{cr359:CFHTKP}, we briefly introduce the framework of the CFHTKP. The 1.6 square degree, wide-field optical survey of Collinder 359 is detailed in \\S\\ref{cr359:obs_cr359}. A description of the candidate cluster member selection process, using an optical ($I-z$, $I$) colour-magnitude diagram, is given in \\S\\ref{cr359:CMCs}. Finally, in \\S\\ref{cr359:IRfollowup}, we present details of a near-infrared follow-up survey of optically-selected candidate members in Collinder 359\\@.  \\begin{table*} \\begin{center} \\caption[Latest status of the bright stars in Collinder 359 (before 2000)]{ This table lists the 13 bright stars within Collinder 359 as listed by \\citet{collinder31}. Column 1 lists the running number of the member, column 2 gives the Henry Draper Catalogue number, columns 3 and 4 list the right ascension and the declination (in J2000), column 5 lists the spectral types \\citep{collinder31}, columns 6, 7, and 8 lists the $V$ magnitude and the $B-V$ and $U-B$ from \\citet{blanco68}, columns 9 and 10 list the $V-R$ and $R-I$, columns 11 and 12 list the proper motion (arcsec per year) of the object according to the SAO catalogue (1966), column 13 gives the HIPPARCOS parallax $\\pi$ \\citep{perryman97}. The membership of the object is given on the last column according to the discussion between \\citet{rucinski80} and \\citet{vantveer80}. } \\begin{tabular}{ c c c c c c c c c c c c c c } \\hline \\hline \\multicolumn{1}{c}{N$^{\\circ}$} & \\multicolumn{1}{c}{HD} & \\multicolumn{1}{c}{RA} & \\multicolumn{1}{c}{Dec} & \\multicolumn{1}{c}{SpT} & \\multicolumn{1}{c}{$V$} & \\multicolumn{1}{c}{$B-V$} & \\multicolumn{1}{c}{$U-B$} & \\multicolumn{1}{c}{$V-R$} & \\multicolumn{1}{c}{$R-I$} & \\multicolumn{1}{c}{$\\mu_{\\alpha}$} & \\multicolumn{1}{c}{$\\mu_{\\delta}$} & \\multicolumn{1}{c}{$\\pi$} & \\multicolumn{1}{c}{M?} \\\\ \\hline 1 & 161868 & 17 47 53.5  & 02 42 26 & A0     & 3.74 &  $+$0.03 & $+$0.14  &    0.01  &    0.00 &  $-$0.0240 & $-$0.074 &   34.42$\\pm$0.99 & NM  \\\\ 2 & 164353 & 17 58 08.3  & 02 55 57 & B5\\,Ib & 3.96 &  $+$0.04 & $-$0.63  &    0.06  &    0.03 &  $-$0.0015 & $-$0.010 &    2.30$\\pm$0.77 & M   \\\\ 3 & 164577 & 18 01 45.2  & 01 18 18 & A2     & 4.43 &  $+$0.04 & $+$0.05  &    0.04  &    0.01 &  $+$0.0090 & $-$0.012 &   12.31$\\pm$0.83 & NM  \\\\ 4 & 164284 & 18 00 15.8  & 04 22 07 & B3     & 4.70 &  $-$0.04 & $-$0.86  &    0.10  &    0.08 &  $+$0.0000 & $-$0.013 &    4.82$\\pm$0.78 & NM  \\\\ 5 & 166233 & 18 09 33.8  & 03 59 35 & F2     & 5.72 &  $+$0.37 & $+$0.02  &    0.22  &    0.21 &  $+$0.0360 & $-$0.007 &   19.62$\\pm$1.22 & NM  \\\\ 6 & 165174 & 18 04 37.3  & 01 55 08 & B3     & 6.14 &  $-$0.01 & $-$0.98  &    0.03  &    0.03 &  $-$0.0045 & $-$0.003 & $-$0.76$\\pm$0.89 & NM  \\\\ 7 & 168797 & 18 21 28.4  & 05 26 08 & B5     & 6.16 &  $-$0.02 & $-$0.64  &    0.00  &    0.01 &  $+$0.0105 & $-$0.004 &    0.97$\\pm$0.83 & NM   \\\\ 8 & 164432 & 18 00 52.8  & 06 16 05 & B3     & 6.35 &  $-$0.08 & $-$0.77  & $-$0.01  & $-$0.01 &  $+$0.0015 & $-$0.003 &    2.17$\\pm$0.87 & M   \\\\ 9 & 163346 & 17 55 37.5  & 02 04 29 & A3     & 6.78 &  $+$0.56 & $+$0.36  &    0.37  &    0.40 &  $-$0.0030 & $+$0.007 &    5.07$\\pm$0.87 & NM   \\\\ 10 & 164097 & 17 59 29.5  & 02 20 37 & A2     & 8.54 &  $+$0.17 & $+$0.15  &    0.12  &    0.15 &  $-$0.0060 & $+$0.003 &         ---      & M    \\\\ 11 & 164283 & 17 57 42.4  & 05 32 37 & A0     & 9.10 &  $+$0.26 & $+$0.19  &    0.16  &    0.21 &  $+$0.0075 & $-$0.014 &    4.11$\\pm$1.27 & M   \\\\ 12 & 164352 & 18 00 41.7  & 03 08 57 & B8     & 9.33 &  $-$0.01 & $-$0.39  &    0.02  &    0.06 &  $-$0.0015 & $-$0.002 &         ---      & M   \\\\ 13 & 164096 & 17 59 34.6  & 02 30 16 & A2     & 9.70 &  $+$0.20 & $+$0.17  &    0.13  &    0.20 &  $-$0.0105 & $-$0.006 &         ---      & M   \\\\ \\hline \\end{tabular} \\label{tab_cr359:old_members} \\end{center} \\end{table*} ",
        "conclusions": "\\label{cr359:summary}  We have presented the first deep optical wide-field imaging survey complemented with near-infrared photometric observations of the young open cluster \\Coll{}. We have surveyed a 1.6 square degree area in the cluster in the $I$ and $z$ filters down to detection and completeness limits of 22.0 and 24.0 mag with the CFH12K wide-field camera on the Canada-France-Hawaii 3.6-m telescope. Based on their location in the optical ($I-z$,$I$) \\CMD{}, we have extracted a total of 506 cluster member candidates in \\Coll{} spanning 1.3--0.030\\,\\Msun{}, assuming a distance of 450 pc and an age of 60 Myr. The uncertainties on the distance and age are 200 pc and 20 Myr, respectively. We have cross-correlated the optically-selected candidates with the 2MASS database for objects brighter than $I$ = 17.0 mag to weed out a proportion of contaminating field stars. Further $K'$-band photometry has been obtained for a subsample of 49 faint cluster candidates to probe the contamination at and below the stellar/substellar boundary.  By comparing the location of the brightest cluster member 67 Oph, with solar metallicity isochrones including moderate overshoot, we have derived an age of 60\\,$\\pm$\\,20 Myr for \\Coll{}. Taking into account the observed differences in ages of open clusters between the turn-off main-sequence method and the lithium test, \\Coll{} is likely older than \\APer{} but younger than the Pleiades. Thus, the expected age of the cluster is at least twice larger than previously thought. Finally, the comparison of the number of sources in a control field with the number of selected cluster candidates indicates that the surveyed fields do not contribute significantly to the cluster population. A new optical survey closer to the cluster centre but avoiding bright stars is needed to assess its extent and properties."
    },
    "0601/astro-ph0601346_arXiv.txt": {
        "abstract": "We identify galaxy groups and clusters in volume-limited samples of the SDSS redshift survey, using a redshift-space friends-of-friends algorithm.  We optimize the friends-of-friends linking lengths to recover galaxy systems that occupy the same dark matter halos, using a set of mock catalogs created by populating halos of N-body simulations with galaxies.  Extensive tests with these mock catalogs show that no combination of perpendicular and line-of-sight linking lengths is able to yield groups and clusters that simultaneously recover the true halo multiplicity function, projected size distribution, and velocity dispersion.  We adopt a linking length combination that yields, for galaxy groups with ten or more members: a group multiplicity function that is unbiased with respect to the true halo multiplicity function; an unbiased median relation between the multiplicities of groups and their associated halos; a spurious group fraction of less than $\\sim 1\\%$; a halo completeness of more than $\\sim 97\\%$; the correct projected size distribution as a function of multiplicity; and a velocity dispersion distribution that is $\\sim 20\\%$ too low at all multiplicities.  These results hold over a range of mock catalogs that use different input recipes of populating halos with galaxies.  We apply our group-finding algorithm to the SDSS data and obtain three group and cluster catalogs for three volume-limited samples that cover 3495.1 square degrees on the sky, go out to redshifts of 0.1, 0.068, and 0.045, and contain 57138, 37820, and 18895 galaxies, respectively.  We correct for incompleteness caused by fiber collisions and survey edges, and obtain measurements of the group multiplicity function, with errors calculated from realistic mock catalogs. These multiplicity function measurements provide a key constraint on the relation between galaxy populations and dark matter halos. ",
        "introduction": "\\label{intro}  Galaxies are gregarious by nature.  Bright galaxies typically reside in groups or clusters, surrounded by less luminous neighbors.  Interactions within the group or cluster environment may have important effects on the star formation history, morphology, dynamics, and other properties of member galaxies. Characterizing the relation between galaxy properties and their group environment is thus a key step in understanding galaxy formation and evolution. At the density thresholds often used to identify groups, most members should belong to the same, gravitationally bound dark matter (DM) halo.\\footnote{Throughout this paper, we use the term ``halo'' to refer to a gravitationally bound structure with overdensity $\\rho/\\bar{\\rho}\\sim200$, so an occupied halo may host a single luminous galaxy, a group of galaxies, or a cluster.  Higher overdensity concentrations around individual galaxies of a group or cluster constitute, in this terminology, halo substructure, or ``sub-halos''.}  Recent approaches to describing the relation between galaxies and DM focus on galaxy populations of DM halos as a function of halo mass.  Specifically, the bias of a particular class of galaxies can be characterized by its Halo Occupation Distribution (HOD), which specifies the probability distribution $P(N|M)$ that a halo of mass $M$ contains $N$ such galaxies, together with relations describing the relative spatial and velocity distributions of galaxies and dark matter within halos (\\citealt{berlind_weinberg_02} and references therein).  A well defined group catalog with well understood properties can play a central role in the empirical determination of this relation.  This paper presents a group and cluster catalog defined from the Sloan Digital Sky Survey (SDSS, \\citealt{york_etal_00}).  While this catalog is useful for many purposes, our overriding objective is to obtain a well understood measurement of the group multiplicity function (the space density of groups as a function of richness), with the goal of determining the HOD in the high mass regime \\citep{peacock_smith_00,berlind_weinberg_02,marinoni_hudson_02,kochanek_etal_03,lin_etal_04}. With this objective in mind, we have adopted a simple group-finding algorithm, friends-of-friends in redshift space \\citep{huchra_geller_82}, and carried out extensive tests on realistic mock catalogs in order to assess its performance and optimize parameter choices.  We apply the group-finding algorithm to volume-limited samples of galaxies so that the resulting group statistics characterize the clustering of well defined populations of galaxies.  Galaxy clusters have been the focus of study since they were first seen on optical photographic plates \\citep{shapley_ames_26}.  \\citet{zwicky_37} pioneered the study of clusters as dynamical objects by using imaging and spectroscopy of the Coma cluster to estimate its mass.  However, the most influential pioneering work on clusters was done by \\citet{abell_58}, who assembled the first large sample of galaxy clusters.  The Abell catalog of rich galaxy clusters \\citep{abell_58,abell_etal_89} was created by eyeball identification in the Palomar Observatory Sky Survey and it spawned numerous follow-up studies.  \\citet{devaucouleurs_71} shifted focus to poorer systems by studying nearby groups of galaxies.  \\citet{gott_turner_77b} made the first measurement of the group multiplicity function using the \\citep{turner_gott_76} catalog of groups selected based on the projected surface density of galaxies.  With the advent of large redshift surveys, group identification became three dimensional and thus less subject to projection effects.  Group-finding in redshift space was pioneered by \\citet{huchra_geller_82} and \\citet{geller_huchra_83}, using the Center for Astrophysics (CfA) redshift survey.  Subsequent versions of the CfA redshift survey were used to identify groups by various authors \\citep{nolthenius_white_87,ramella_etal_89,moore_etal_93,ramella_etal_97}. Other redshift surveys that spawned group catalogs were the Nearby Galaxies Catalog \\citep{tully_87}, the ESO Slice Project \\citep{ramella_etal_99}, the Las Campanas Redshift Survey (LCRS) \\citep{tucker_etal_00}, the Nearby Optical Galaxy Sample (NOG) \\citep{giuricin_etal_00}, the Southern Sky Redshift Survey (SSRS) \\citep{ramella_etal_02}, the 2dF redshift survey \\citep{merchan_zandivarez_02,eke_etal_04a,yang_etal_05}, and even the high redshift DEEP2 survey \\citep{gerke_etal_05}.  There have been several efforts to detect clusters in the SDSS to date, most of them using the photometric data rather than the redshift data.  \\citet{annis_etal_99} developed the maxBCG technique, where Brightest Cluster Galaxy (BCG) candidates are identified based on their colors and magnitudes and other cluster members are selected from nearby galaxies that have the colors of the E/S0 ridgeline. \\citet{kim_etal_02} developed a hybrid matched filter (HMF) technique that assumes a radial profile for clusters and convolves the data with that filter. \\citet{goto_etal_02} developed the cut-and-enhance (CE) method, which selects overdensities of galaxies that have similar colors.  All these techniques were applied to the early SDSS commissioning data \\citep{bahcall_etal_03,goto_etal_02}. \\citet{lee_etal_04} identified compact groups by looking for small and isolated concentrations of galaxies in the SDSS Early Data Release (EDR; \\citealt{stoughton_etal_02}).  Cluster searches in the SDSS redshift survey have also been carried out.  \\citet{goto_etal_05} used a friends-of-friends algorithm (though with linking lengths that do not scale with the changing number density of galaxies due to the flux limit) to identify clusters in the SDSS Data Release 2 (DR2; \\citealt{abazajian_etal_04}).  \\citet{merchan_zandivarez_05} used a friends-of-friends algorithm to identify groups in the SDSS Data Release 3 (DR3; \\citealt{abazajian_etal_05}).  \\citet{weinmann_etal_05} used the \\citet{yang_etal_05} algorithm to identify groups in SDSS DR2.  \\citet{miller_etal_05} developed the C4 algorithm for finding clusters in redshift space and also applied it to the SDSS DR2. The C4 algorithm looks for concentrations of galaxies in a seven-dimensional position and color space.  It takes advantage of the color similarity of cluster member galaxies and thus minimizes contamination due to projection.  However, some correlations are built into the method, and modeling it in order to understand the properties of the resulting cluster catalog requires a complete model of the galaxy population (including colors and luminosities).  Our method complements the C4 catalog by applying a simple and easily modeled algorithm to volume-limited samples with homogeneous properties.  In \\S~\\ref{data} we describe the SDSS data that we use.  In \\S~\\ref{mocks} we describe the mock catalogs that we use to optimize our group-finder and to estimate uncertainties for our measured group statistics.  In \\S~\\ref{groupfinder} we outline our group-finding algorithm and choice of parameters.  We present a detailed discussion of tests with mock catalogs in the Appendix, with the key points summarized in the main text. We discuss incompleteness in our group catalogs due to fiber collisions and survey edges in \\S~\\ref{incompleteness}.  The group catalogs are published in electronic tables and their contents are described in \\S~\\ref{catalog}.  Finally, in \\S~\\ref{multiplicity}, we present our measured group multiplicity function.  We will use this to constrain the HOD in future work.  We summarize our results in \\S~\\ref{summary}. ",
        "conclusions": "\\label{summary}  We have used a simple friends-of-friends algorithm to identify galaxy groups in volume-limited samples of the SDSS redshift survey.  We have selected FoF linking lengths that are best at grouping together galaxies that occupy the same dark matter halos.  We based this choice on extensive tests with mock galaxy catalogs, which we constructed by populating halos in N-body simulations with galaxies.  The result of our mock tests is that no combination of perpendicular and line-of-sight linking lengths can yield groups that successfully recover all aspects of the parent halo distribution, even for large richness systems. Specifically, FoF cannot identify groups that simultaneously have unbiased abundances, projected sizes, and velocity dispersions.  The ideal group-finding parameters for a given study depend on its scientific objectives.  Given our objective of using the multiplicity function to constrain the HOD, it makes sense to sacrifice velocity dispersions and obtain groups with unbiased abundances and projected sizes.  Our choice of linking lengths results in a group catalog that, for groups of ten or more members, has an unbiased multiplicity function, an unbiased median relation between the multiplicities of groups and their parent halos, an unbiased projected size distribution as a function of multiplicity, and a velocity dispersion distribution that is $\\sim 20\\%$ too low for all multiplicities.  We correct for fiber collisions and survey edge effects and present three SDSS group catalogs (for three different volume-limited samples) and their measured multiplicity functions.  It is important to recognize that our adopted group finder has the above properties only for halos defined using FoF with a linking length of 0.2 times the mean interparticle separation, since this is how halos were identified in our mock catalogs.  A different halo definition (such as FoF with a different linking length, or spherical overdensity halos) would require a different set of optimal group-finding parameters.  This is not a problem as long as the same halo definition is used consistently.  For example, an HOD measured from these group catalogs will hold for this halo definition, and any theoretical model should use the same halo definition to compare its predictions to the measured HOD. We chose this particular halo finder because it has been widely used and tested, and the properties of the resulting halo distribution (e.g., mass function) are well understood.  The groups and clusters that we present here are intended to be systems of galaxies that belong to the same virialized dark matter halo.  We can test whether these systems are virialized by computing crossing times for the groups and checking if they are sufficiently less than the Hubble time.  We define the crossing time divided by the hubble time as \\begin{equation} \\frac{\\tcross}{\\tH} = \\frac{(\\Rrms/\\hmpc)}{(\\sigv/100\\kms)}, \\end{equation} where $\\Rrms$ is the one-dimensional group radius, which is equal to the projected (two-dimensional) radius, $\\Rproj$, divided by the square root of two.  We correct for the velocity dispersion bias revealed in our mock tests by applying a 20\\% upward correction to all group velocity dispersions, and we compute $\\tcross/\\tH$ for all groups.  We find that, for all three group catalogs, the median value of $\\tcross/\\tH$ is $\\sim0.15$, and 80\\% of all groups have values less than $\\sim0.29$. These numbers can be interpreted in terms of the spherical infall model \\citep{gunn_gott_72,gott_turner_77a}, or other analytic or numerical models. However, at a first glance, the numbers are encouraging and suggest that most of our groups are likely virialized systems.  The group and cluster catalogs presented here are well-suited for testing many of the predictions and assumptions made by galaxy formation models regarding the relationship between galaxies and their underlying dark matter halos.  We will investigate several of these issues in subsequent papers."
    },
    "0601/nucl-ex0601035_arXiv.txt": {
        "abstract": "Recent  Wilkinson Microwave Anisotropy Probe (WMAP) measurements have determined the baryon density of the Universe $\\Omega_b$ with a precision of about 4\\%\\@. With $\\Omega_b$ tightly constrained, comparisons of Big Bang Nucleosynthesis (BBN) abundance predictions to primordial abundance observations can be made and used to test BBN models and/or to further constrain abundances of isotopes with weak observational limits. To push the limits and improve constraints on BBN models, uncertainties in key nuclear reaction rates must be minimized. To this end, we made new precise measurements of the \\ddp and \\ddn total cross sections at lab energies from 110~keV to 650~keV\\@. A complete fit was performed in energy and angle to both angular distribution and normalization data for both reactions simultaneously.  By including parameters for experimental variables in the fit,  error correlations between detectors, reactions, and reaction energies were accurately tabulated by computational methods.  With uncertainties around $\\rm 2\\% \\pm 1\\%$ scale error,  these new measurements significantly improve on the existing data set.  At relevant temperatures, using the data of the present work, both reaction rates  are found to be about 7\\% higher than those in the widely used Nuclear Astrophysics Compilation of Reaction Rates (NACRE).  These data will thus lead not only to reduced uncertainties, but also to modifications in the BBN abundance predictions. ",
        "introduction": "\\label{chap:intro}  The standard Big Bang model explains remarkably well many features of the Universe which are otherwise difficult to reconcile. The standard Big-Bang nucleosynthesis (BBN) model consists of a small network of nuclear reactions occuring at energies easily obtained in the lab. The outcome of Standard BBN is determined almost entirely by the nuclear reaction rates and the baryon to photon ratio, $\\eta$, of the Universe, which is directly related to the baryon density $\\Omega_b$ ($\\Omega_bh^2=3.66\\times10^7 \\eta$, where $h$ is the Hubble constant in units of $\\rm 100 km\\times s^{-1}\\times Mpc^{-1}$~\\cite{PDG} ).  With detailed network calculations and knowledge of the nuclear cross-sections, all of the primordial abundances can be precisely calculated as a function of $\\eta$. Knowledge of any one abundance or a separate determination of $\\eta$ thus allows all other primordial abundances to be inferred from the Standard BBN model.  Knowledge of any two of these observables produces a check of the model itself.  In the past, quantitative understanding of BBN has been limited by uncertainties in the observed primordial abundances and the value of $\\eta$.  Until recently $\\eta$ was treated as a free parameter.  As observations have improved, the value of $\\eta$ has become a well determined input and is instead used, along with the nuclear reaction rates, to predict the primordial abundances.  Regardless of which values are assumed and which are predicted, the uncertainties of the cross sections themselves are becoming a significant factor in precision tests of BBN models.  Improved analyses depend on having accurate and precise knowledge of the nuclear reaction rates, their uncertainties, and the correlations in those uncertainties.  \\subsection{Primordial Deuterium Observations} Recent measurements of absorption lines in high red-shift, metal poor, QSO-back-lit gas clouds have constrained the primordial deuterium abundance $D$ to the impressive interval of ${D/H}=2.78_{-.38}^{+.44}\\times10^{-5}$~\\cite{kirkman}, expressed relative to the hydrogen abundance $H$\\@.  The measurement and analysis procedure is well described in the review article of Tytler {\\it et al.}~\\cite{Tyt00}\\@. Currently systematic scatter limits the precision of the deuterium abundance observations~\\cite{kirkman}, but as more data arrive and the systematics become better understood this could quickly change.  Measurements of primordial deuterium abundances \\cite{bur98a,bur98b}, along with recent measurements of $\\eta$ from WMAP data~\\cite{cyburt,spergelwmap}, bring the nuclear reaction rates increasingly closer to being the limiting factors in testing the consistency of the Standard BBN model.    \\subsection{Cosmology Enters the Lab}  The value of $\\eta$ determined from the WMAP results is in good agreement with primordial deuterium abundance measurements of Kirkman~{\\it et al.}~\\cite{kirkman}\\@.  However limited amounts of data and the systematic uncertainties in abundance measurements and in reaction rates have limited the level of precision at which the BBN models can be tested to around 10\\% for most observables.  With the uncertainties of previously existing data, the \\ddp and \\ddn reaction cross-sections at energies in the range of a few hundred keV make large contributions to the uncertainties in deuterium abundances as predicted by network calculations~\\cite{nollettburles,bur99a,bur99b,sch98}\\@.  Nollett and Burles~\\cite{nollettburles} provide sensitivity functions estimating contributions to the deuterium and $\\rm ^7Li$ abundances of several reactions as a function of energy.  The most relevant energy range of both the \\ddn and \\ddp reactions extends from roughly $\\rm {\\it E_d}=100~keV$ up to about $\\rm {\\it E_d}=700~keV$.  Prior to our measurements, high precision data for these reactions were very limited in this range, as shown in Fig.~\\ref{fig:ddndata}\\@.    The \\ddp reaction data is qualitatively very similar.   From $\\rm {\\it E_d}=325~keV$ to the top of the BBN energy range there were very few data points, all having large uncertainties.  Even at lower energies of significance to BBN, the uncertainties of previously existing data were around the 10\\% level or only slightly better.   For the energy range above 325~keV, one of the more complete data sets was that of Ganeev {\\it et al.}~\\cite{ganeev}\\@.   These data are not included in the NACRE data compilation~\\cite{NACRE},  considered the most prominent collection of experimental rates for reactions of astrophysical significance.  \\begin{figure}[htb]\\centering \\includegraphics[width=\\figwidth]{ddndata.eps} \\caption[Prominent data sets for \\ddn cross-sections at BBN energies.]{(Color) Prominant data sets for \\ddn cross-sections at BBN energies} \\label{fig:ddndata} \\end{figure}  In order to aid in tightening the BBN constraints, we measured total cross sections for both of these deuterium burning reactions at lab energies ranging from about 112~keV to 646~keV\\@.  By carefully establishing experimental procedures in order to maximally cancel the dependencies of yields on experimental parameters, we have obtained about 2\\% statistical uncertainties plus or minus a 1\\% systematic scale uncertainty.  Furthermore our $\\chi^2$ analysis techniques have allowed optimal use of the data and  provided correlations in the uncertainties of the two reactions across the range of energies.  The new data form a significant improvement of the inputs to the BBN calculations and facilitate an emerging era of high-precision BBN. ",
        "conclusions": "\\label{chap:results}  Here we present our data and address the effect these data will have on Big-Bang Nucleosynthesis calculations. The cross-section results obtained from the parameterization described in Eq.~\\ref{eq:discretepar} are given in Table~\\ref{tab:totalsfact}\\@.   The coefficients of the Legendre expansions of the differential cross-sections are given in Table~\\ref{tab:dsigs}\\@.  The errors shown are 1$\\sigma$ equivalents which include statistical and background uncertainties, energy uncertainties discussed in Sec.~\\ref{sec:anal:error:zero}, fitting uncertainties described in Sec.~\\ref{sec:fitting_error}, and which account for scatter in the data sets used to derive the total cross sections as explained in Sec.~\\ref{sec:error:scatter}\\@.  The uncertainties are generally around  the 2\\% level except at the lowest energies where peak fitting and energy uncertainties become somewhat larger.   An overall scale error of 1\\% is estimated arising from uncertainties in the $p$--$d$ elastic-scattering cross section to which these data were normalized.  This normalization error is not included in the uncertainties.  Finally, the error matrix for the total cross-section parameters is given in Table~\\ref{tab:discmat}\\@.  The complete error matrix for the Legendre coefficients is too large to present here and is not relevant to our primary goal of providing BBN inputs.  The cross sections of Table~\\ref{tab:totalsfact} are plotted in Fig.~\\ref{fig:bothsigs}, along with the continuous parameterization mentioned in Sec.~\\ref{sec:enchilada} and with the recent cross section compilation of Cyburt~\\cite{cyburtprd04}\\@.  Some scatter within the uncertainties is detectable, which is expected since our statistical and systematic errors are of comparable magnitudes.    The continuous parameterization appears to be systematically somewhat higher than the results of the discrete fit.  This is not surprising since cross sections for the lowest four energies are all normalized directly to the differential cross sections at 480~keV\\@. These data thus inherit all of the uncertainty of the 480~keV points, but in a systematic manner.  The continuous fit is then free to systematically renormalize these points within constraints.  This systematic uncertainty is quantified in the error matrix of the discrete fit.  The Cyburt compilation agrees well with the present data except at the high-energy end  the \\ddp reaction cross sections.  At the lowest energy we can compare directly to the high-precision data of Brown and Jarmie~\\cite{browndd}, which falls 11\\% and 8\\% below the present work for \\ddp and \\ddn respectively, corresponding to discrepancies of 2.1 $\\sigma$ and  1.6 $\\sigma$ respectively when systematic and statistical uncertainties of both experiments are included.  The data of Greife {\\it et al.}~\\cite{greife}, having uncertainties of only about 2.8\\%, agree well with our results at higher energies, up to 256~keV, for both reactions.  There have been several published compilations and analyses of BBN data and network calculations and we cannot review or even acknowledge all of the major efforts here.    Most, including Refs.~\\cite{cuoco,cyburt,redux,Des05}, have focused a significant amount of attention on consistency of cosmological observables including primordial abundances and CMBR observations.  Many, including some which have now acquired benchmark status, have used Monte Carlo techniques to analyze the relationship between the uncertainties in the reaction data to the uncertainties in the astrophysical constraints~\\cite{skm,romanelli,nollettburles}\\@.  The work of the latter produced estimates of energy regions having the highest sensitivities to cross-section data, estimates which inspired the present work.  Finally there are even some works that have used traditional first-order error propagation to understand the effects of the data on network predictions.  For example, Fiorentini {\\it et al.}~\\cite{fior} achieved impressively good agreement with the computationally expensive techniques used elsewhere and have the advantage of producing simple functional understandings of various uncertainty relationships via tabulated derivatives.  The NACRE compilation~\\cite{NACRE} has become widely accepted as a standard in adopted values for reaction-rates of nuclear reactions of astrophysical significance.  NACRE does not address network calculations since it is intended as a broad resource for general astrophysical use.  It has though become the data source of choice for several BBN analyses~\\cite{cuoco,flam,Coc04}\\@.  The NACRE reaction-rate values cannot be compared directly to the cross-section data of the present work.  The reaction rate is an integral of the total cross section weighted by a Maxwellian distribution and taken over all energies.  In order to calculate the reaction-rate integral for the purpose of this comparison, we used the cross section results from the continuous parameterization of the present work. For energies given in MeV and cross sections in barns, the reaction rate \\rate in $\\rm cm^3mol^{-1}s^{-1}$ is given by \\begin{equation} \\rate=3.731\\times10^{10}\\mu^{-1/2}T_9^{-3/2}\\int_0^\\infty\\sigma E\\,e^{-11.604\\frac{E}{T_9}}dE\\,, \\label{eq:rate} \\end{equation} where $\\mu$ is the reduced mass in amu and $T_9$ is the temperature in units of $\\rm 10^9 K$.  For the integral to converge at temperatures relevant to BBN, $T_9<2$, it can be cut off at $E_{\\rm c.m.}$ of about 2.5 MeV but not much lower.  Thus to calculate reaction rates accurately, we must include data at energies higher than those of the present work, although data at these energies contribute relatively little to the integral at the relevant temperatures.  For this purpose and to make a fair judgment of how the new data compare to the NACRE compilation, we have used the same data at high energies as those used in NACRE compilation, the data of Schulte {\\it et al.}~\\cite{schulte}\\@.  We have included this data up to $ E_{\\rm c.m.}=2.75~{\\rm MeV}$.  These data were included in the continuous parameterization by using the total cross sections to appropriately constrain the zero order Legendre coefficients at the high energies.   The data of Ref.~\\cite{schulte} have scatter which is significantly larger than their quoted statistical errors.  To prevent this data set from unduly constraining the fit, we have multiplied their statistical uncertainties by 5, leaving the smallest uncertainties in these data still below one percent of the value.  This multiplication is a coarse procedure but suffices for the present purpose. The results of these fits are shown in Figs.~\\ref{fig:schultfit} and ~\\ref{fig:heschultfit}\\@.  The \\ddn reaction curve does not follow the Schulte data impressively well.  This is probably due to a limitation of the parameterization in representing changes in curvature over these large energy regions.  It is not a significant problem for the purpose of constraining the tail of the reaction-rate integral.  The NACRE compilation gives coefficients for a quadratic cross-section function which was used in calculating the rates for the \\ddn and \\ddp reactions at ''low energies''\\@.  NACRE does not specify exactly what is meant by low energies.  We have used these cross-section functions for energies below the present work in order to complete the integrals. As was the case for the high energy data, this low energy data contributes only a small amount to the integral for temperatures significant to BBN.  Using these total cross-section curves, we then calculated the reaction rates by computing the right hand side of Eq.~\\ref{eq:rate} with a high-energy cutoff on the integral at $ E_{\\rm c.m.}=2.5{\\rm MeV}$.  The results are compared to the NACRE reaction rates in Figs.~\\ref{fig:ratecomp} and~\\ref{fig:heratecomp}\\@.   Judging by plots of abundances as a function of temperature for network calculations of Nollett and Burles~\\cite{nollettburles}, all significant standard BBN occurs at temperatures near $T_9=1$.  At this temperature the rates derived using the present work are about 7\\% higher than the NACRE rates for  both \\ddn and \\ddp\\@.  These discrepancies are systematic, remaining at similar levels for a large range of temperatures.  Our derived \\ddp rate is well beyond the NACRE quoted upper bounds, and for \\ddn our value is marginally within the NACRE bounds.  It is important to emphasize that there is nothing inherently controversial about disagreement between the present work and NACRE or any other compilation.  The NACRE curves are derived from a limited set of data and in fact include no data between $E_d=325$~keV and $ E_d=2.0~{\\rm MeV}$.  Furthermore, the data of Krauss {\\it et al.}\\@.~\\cite{krauss} were the only data used in the  NACRE compilation in the energy range of 250 to 325 keV, and although these are respectable data, they have total uncertainties of about 8\\% (see  Fig.~\\ref{fig:ddndata}).   The NACRE compilation incorporates all data in these plots with the exception of the Ganeev data.   Our good agreement with Cyburt~\\cite{cyburtprd04} is likely due to his inclusion of the Ganeev data set which, although it has large uncertainties, also agrees well with our data and covers the energy range with the least data available.  There was indeed a previous shortage of data and our results will greatly contribute to the cross-section information available in this energy region.  Many recent compilations such as that of Descouvemont {\\it et al.}~\\cite{Des04} used in the analysis of Coc {\\it et al.}~\\cite{Coc04}, as well as the compilation described by Serpico {\\it et al.}~\\cite{Ser04}, have a focus on BBN applications and strive to improve on the extrapolation techniques of the NACRE compilation in the BBN energy regions.  However, these compilations suffer from a lack of data in the same energy region as NACRE.    In spite of this, Serpico {\\it et al.}  quote unusually low uncertainties at the levels of 1.3\\% and 1\\% for the \\ddn and \\ddp reaction rates respectively.  They do however acknowledge the severe lack of data and \"strongly recommend a new experimental campaign\" which we have now furthered.  Using the logarithmic derivatives tabulated in Ref.~\\cite{fior}, we can estimate the differences in primordial abundances calculated using our new data vs. using the NACRE compilation.   These derivatives of abundances with respect to reaction rates are given in tables of coefficients of polynomial expansions in the baryon to photon ratio $\\eta$.  For our estimates we use a value of $\\eta$ of $6.14 \\times 10^{10}$ from Ref.~\\cite{cyburt} which was derived directly from an analysis of recent WMAP data performed in Ref.~\\cite{spergelwmap}\\@.  The values of the logarithmic derivatives calculated for $\\eta=6.14 \\times 10^{10}$ are given in Table~\\ref{tab:logder}\\@.  It is immediately evident that a roughly 7\\% change in the cross sections will have essentially no effect on the primordial abundance of \\hethree\\@.   The relative signs of the derivatives for the two reactions have important significance.  Although the derivatives shown for the $\\rm ^3He$ abundance are of significant magnitude, the changes in the abundance induced by the changes in the two reactions nearly cancel and the net effect is small.   However, because the two $D/H$ derivatives have the same sign as do the reaction rate changes inferred from the present work, the total change estimated in the deuterium abundance is roughly $-7$\\%\\@.  The change calculated for $\\rm ^7Li/{\\it H}$ is +5\\% and is almost entirely attributable to the changes in the \\ddn reaction rates.  The lithium abundance observations are currently in severe disagreement with the Standard BBN results~\\cite{kirkman}\\@.  This 5\\% increase in the predicted lithium abundance makes an unresolved problem slightly worse.  Currently the $\\rm ^7Li$ abundances predicted using CMBR values of $\\eta$ are at least a factor of 2 higher than the observed abundances~\\cite{cuoco}\\@.  The 5\\% change is small compared with the uncertainties plaguing this comparison but it certainly does not improve the hopes of finding agreement between standard BBN and observed $\\rm ^7Li$ abundances.  We look at this not as a failure but as an opportunity to explore possibilities of non-standard Big Bang Models and thus new physics.  The current observational value of the primordial deuterium abundance, $ {D/H}=2.78_{-0.38}^{+0.44}\\times 10^{-5}$~\\cite{kirkman}, obtained from QSO absorption spectra, is in good agreement with predictions of BBN using $\\eta$ from the CMBR data, ${D/H}=2.56_{-0.24}^{+0.35}\\times 10^{-5}$~\\cite{cuoco} or $ {D/H}=2.55_{-0.20}^{+0.21}\\times 10^{-5}$~\\cite{cyburtprd04}\\@.  However, as more data have arrived, the statistical uncertainty estimates on the abundance observations have not held up; the uncertainty is now dominated by scatter in the data.  The change which we predict in the $D/H$ value derived from BBN+CMBR will strain the current agreement slightly, but the results should still be well within the current mutual uncertainties.  The  reduced uncertainties of the \\ddn and \\ddp cross-sections from the present work, future improvements in precision of the \\dpg cross-section measurements, and more QSO observations may soon provide significantly more stringent tests of this comparison, again yielding insight into the validity of the details of the standard BBN model.  Although other aspects of the BBN picture are still limiting the comparisons of theory and experiment to around the level of 10\\% or worse, the new data of the present work will pave the way for future developments in precision cosmology.  As other measurements are improved, BBN predictions will become sensitive to many details of the model including possible inhomogeneities and neutrino properties.   The new data should put to rest concerns and claims that the $d$--$d$ reaction rates may be inaccurate, and should ultimately result in significant modifications to the reaction rates used for current astrophysical applications including BBN\\@.  The present data verify the limited, not always trusted, and often overlooked data sets which previously existed in this region of energy~\\cite{ganeev} and with roughly 2\\% to 3\\% uncertainties, this work represents a significant improvement in precision, which along with improvements in other cross-sections, will translate directly into reductions in uncertainties of BBN predictions.  These data will truly help to usher in the ''new era''~\\cite{cyburt} of precision cosmology."
    },
    "0601/astro-ph0601616_arXiv.txt": {
        "abstract": "{The chemically peculiar HgMn stars are a class of Bp stars which have historically been found to be both non-magnetic and non-variable. Remarkably, it has recently been demonstrated that the bright, well-studied HgMn star $\\alpha$~And exhibits clear Hg~{\\sc ii} line profile variations indicative of a non-uniform surface distribution of this element.}{With this work, we have conducted an extensive search for magnetic fields in the photosphere of $\\alpha$~And.}{We have acquired new circular polarisation spectra with the MuSiCoS and ESPaDOnS spectropolarimeters. We have also obtained FORS1 circular polarisation spectra from the ESO Archive, and considered all previously published magnetic data. This extensive dataset has been used to systematically test for the presence of magnetic fields in the photosphere of $\\alpha$~And. We have also examined the high-resolution spectra for line profile variability.} {The polarimetric and magnetic data provide no convincing evidence for photospheric magnetic fields. The highest-S/N phase- and velocity-resolved Stokes $V$ profiles, obtained with ESPaDOnS, allow us to place a $3\\sigma$ upper limit of about $100$~G on the possible presence of any undetected pure dipolar, quadrupolar or octupolar surface magnetic fields (and just 50~G for fields with significant obliquity). We also consider and dismiss the possible existence of more complex fossil and dynamo-generated fields, and discuss the implications of these results for explaining the non-uniform surface distribution of Hg. The very high-quality ESPaDOnS spectra have allowed us to confidently detect variability of Hg~{\\sc ii} $\\lambda 6149$, $\\lambda 5425$ and $\\lambda 5677$. The profile variability of the Hg~{\\sc ii} lines is strong, and similar to that of the Hg~{\\sc ii}\\ $\\lambda 3984$ line.  On the other hand, variability of other lines (e.g. Mn, Fe) is much weaker, and appears to be attributable to orbital modulation, continuum normalisation differences and weak, variable fringing. } {} ",
        "introduction": "Although Ap HgMn stars have generally been found to be both non-variable (e.g. Adelman 1998) and non-magnetic (Shorlin et al. 2002), Adelman et al. (2002) reported convincing evidence that the bright, well-established HgMn star $\\alpha$ Andromedae A (HD 358) exhibits clear ($>30$\\%) variability of the equivalent width of the Hg~{\\sc ii} $\\lambda$3984 line profile.  The Hg line variations observed by Adelman et al. (2002) were consistent with a period of $2.38236\\pm 0.00011$ days, a value in good agreement with the rigid rotation period $P=2.53\\pm 0.4$~days calculated using the measured $v\\sin i$ of $52\\pm 2$ km/s, the inferred rotational axis inclination $i=74^{\\rm o}$ and the stellar radius $R=2.7\\pm 0.4\\ R_\\odot$ (Ryabchikova et al. 1999).  Interpreting the Hg line variability as the consequence of rotational modulation of a nonuniform photospheric abundance distribution, Adelman et al. (2002) inverted the variations using the Doppler Imaging technique to recover a map of the Hg surface distribution. The derived distribution is characterised by a strong ($\\sim 2.5-4.5$ dex) enhancement of the Hg abundance in the {rotational} equatorial regions, from about $-30^{\\rm o}$ to $+30^{\\rm o}$ rotational latitude. The equatorial Hg enhancement appears to be structured into a series of three spots of very high abundance, at longitudes of $90\\degr$, $170\\degr$ and $270\\degr$, and connected by ``bridges'' of somewhat lower Hg abundance.  In the atmospheres of classical Ap stars, magnetic fields seem to be a necessary ingredient in the production of non-uniform surface abundance distributions (see, e.g. Auri\\`ere et al. 2004). Could the inferred non-uniform surface distribution of Hg imply that $\\alpha$~And is a magnetic HgMn star? Such a discovery would run counter to the conclusions of many studies (e.g. Borra \\& Landstreet 1980, Shorlin et al. 2002), and could lead to important new insights into the processes responsible for the production of atmospheric chemical peculiarities.  Longitudinal magnetic field observations of $\\alpha$~And have been obtained previously by Borra \\& Landstreet (1980, 5 measurements), Glagolevskij et al. (1985, 6 measurements), and Chountonov (2001, 2 measurements), with best formal error bars of about 40-50 G and with no detection of a magnetic field. Although all of these authors obtained measurements of the longitudinal magnetic field using circular polarisation, none obtained the Stokes $V$ spectrum, and none obtained a complete sampling of the 2.38-day rotational cycle.  Given the remarkable behaviour of $\\alpha$~And and the limitations of previous studies of its magnetic properties, we have decided to explore further the hypothesis that $\\alpha$~And is a magnetic HgMn star. In this paper we discuss available magnetic field measurements of this star, including new high-S/N circular spectropolarisation observations, well-distributed according to the rotational ephemeris of Adelman et al. (2002), along with a detailed examination of the optical line profiles and spectrum variability. ",
        "conclusions": "Until recently, $\\alpha$ And represented the only convincing example of a HgMn star showing intrinsic line profile variability. However, Kochukhov et al. (2005) have reported asymmetry and variability of the Hg~{\\sc ii} $\\lambda3984$ line in spectra of two additional HgMn stars, with temperatures similar to that of $\\alpha$~And. This appears to establish a new class of spotted, spectrum variable stars, of which $\\alpha$~And is the prototypical example.  From our analysis of the magnetic and polarisation measurements of $\\alpha$~And, we find that any oblique low-order multipolar magnetic field present in the photosphere must be {weaker} than about 100~G (for all field geometries), and weaker than about 50~G (for geometries with significant obliquity). Any undetected field is therefore constrained to be significantly weaker than the equipartition threshold at $\\log \\tau_{\\rm 5000}=0.0$. If we assume that a magnetic field of at least this order is required to support horizontal abundance nonuniformities { (e.g. by suppressing mixing due to meridional circulation)}, then {an oblique low-order multipolar magnetic field does not appear to be responsible for the nonuniform distributions of Hg and Mn of $\\alpha$ And,} at least in atmospheric layers around $\\log \\tau_{\\rm 5000}=0.0$. Observational support for this { assumption} is provided by Auri\\`ere et al. (2004), who have concluded that all magnetic Ap/Bp stars host dipolar magnetic field components that are {stronger} than the photospheric equipartition field (at $\\log \\tau_{\\rm 5000}=0.0$). This supports the assumption that a field of at least this strength is required for magnetic maintenance of photospheric chemical abundance non-uniformities.  Could it be that a dipolar magnetic field does in fact exist in the photosphere of $\\alpha$~And, but which is weaker than the local equipartition field in layers around $\\log \\tau_{\\rm 5000}=0.0$? Due to the exponential decrease of the gas density with height in the outer layers of the star, such a field could in fact be {stronger} than the equipartition field in atmospheric layers well above (say) $\\log \\tau_{\\rm 5000}=-1$. If the abundance non-uniformities were confined to these high layers, an undetected magnetic field (i.e. with polar intensity weaker than our 3$\\sigma$ upper limit) could in principle be responsible { for preventing horizontal mixing of the non-uniform abundance distributions.}  The concentration of Hg, Mn and other elements in the photospheres of HgMn stars in thin layers at small optical depths has been suggested based on both theoretical and observational grounds (e.g. Proffitt et al. 1999,  Sigut et al. 2000). Most recently, the discovery of emission lines in the optical spectra of some HgMn stars (e.g. Sigut et al. 2000, Wahlgren \\& Hubrig 2000) has provided additional evidence of stratification of these elements. Sigut (2001), interpreting the presence of Mn~{\\sc ii} emission lines in the spectra of HD 186122 and HD 179761 as the result of interlocked non-LTE effects combined with chemical stratification, found that Mn must be confined to column masses above $\\log \\Sigma=-1.5$, equivalent to about $\\log\\tau_{\\rm 5000}\\ltsim -1.1$. Although we observe no such emission lines in the spectrum of $\\alpha$~And, were Hg confined to very small optical depths in the atmosphere of this star (i.e. optical depths at which the equipartition field satisfies our field constraints - for $\\alpha$~And, $B_{\\rm eq}$ should be below 100~G above $\\log \\tau_{\\rm 5000}=-1.2$, and below 50~G above $\\log \\tau_{\\rm 5000}=-2.2$) it might be possible for an undetected magnetic field to { sustain} the lateral abundance non-uniformities.  In order to test this possibility, {we have performed some comparisons between the observed Hg line profiles and those predicted by a single-step stratified Hg distribution, in which Hg is 100 times more abundant above $\\log\\tau_{\\rm 5000}=-1.2$ than below, and in which the lower zone abundance { (we employed $\\epsilon_{\\rm Hg}=-5.0$)} has been chosen to approximately fit the profile of Hg~{\\sc ii}~$\\lambda 3984$. Unfortunately, the combination of line asymmetry and variability, coupled with the weakness of many of the Hg lines (1-2\\% deep) and local continuum normalisation errors, makes it impossible to realistically choose between the stratified and unstratified models. }    { Ultimately, we are led to conclude that models seeking to explain the nonuniform Hg surface distribution of $\\alpha$~And which invoke the presence of a magnetic field of $B\\gtrsim 100$~G are almost certainly falsified by the results presented in this paper. It may still be possible to implicate a dynamically-important magnetic field if Hg is strongly stratified, and confined primarily to upper layers of the atmosphere (where the equipartition field is weaker than the upper limits presented here). Based on our experiments, modeling of the vertical Hg distribution will require very high quality data (of quality similar to or better than that discussed here), coupled with a model which simultaneously considers the lateral and vertical abundance distribution. However, in the absence of evidence of strong Hg stratification, we conclude that { a non-magnetic} mechanism is almost certainly responsible for this phenomenon { (see e.g. Kochukhov et al. 2005).}}"
    },
    "0601/astro-ph0601085_arXiv.txt": {
        "abstract": "We analyze photometric properties of 1384 cluster galaxies as a function of the normalized distance to cluster center. These galaxies were selected in the central region ($r/r_{200} \\leq$ 0.8)  of 14 southern Abell clusters chosen from the Southern Abell Cluster Redshifts Survey (SARS). For 507 of these galaxies we also obtained their luminosity profiles. We have studied the morphology-clustercentric distance relation on the basis of the shape parameter $n$ of the S\\'ersic's law. We also have analyzed the presence of a possible segregation in magnitude for both, the galaxy total luminosity and that of their components (i.e. the bulge and the disk).  Results show a marginal ($2\\sigma$ level) decrease of the total luminosity as a function of normalized radius. However, when bulges are analyzed separately, a significant luminosity segregation is found ($3\\sigma$ and $2\\sigma$ for galaxies in projection and member galaxies respectively). The fraction of bulges brighter than $M_B \\leq -22$ is three times larger in the core of clusters than in the outer region.  Our analysis of the disk component suggests that disks are, on average, less luminous in the cluster core than at $r/r_{200} \\sim 0.8$. In addition, we found that the magnitude-size relation as a function of $r/r_{200}$ indicates (at $2\\sigma$ level) that disks are smaller and centrally brighter in the core of clusters. However, the Kormendy relation (the bulge magnitude-size relation) appears to be independent of environment. ",
        "introduction": "It is well known that environment affects galaxy properties as: morphology, luminosity, color, star formation rate, gas content and  structure of the subsystems. Different mechanisms have been proposed to explain how these properties can be affected by the environment. A large number of galaxies can be well represented by two major components, the bulge and the disk. These two components can be affected in many ways when the galaxy is moving into the environment of a rich cluster. Moore et al.(1998) follow the evolution of disks galaxies in a rich cluster (galaxy harassment) and find that the result of close encounters is a transformation from disks to spheroids. Fujita and Nagashima (1999) suggested that ram pressure stripping (Gunn and Gott 1972, Abadi et al. 1999) increases the bulge to disk luminosity ratio $B/D$ of normal spiral galaxies due to the suppression of the star formation and hence favoring the transformation into earlier Hubble types.  Luminosity segregation was detected by Rood \\& Turnrose (1968), Quintana (1979), Capelato et al. (1980), Yepes et al. (1991) and Kashikawa et al. (1998). Moreover, this effect was detected when galaxies are considered through their clustercentric distances or their velocity dispersions (most luminous galaxies have smaller velocity dispersion) (Rood et al. 1972, Biviano et al. 1992). However, luminosity segregation have also found opponents like Noonan (1961), Bahcall (1973) and Sarazin (1980), who suggested that evidences for luminosity segregation are spurious, and mostly due to poor background subtraction. However, his optimized fitting procedure was applied to the Coma cluster data with little background data.   In this paper we analyze the relations between galaxy structure and photometric parameters vs. cluster's environment. Dom\\'{\\i}nguez et al. (2001)  found that parameters defined as a function of the distance to the cluster centre are the most appropriated to represent the morphological segregation of galaxies in the inner relaxed region of nearby clusters. Since the present work concentrates on that region of clusters ($r<r/r_{200}$),  we analyze the galaxy properties as a function of the clustercentric distance normalized to $r_{200}$. Redshift confirmed members and galaxies seen in projection are analyzed separately. The goal of this paper is to quantify the environment dependence of the structural and photometric properties of galaxies in clusters. Our sample consists of 507 galaxies in 14 southern Abell clusters of the SARS sample (Way et al. 2005, Hearafter Paper I). The structural and photometric parameters were obtained in Coenda et al. (2005, Hereafter Paper III). The paper is structured as follows: observations and photometric analysis are described in section \\ref{photo}. In section \\ref{res} we derive and analyze our results, and the conclusions are given in section \\ref{concl}. ",
        "conclusions": "We have analyzed the correlation between galaxy photometric parameters and the normalized clustercentric radius $r/r_{200}$ for 507 galaxies in the central region of 14 Abell cluster. All the analysis performed in this work were applied to two different samples: i) all galaxies in projection, for which we have taken into account the standard corrections, and ii) redshift confirmed members corrected by completeness.  Based on the S\\'ersic index $n$ we analyzed the morphological-$r/r_{200}$ relation. We found that the $n$ parameter is a good alternative to measure the morphological segregation. The use of $n$ has the advantage that it is a continuous parameter that can be estimated in a more reproducible procedure.  In order to test for a possible luminosity segregation, we analyzed the correlation between the fraction of galaxies with $M_{t}\\le-21$ and $r/r_{200}$. Our results show a marginal (two sigma level) decrease of the total luminosity as the normalized radius increases. It should be noted that this effect is only present when redshift-confirmed members are considered. The same analysis was repeated for bulge and disk sub-systems. Our results indicate a segregation in bulge-magnitude when both confirmed members galaxies ($2\\sigma$ level) and projected galaxies ($3\\sigma$ level) are used. We found that the fraction of bulges brighter than  $M_{t}\\le-22$ is approximately three times larger in the inner cluster region than in the outer cluster region. This analysis was performed for a total-magnitude-complete sample and for a bulge-magnitude-complete sub-sample. In both cases, our results are consistent with a segregation in the bulge luminosity.  On the other hand, the absolute magnitude of disks presents a dependence on $r/r_{200}$ in the sense that disks tend to have on average lower luminosities as they are closer to the core of the parent cluster of galaxies. It should be noted that this effect is only present when redshift-confirmed members are considered.  If disk galaxies are selected in projection, the effect is only statistically significant for the nearest clusters. Disk luminosity segregation support the idea that disks are tidally affected by the cluster potential or by high speed encounters with other cluster member galaxies. However, there are other phenomena that could affect star formation in galaxy disks such as ram pressure, gas evaporation, or lack of gas accretion from the intra-cluster medium. All these mechanisms can also be responsible for the observed disk luminosity segregation.  We have analyzed the scaling relations for bulges and disks as a function of $r/r_{200}$. We did not find any statistically significant dependence of the $M_B-log(r_e)$ relation with the clustercentric distance. Our results suggest that the physical conditions responsible for the Kormendy relation are sufficiently robust to support the extreme conditions found in the core of the clusters of galaxies. On the other hand, the $M_D-log(r_0)$ relation is consistent with a marginal dependency (two sigma level) of this relation as a function of the normalized radius. The  correlation between these two parameters appears weaker in the cluster core when compared with larger clustercentric distances. This result indicates that disks of galaxies in the central region of clusters are more compact and have a brighter central luminosity. Nevertheless, it should be remembered that, on average, disks are less luminous in the central region than in the outskirts. In summary, disks in the core of clusters are less luminous, more compact and present a higher central surface brightness. The simulations performed by Moore et al. (1999) indicate that galaxy harassment is particularly strong for low surface brightness galaxies, which suggests that disks that can survive in the cluster core are the most compact ones, a scenario that is consistent with our results.  Finally, it is important to notice that for most of the analysis made in this work the results clearly differ depending on whether redshift-confirmed members or galaxies in projection are considered. This indicates the importance of using redshift confirmed members."
    },
    "0601/astro-ph0601566_arXiv.txt": {
        "abstract": "GW~Vir variables are the pulsating members in the spectroscopic class of the PG~1159 stars. In order to understand the characteristic differences between pulsating and non-pulsating PG\\,1159 stars, we analyse FUSE spectra of eleven objects, of which six are pulsating, by means of state-of-the-art NLTE model atmospheres. The numerous metal lines in the FUV spectra of these stars allow a precise determination of the photospheric parameters. We present here preliminary results of our analysis. ",
        "introduction": "\\label{sec:introduction} GW~Vir variables \\index{GW~Vir variables} belong to the spectroscopic class of the PG~1159 stars \\index{PG~1159 stars} (Wesemael, Green \\& Liebert 1985), which is named after the prototype \\index{PG~1159$-$035} PG~1159$-$035. These objects are strongly hydrogen-deficient post-AGB stars \\index{post-AGB stars} which pass through the hottest stage of stellar evolution. Their effective temperatures range between 75~000 and 200~000\\,K, surface gravities vary from $\\log g = 5.5-8.0\\,\\,[\\mathrm{cm\\,s}^{-2}]$. The so-called born-again scenario (a late thermal pulse which transferred these objects back to the AGB followed by a second post-AGB evolution, Iben et al\\@. 1983) is mainly accepted as an explanation for the H-deficiency and can reproduce well the observed abundances. PG~1159 stars have spectra which are dominated by lines of \\ion{He}{ii}, \\ion{C}{iv}, and \\ion{O}{vi} (Werner et al. 2004), their atmospheres show a typical surface composition of He:C:O = 33:50:17 by mass. Beside these main constituents there are several lines of trace elements, such as neon, nitrogen, silicon, sulfur, phosphorus, and fluorine (Reiff et al. 2005, Werner et al. 2005). Presently 37 PG~1159 stars are known, eleven of them proved to be pulsators. The pulsating members of the PG~1159 class are referred as GW~Vir variables. They are non-radial g-mode pulsators with periods from 300\\,s up to 1000\\,s, in some cases exceeding even 2000\\,s (Nagel \\& Werner 2004). The favored excitation mechanism for the pulsations is the $\\kappa$-mechanism associated with cyclic ionization of carbon and oxygen (Quirion, Fontaine \\& Brassard 2004). In the $\\log T_{\\mathrm{eff}} - \\log g$ diagram the GW~Vir variables are located among the PG~1159 stars in the so-called GW~Vir instability strip. Spectral analyses of pulsating and non-pulsating PG~1159 stars were used by Dreizler \\& Heber (1998) to define empirically the edges of this instability strip. But it is still puzzling that also non-pulsating PG~1159 stars are located within the instability strip. In our analysis we try to find more characteristic properties to distinguish between pulsating and non-pulsating members of this class.  \\begin{figure}[t] \\centerline{\\hbox{\\psfig{figure=fig1.eps,width=13.4cm}}} \\caption[]{FUSE\\index{FUSE} spectrum and synthetic spectrum of PG\\,1159$-$035 (GW~Vir, $T_{\\mathrm{eff}} = 140~000\\,$K, $\\log g = 7.0$).} \\label{fig:pg1159} \\end{figure} ",
        "conclusions": ""
    },
    "0601/hep-ph0601053_arXiv.txt": {
        "abstract": "We revisit a scenario in which the cosmological constant is cancelled by the potential energy of a slowly evolving scalar field, or ``cosmon\". The cosmon's evolution is tied to the cosmological constant by a feedback mechanism. This feedback is achieved by an unconventional coupling of the cosmon field to the Ricci curvature scalar. The solutions show that the effective cosmological constant evolves approximately as $t^{-2}$ and remains always of the same order as the density of ordinary matter and radiation. Newton's constant varies on cosmological time scales, with $\\dot{G}_N/G_N \\ll 1/t$. $G_N$ could have been somewhat different, and possibly smaller, at the time of Big Bang nucleosynthesis. ",
        "introduction": "Many ideas have been suggested for solving the cosmological constant problem \\cite{review}, but none so far has proven to be completely satisfactory. An old idea is that the vacuum energy of ordinary fields is cancelled by some compensating field, sometimes called a ``cosmon\". The cosmon would roll down a potential hill until it had lowered the total vacuum energy to an acceptable level \\cite{barr, ford, cosmon}. The question is how the cosmon field would ``know\" when to stop rolling. If the cosmon potential happened to have a minimum at exactly the right height, it can cancel off other contributions to $\\Lambda$. However, that would obviously be nothing other than a manual setting of the total $\\Lambda$ to zero at the minimum, i.e. just the old fine-tuning of $\\Lambda$ that one is trying to avoid.  If a conspiracy of parameters is to be avoided, it is imperative that the cosmon action does not directly depend on the contributions to $\\Lambda$ coming from other sectors of the theory. But how, in that case, is the cosmon to know when to stop rolling? Some sort of feedback mechanism is required. Several such mechanisms have been proposed over the years \\cite{barr, ford, cosmon, feedback}. The one we will explore in this paper was proposed in the papers of Ref. \\cite{barr,ford} in the mid 1980s, and is based on the idea that the cosmon field couples to the scalar curvature $R$. Since the trace of the Einstein field equations tells us that $-R = 2 \\Lambda + \\frac{1}{2} (\\rho - 3 p)$, the cosmon can learn about $\\Lambda$ through its effect on $R$. The conventional wisdom, based on a ``no-go theorem\" stated in Weinberg's well-known 1989 review of the cosmological constant problem, is that this kind of feedback mechanism cannot work. However, this ``theorem\", which originally comes from the papers \\cite{barr,ford}, does not apply to the kind of model we discuss in this paper, as we shall see. (Note: we will be using units where $16 \\pi G_N = 1$, signature -+++, and the sign convention in which a sphere has negative $R$.)  Before exploring this idea further, there is a general point that may be important. In any realistic feedback mechanism, it would seem that the ``effective vacuum energy\" $\\lambda$ --- by which we mean the vacuum energy from other sectors of the theory ($\\Lambda$) plus the potential energy in the cosmon field ($V_{cosmon}$) --- must fall to zero as time passes in such a way that it is always of roughly the same order as the matter energy density, $\\rho_{matter}$. By ``matter\" here we mean all forms of particles, including baryons, photons, neutrinos, dark matter particles, etc. The reason for this is as follows. If $\\lambda$ falls off significantly more {\\it slowly} than $\\rho_{matter}$, then it will quickly come to dominate the universe, which will end up with $\\Omega_{matter}$ being negligibly small. If, on the other hand, $\\lambda$ falls off significantly {\\it faster} than $\\rho_{matter}$, then very early in the history of the universe the effective cosmological constant would be much smaller than $T^4$, where $T$ is the temperature of the universe. However, that would mean that any subsequent phase transition (such as one associated with the symmetry breaking of grand unification, the weak interactions, or the chiral symmetry of the quarks) would drive $\\lambda$ very negative, since it would lower $\\Lambda$ by an amount of order $T^4$. Given the arguments of Coleman and De Luccia \\cite{coleman} and more recently from a string theoretic viewpoint, Banks \\cite{banks}, it is questionable whether such a transition could occur. But if it did, once the universe has negative $\\lambda$, it is not clear how it could get out of that hole, as it would have to in order to be consistent with present observations.  The upshot is that any successful feedback mechanism would seem to require that the hidden energy density of $\\Lambda$ plus ``cosmon\" field energy must be roughly of order the density of ordinary matter at all times. This is very suggestive, given the puzzling fact that a dark energy has been observed that is comparable to the matter energy today.  The feedback model we discuss does indeed have the desired feature that $\\lambda \\sim \\rho_{matter}$ at all times. As we shall see, however, this does not necessarily account for the dark energy that has been observed. The problem is that any mechanism that is designed to erase $\\Lambda$ will tend to erase dark energy too. In the model we discuss there is indeed energy that is ``dark\" (namely $\\lambda = \\Lambda + V_{cosmon}$), which generally falls to zero roughly as $t^{-2}$ and has $w \\cong 1/3$. However, after a phase transition $\\lambda$ tends to fall more slowly for a period of time, so that $w$ can be smaller than $1/3$ and perhaps even close to $-1$. It is possible that this fact can be used to account for the nearly constant dark energy that has recently been observed.  To return to the specific idea that the feedback occurs through coupling of the cosmon to the scalar curvature $R$, the obvious way for this to happen is through the cosmon potential. To illustrate how this might work, consider a toy model \\cite{barr,ford} with $V_{cosmon} = - \\alpha \\phi R$, where $\\phi$ always denotes the cosmon in this paper. The effective vacuum energy, as we defined it above, is then $\\lambda = \\Lambda - \\alpha \\phi R$. As the cosmon rolls down its linear potential hill, $\\lambda$ will decrease; and as that happens, the scalar curvature will decrease as well. This decrease of the scalar curvature reduces the slope ($- \\alpha R$) of the cosmon potential, and therefore the cosmon rolling will decelerate (since there is friction due to the Hubble expansion). If we assume that $R$ receives contributions only from $\\lambda$, then as $\\lambda$ approaches zero, the slope of the cosmon potential will approach zero also, and the velocity of the cosmon's rolling will approach zero. In other words, whatever the value of $\\Lambda$ is, the quantity $\\lambda$ will asymptotically approach zero as time passes. That is, the potential energy of the cosmon, $- \\alpha \\phi R$ will progressively erase $\\Lambda$.  In fact, this is what actually would happen in this toy model, as an explicit solution of the field equations shows. Unfortunately, however, Newton's constant $G_N$ suffers a disaster. Instead of the usual Einstein-Hilbert action $-R$, one has $-(1 - \\alpha \\phi) R$. Since the quantity $\\alpha \\phi R_U$ (where $R_U$ is the background scalar curvature of the universe) is asymptotically approaching $\\Lambda$, that means that $(16 \\pi G_N)_{eff}^{-1} = 1 - \\alpha \\phi$ is asymptotically approaching $1 - \\Lambda/R_U$, which is exponentially large at present ($\\sim 10^{120}$). In other words, Newton's ``constant\" turns off in a drastic way.  This is the problem that seems to afflict any attempt to construct a feedback mechanism through a coupling of the cosmon to the scalar curvature through the cosmon potential. (This is the conclusion that was reached in the papers \\cite{barr,ford} and repeated in Weinberg's 1989 review \\cite{review}.) This led to an attempt in \\cite{barr} to achieve a viable feedback through the {\\it kinetic} term of the cosmon. What is required is that negative powers of $R$ appear in the cosmon kinetic term, which, of course, is a peculiar thing indeed \\cite{overR}. Nevertheless, a model that erases the cosmological constant without any fine-tuning of parameters at the classical level can easily be constructed as shown in \\cite{barr}. In the paper we correct and extend the analysis in that earlier paper, and address several important issues that were not considered there.  This paper is organized as follows. In section 2 we present the toy model that we are discussing and find cosmological solutions for the cases where cosmon energy is dominant and where radiation is dominant. (In \\cite{barr} only the cosmon-dominated case was considered, and the analysis of it was faulty, since certain assumptions were made that were not self-consistent.) We also discuss the stability of the solutions, and what the cosmon equation of state looks like. We find that while there are modifications to the effective Newton's constant (because of the appearance of $R$ in the cosmon kinetic term), they are of order 1 and harmless.  In section 3, we look at what happens when baryonic matter (or other matter whose stress-energy tensor is not traceless) is present. We confirm the conclusion of \\cite{barr} that this does not disturb the feedback mechanism for erasing the cosmological constant. However, there is an interesting effect: the cosmon field does wipe out the trace of the stress-energy tensor inside baryonic matter (contrary to what was claimed in \\cite{barr}). That is, inside a lump of baryonic matter (e.g. a nucleon, atomic nucleus, or neutron star) the cosmon field has a pressure and density such that the total stress-energy tensor is traceless. However, this has less dramatic effects than one might suppose. This will be analyzed in section 3.  In section 4, we discuss other issues, such as stability of solutions, what happens if the vacuum energy is suddenly reset by a phase transition, and quantum corrections to the model and fine-tuning.  In the Appendix we point out certain mistakes and deficiencies in the analysis in \\cite{barr}. ",
        "conclusions": "We have analyzed a toy model that realizes the idea of a feedback mechanism for cancelling the cosmological constant at the classical level. The model has several appealing features. First, the total ``dark energy\" of vacuum energy plus cosmon field energy is always of the same order as the energy in radiation and matter. This is certainly suggestive, given the observed ``cosmic coincidence\" that the dark energy is currently of the same order as the energy in particles. Second, the mechanism works even if there are several phase transitions throughout the history of the universe that reset the cosmological constant. As explained in the introduction, any successful cancellation mechanism for the cosmological constant must have these features. There are also two interesting modifications of gravity. (a) The effective Newton's constant changes slightly over cosmological time scales during periods when cosmon energy is subdominant. (b) The cosmon stress energy acts to make the total stress energy tensor approximately traceless everywhere, even inside baryonic matter. There are several aspects of the model that are less attractive. The most serious is that the scalar curvature appears (at least effectively) in the denominator of the cosmon kinetic term. This makes it questionable whether this model would make sense at the quantum level. Moreover, naive power counting suggests that the model must be fine tuned at the one loop level by the same amount that is involved in solving the cosmological constant ``by hand\". There is also the difficulty that the rolling of the cosmon field tends to erase any dark energy that has $w \\cong -1$ just as it erases a true cosmological constant. In other words, the cancellation mechanism works {it too} well! However, after phase transitions, the cosmon field tends to stop rolling for a while, so that the ``effective cosmological constant\" is approximately constant and may explain the observed dark energy if the universe has recently undergone a phase transition. This deserves further study.    There are several other aspects of the model that also require further study. It must be seen what the implications of the tracelessness of the stress-energy tensor are. The cosmology of the model must be studied to see how close to being realistic it can be made. And whether the model makes sense as a quantum theory needs to be investigated. Finally, it would be interesting to see if the feedback idea can be made to work in some other way that does not involve such an exotic cosmon kinetic term.  \\newpage"
    },
    "0601/astro-ph0601221_arXiv.txt": {
        "abstract": "We present Gemini GMOS longslit spectroscopy of the isolated S0 galaxy NGC 3115. We have determined kinematical data and Lick/IDS absorption line-strength indices for the major axis out to around 9 kpc and for the minor axis out to around 5 kpc (around 2R$_e$). Using stellar population models which include the effects of variable [$\\alpha$/Fe] ratios we derive metallicities, abundance ratios and ages for the stellar population of NGC 3115. We find that [$\\alpha$/Fe] remains fairly constant with increasing radius at around [$\\alpha$/Fe] = 0.17 for the major axis but increases rapidly for the minor axis to around [$\\alpha$/Fe] = 0.3. We also find that to first order this behaviour can be explained by a simple spheroid + disc model, where the spheroid has [$\\alpha$/Fe] = 0.3 and the disc shows close to solar abundance ratios. The disc also appears considerably younger than the spheroid, having an age of around 6 Gyr compared to 12 Gyr for the spheroid. We compare these results to those previously presented for the globular cluster system of NGC 3115. ",
        "introduction": " ",
        "conclusions": ""
    },
    "0601/astro-ph0601017_arXiv.txt": {
        "abstract": "{We study the charge screening effect in the hadron-quark mixed phase. By including the charge screening effect, rearrangement of charged particles occurs and some part becomes locally charge-neutral. As a result the equation of state for the mixed phase becomes close to that given by the Maxwell construction, which means that the Maxwell construction would effectively gain the physical meaning again even in the system with two or more chemical potentials. We also discuss the interplay of the surface tension and the Coulomb interaction. Both effects would restrict the region of the mixed phase in the core of hybrid stars. }  \\FullConference{29th Johns Hopkins Workshop on current problems in particle theory: strong matter in the heavens\\\\ 1-3 August\\\\ Budapest}  \\begin{document} ",
        "introduction": "It is widely believed that quark matter exists in high-temperature and/or high-density situations like in the relativistic heavy-ion collisions (RHIC) \\cite{rhic1} or in the core of neutron stars \\cite{mad3,chen}, and the ``deconfinement transition'' has been actively searched. Theoretical studies using model calculations or based on the first principle, lattice QCD \\cite{rev} have been also carried out by many authors to find the critical temperature. Although many exciting results have been reported, the deconfinement transition has not been clearly understood yet.  We, hereafter, consider the phase transition from hadron phase to three-flavor quark phase in neutron-star matter in a semi-phenomenological way. Since many theoretical calculations have suggested that the deconfinement transition should be of first order in low-temperature and high-density area \\cite{pisa,latt}, we assume the first-order phase transition in this paper. If the deconfinement transition is of first order, one may expect a {\\it mixed phase} during the transition. Actually, the hadron-quark mixed phase has been considered during the hadronization era in RHIC \\cite{rhic21,rhic22,rhic23} or in the core region of neutron stars \\cite{gle2,web1,web2}.  There is an issue about the mixed phase for the first-order phase transitions with two or more chemical potentials \\cite{gle1}. We often use the Maxwell construction (MC) to derive the equation of state (EOS) in thermodynamic equilibrium, as in the liquid-vapor phase transition of water, where both phases consist of single particle species. However, if many particle species participate in the phase transition as in neutron-star matter, MC is no more an appropriate method. Before Glendenning first pointed out \\cite{gle1}, many people have applied MC to get EOS for the first-order phase transitions \\cite{weis,migd,elli,rose} expected in neutron stars, such as pion or kaon condensation and the deconfinement transition.  Let us consider a deconfinement transition in neutron-star matter. We must introduce many chemical potentials for particle species, but the independent ones in this case are reduced to two, i.e.\\ baryon-number chemical potential $\\mu_\\mathrm{B}$ and charge chemical potential $\\mu_\\mathrm{Q}$, due to chemical equilibrium and total charge neutrality. They are nothing but the neutron and electron chemical potentials, $\\mu_n$ and $\\mu_e$, respectively. In the mixed-phase these chemical potentials should be spatially constant. When we naively apply MC to get EOS in thermodynamic equilibrium, we immediately notice that $\\mu_\\mathrm{B}$ is constant in the mixed-phase, while $\\mu_\\mathrm{e}$ is different in each phase because of the difference of the electron number density in these phases. This is because MC uses EOS of bulk matter in each phase, which is of locally charge-neutral and uniform matter; many electrons are needed in hadron matter to compensate the positive charge of protons, while in quark matter charge neutrality is almost fulfilled without electrons. Thus \\begin{equation} \\mu_{\\mathrm{B}}^{\\mathrm{Q}} = \\mu_{\\mathrm{B}}^{\\mathrm{H}}, \\hspace{10pt} \\mu_{\\mathrm{e}}^{\\mathrm{Q}} \\neq  \\mu_{\\mathrm{e}}^{\\mathrm{H}}, \\label{chemeqMC} \\end{equation} in MC, where superscripts ``Q'' and ``H'' denote the quark and hadron phases, respectively. Glendenning emphasized that we must use the Gibbs conditions (GC) in this case instead of MC, which relaxes the charge-neutrality condition to be globally satisfied as a whole, not locally in each phase \\cite{gle1}. GC imposes the following conditions, \\begin{eqnarray} \\mu_{\\mathrm{B}}^{\\mathrm{Q}} &= \\mu_{\\mathrm{B}}^{\\mathrm{H}}, \\hspace{5pt} \\mu_{\\mathrm{e}}^{\\mathrm{Q}} = \\mu_{\\mathrm{e}}^{\\mathrm{H}}, \\nonumber \\\\ P^{\\mathrm{Q}} &=P^{\\mathrm{H}} , \\hspace{5pt} T^{\\mathrm{Q}} = T^{\\mathrm{H}}. \\end{eqnarray} He demonstrated a wide region of the mixed phase, where two phases have a net charge but totally charge-neutral: EOS thus obtained, different from that given by MC, never exhibits a constant-pressure region. He simply considered a mixed phase consisting of two bulk matters separated by a sharp boundary without any surface tension and the Coulomb interaction, which we call ``bulk Gibbs'' for convenience. ``Bulk Gibbs'' requires that each matter can have a net charge but the total charge is neutral, \\begin{equation} f_V \\rho_{\\mathrm{ch}}^\\mathrm{Q} + (1-f_V) \\rho_{\\mathrm{ch}}^\\mathrm{H} = 0, \\end{equation} where $f_V$ means the volume fraction of quark matter in the mixed phase and ``$\\rho_\\mathrm{ch}^\\mathrm{Q,H}$'' means charge density in each matter. Figure \\ref{bulk} shows the phase diagram in the $\\mu_\\mathrm{B}$ - $\\mu_e$ plane. We can see that there is a discontinuous jump in $\\mu_e$ for the case of MC, while the curve given by ``bulk Gibbs'' smoothly connects uniform hadron matter and uniform quark matter; the mixed phase can appear in a wide $\\mu_\\mathrm{B}$ region in ``bulk Gibbs'', in contrast with MC \\cite{alf2}. \\begin{figure}[h!]\\begin{center} \\includegraphics[width=80mm]{mubmue00color.eps} \\caption{ Phase diagram in the $\\mu_\\mathrm{B}-\\mu_\\mathrm{e}$ plane. There appears no region of the mixed phase by the calculation with MC, while a wide region of the mixed phase by the ``bulk Gibbs'' calculation.} \\label{bulk} \\end{center} \\end{figure}  However, this ``bulk Gibbs'' is too simple to study the mixed phase, since we must consider matter with non-uniform structures instead of two bulk uniform matters; the mixed phase should have various geometrical structures where both the number and charge densities are no more uniform. Then we have to take into account finite-size effects like the surface and the Coulomb interaction energies. We show that the mixed phase should be narrow in the $\\mu_\\mathrm{B}$ space by the charge screening effect, and derive EOS for the deconfinement transition in neutron-star matter. We shall see EOS results in being similar to that given by MC. We also discuss the interplay of the Coulomb interaction effect and the surface effect in the context of the hadron-quark mixed phase. Preliminary results for the droplet case has been already reported in Ref.\\ \\cite{end1}. ",
        "conclusions": "\\label{summary}  We have studied the charge screening effect in the hadron-quark mixed phase in neutron-star matter. We have elucidated the charge screening effect comparing the results of ``no screening'' with those of ``screening''  which includes non-linearity of the Poisson equation. The density profiles of the charged particles are drastically modified by the screening effect, while the thermodynamic potential is not affected so much; the charge rearrangement induced by the screening effect tends to make the net charge smaller in each phase. Consequently the system tends to be locally charge-neutral, which suggests that the Maxwell construction (MC) is effectively justified even if it is thermodynamically incorrect. In this context, it would be interesting to refer to the work by Heiselberg \\cite{hei}, who  studied the screening effect on a quark droplet (strangelet) in the vacuum, and suggested the importance of the rearrangement of charged particles.  We have seen that thermodynamic quantities such as thermodynamic potential and energy become close to those derived from MC by the screening effect, which also suggests that MC is effectively justified due to the screening effect. As another case of the system with two or more chemical potentials, kaon condensation  has been also studied \\cite{maru1} and the results are shown to be similar to those in the present study. Thus the importance of the screening effect should be a common feature for the first-order phase transitions in high-density matter.  We have included the surface tension at the hadron-quark interface, while its definite value is not clear at present. An uncertainty in the value of the surface tension does not allow to conclude whether the mixed phase exists or not. There are many estimations for the surface tension at the hadron-quark interface in lattice QCD \\cite{kaja,huan}, in shell-model calculations \\cite{mad1,mad2,berg} and in model calculations based on the dual-Ginzburg Landau theory \\cite{mond}. Our parameter is in that reasonable range.  We have considered some implications of our results for neutron star phenomena. The screening effect would restrict the allowed region of the structured mixed phase in neutron stars, in contrast with a wide region given by ``bulk Gibbs''\\cite {gle2,gle3}. It could be said that they should change the bulk property of neutron stars, especially the structure of the core region.  Compact stars have the strong magnetic field and its origin is not well understood. One possibility is that it comes from ferromagnetism of the quark matter in the core \\cite{tat1,tat2,tat3}. Therefore it should be interesting to include the magnetic field contribution in our formalism. We have assumed zero temperature here. It would be  much interesting to include the finite-temperature effect. Then it is possible to draw the phase diagram in the $\\mu_\\mathrm{B}$ - $T$ plane and we can study the properties of the deconfinement phase transition; our study may be extended to treat the mixed phase to appear during the hadronization of QGP in the nucleus-nucleus collisions and supernova explosions.  In this study we have used a simple model for quark matter to figure out the finite-size effects in SMP. However, it has been suggested that the color superconductivity would be a ground state of quark matter \\cite{alf2,alf1,alf3}. Hence we will include it in a further study. The hadron phase should be also treated more realistically; for example, we should include the hyperons or kaons in hadron matter. In the recent studies the mixed phase has been also studied \\cite{neum,shov,redd} in the context of various phases in the color superconducting phase."
    },
    "0601/astro-ph0601684_arXiv.txt": {
        "abstract": "We present 39 nights of optical photometry, 34 nights of infrared photometry, and 4 nights of optical spectroscopy of the Type Ia SN 1999ac.  This supernova was discovered two weeks before maximum light, and observations were begun shortly thereafter.  At early times its spectra resembled the unusual SN~1999aa and were characterized by very high velocities in the \\ion{Ca}{2} H\\&K lines, but very low velocities in the \\ion{Si}{2} $\\lambda$ 6355 line. The optical photometry showed a slow rise to peak brightness but, quite peculiarly, was followed by a more rapid decline from maximum. Thus, the $B$- and $V$-band light curves cannot be characterized by a single stretch factor.  We argue that the best measure of the nature of this object is {\\em not} the decline rate parameter \\dmm.  The $B-V$ colors were unusual from 30 to 90 days after maximum light in that they evolved to bluer values at a much slower rate than normal Type Ia supernovae.  The spectra and bolometric light curve indicate that this event was similar to the spectroscopically peculiar slow decliner SN~1999aa. ",
        "introduction": "Due to their considerable potential as extragalactic standard candles, Type Ia supernovae (SNe~Ia) have been the subject of intense study during the last 15 years.  Although it has by now been amply demonstrated that SNe~Ia cover a range in instrinsic brightness of a factor of two or more \\citep[see][and references therein] {Lei00}, the fortuitous existence of a correlation between the peak luminosity and the rate of decline from maximum of the $B$ light curve, including prescriptions for color corrections \\citep{Phi93, Ham_etal96b, Rie_etal96, Phi_etal99, Nob_etal03, Guy_etal05}, has allowed distances to be measured to better than 10 percent to redshifts $z \\leq 0.1$ \\citep{Ham_etal96a, Rie_etal96}. Because they can be observed to very large distances, SNe~Ia have become the tool of choice for measuring the Hubble constant \\citep[e.g.,][]{Ham_etal96a, Rie_etal96, San_etal96, Jha_etal99, Sun_etal99, Phi_etal99, Tri_Bra99, Gib_etal00, Par_etal00, Fre_etal01, Phi_etal03, Rie_etal05} and, in combination with measurements of fluctuations in the microwave background radiation \\citep{Ben_etal03}, have provided compelling evidence for an accelerating universe \\citep{Gar_etal98, Rie_etal98, Per_etal99, Ton_etal03, Kno_etal03, Bar_etal04, Rie_etal04, Kri_etal05}.  Although the majority of SNe~Ia display very similar spectral evolution, roughly one-third of all events show spectroscopic peculiarities \\citep{Li_etal01a}.  Three well-observed examples -- SN~1991T, SN~1991bg, and SN~1986G -- have served as the prototypes of spectroscopically-peculiar SNe~Ia \\citep{Bra_etal93}.  The pre-maximum optical spectra of SN~1991T showed unusually weak lines of \\ion{Si}{2}, \\ion{S}{2}, and \\ion{Ca}{2}, while at the same time displaying prominent high-excitation features of \\ion{Fe}{3} \\citep{Fil_etal92a, Phi_etal92}. The \\ion{Si}{2}, \\ion{S}{2}, and \\ion{Ca}{2} lines grew quickly in strength following maximum until the spectrum appeared essentially normal a few weeks past maximum.  SN~1991T dimmed slowly from maximum and was originally thought to be considerably more luminous than ``normal'' SNe~Ia \\citep{Fil_etal92a, Phi_etal92}, although the Cepheid distance to the host galaxy, NGC~4527, obtained later with the Hubble Space Telescope, implies that this supernova was only moderately over-luminous \\citep{Sah_etal01, Gib_Ste01}.  The main spectroscopic peculiarity of SN~1991bg was the presence at maximum light of a broad absorption trough at 4100-4400~\\AA~due mostly to \\ion{Ti}{2} lines \\citep{Fil_etal92b, Lei_etal93, Tur_etal96, Maz_etal97}, reflecting a lower effective temperature. As a consequence, SN~1991bg was unusually red ($B-V \\simeq 0.75$) at maximum light.  SN~1986G was similar to 1991bg, but less extreme \\citep{Phi_etal87, Cri_etal92}. Both events were substantially sub-luminous.  \\citet{Nug_etal95} have shown that the photometric sequence of SNe~Ia (i.e., the luminosity vs. decline rate relation) also manifests itself as a spectroscopic sequence which can be modelled in terms of a range in the effective temperature at maximum.  The luminous, slowly-declining SN~1991T represents the high-temperature extreme of the sequence, while the fast-declining, sub-luminous SN~1986G and SN~1991bg correspond to the low-temperature limit.  Under this scenario, none of these three SNe~Ia should be considered ``peculiar''. However, not all luminous, slow-declining SNe~Ia display SN~1991T-like pre-maximum spectra -- e.g., SN~1992bc \\citep{Maz_etal94} and SN~1999ee \\citep{Ham_etal02} -- so it is not clear if this interpretation is fully correct.  Moreover, the last few years have revealed the existence of several SNe~Ia which are of an intermediate type between SN~1991T events and ``normal'' SNe~Ia.  The prototype of these objects, which may be more common than SN~1991T-like events, is SN~1999aa \\citep{Kri_etal00, Li_etal01a, Gar_etal04}. SN~1999aa was similar to SN~1991T in displaying very weak \\ion{Si}{2} 6150~\\AA~absorption and prominent \\ion{Fe}{3} lines in its pre-maximum spectra, but it showed strong \\ion{Ca}{2} H~\\&~K absorption at the same epoch in which this feature was weak or absent in SN~1991T \\citep{Fil_etal99}.  By maximum light the spectrum of SN~1999aa had evolved to that of a ``normal'' Type~Ia supernova.  Hence, SN~1999aa-like events would be difficult, if not impossible, to distinguish spectroscopically unless a spectrum were obtained at least a week before maximum.  In order to more fully understand the range of spectroscopic and photometric characteristics of SNe~Ia, we have organized a number of observing campaigns to obtain optical photometry, infrared photometry, and optical spectroscopy of nearby events ($z \\leq 0.05$).  In this paper, we present observations obtained of the SN~1999aa-like object SN~1999ac. SN~1999ac was discovered in the Sc galaxy NGC~6063 by \\citet{Mod_etal99} from images obtained on 1999 February 26.5 and 27.5 UT. It was located at RA = 16:07:15.0, DEC = +07:58:20 (equinox 2000), some 24 arcsec east and 30 arcsec south of the core of its host. SN~1999ac was confirmed to be a Type~Ia supernova by \\citet{Phi_Kun99} from a spectrum taken on February 28 UT, who also noted that the spectrum was similar to that of the slow decliner SN~1991T \\citep{Lir_etal98}. Like SN~1991T, the \\ion{Fe}{3} lines at 4300 and 5000~\\AA~in the spectrum of SN~1999ac were observed to be strong and the \\ion{Si}{2} 6355~\\AA~absorption weak; but unlike 1991T, the \\ion{Ca}{2} H~\\&~K absorption in SN~1999ac was strong and well-developed.  Phillips \\& Kunkel speculated that SN~1999ac had been caught near, or a few days before, maximum; our observations show that $B$ maximum did not actually occur until 13 days after this first spectrum was obtained.  In \\S2 we present our optical and near-infrared photometry of SN~1999ac. The resulting light curves, which are among the most complete ever obtained of a SN~Ia for the first few months following explosion, are contrasted with the light curves of other well-observed SNe~Ia. In \\S3 we discuss the optical spectra obtained, likewise comparing these with similar observations of both spectroscopically ``normal'' and ``peculiar'' SNe~Ia.  In \\S4 we derive the most likely host galaxy reddening for SN~1999ac, and compare the peak luminosity and bolometric light curve with those of more typical SNe~Ia. Finally, in \\S5 the conclusions of this investigation are summarized. ",
        "conclusions": "The Type Ia supernova 1999ac was peculiar in several respects.  As shown by \\citet{Gar_etal05} and by our data, the spectra of SN~1999ac resembled the unusual SN~1999aa at early times.  From maximum light onward, however, the spectrum was essentially normal.  This underscores a thorny issue of SN classification: if early-time spectra are not obtained, there may be no way to determine if a particular Type Ia supernova is unusual \\citep[see also][] {Li_etal01a}.  The optical photometry of SN~1999ac was unusual in that it had a slow rise rate (like SN~1991T), but a much more rapid $B$-band decline rate (like SN~1994D). In this sense it was the opposite of SN~2000cx, which was a rapid riser, but slow decliner.  This is in striking contrast to the $B$- and $V$-band light curves of most SN~Ia which can be characterized by a single stretch factor \\citep{Gol_etal01}. In the case of SN~1999ac, our best estimate of the extinction is obtained from the $V-K_s$ colors during the first 20 days after T($B_{max}$).  We find A$_V$(tot) = 0.42 mag, implying a non-zero host galaxy reddening of E($B-V$)$_{host}$ = 0.09 mag.  In spite of the peculiar photometric properties of SN~1999ac, because near-infrared extinction corrections are an order of magnitude smaller than at optical wavelengths, the infrared absolute magnitudes of SN~1999ac cannot be systematically in error by more than a few hundredths of a magnitude. The derived infrared absolute magnitudes \\citep{Kri_etal04a} are statistically equal to the mean of all other SN~Ia with slower decline rates.  As a result, the unusual SN~1999ac can still be regarded as an infrared standard candle.  Its $J_sHK_s$ light curves were very similar to those of the normal SN~2001el.  The bolometric light curve of SN~1999ac closely resembled that of the spectroscopically peculiar slow decliner SN~1999aa.  The differences of their bolometric luminosities at maximum compared to 90 days after maximum are almost identical.  As measured by this parameter, SN~1999ac behaved as if it were a slow decliner.  Hence, in terms of its energetics and pre-maximum spectral characteristics, there is little doubt that SN~1999ac was closely related to other peculiar, slow-declining events such as SNe~1999aa and 1991T.  Finally, we have also examined the luminosity vs. decline rate relation for SN~Ia, but considered according to their bolometric light curves.  Most events occupy a rather small parameter space, suggesting a uniform explosion mechanism for these supernovae.  \\vspace {1 cm}"
    },
    "0601/astro-ph0601637_arXiv.txt": {
        "abstract": "We summarise the science cases for an ELT that were presented in the parallel session on the intergalactic medium, and the open discussion that followed the formal presentations. Observations of the IGM with an ELT provides tremendous potential for dramatic improvements in current programmes in a very wide variety of subjects. These range from fundamental physics (expansion of the Universe, nature of the dark matter, variation of physical constants), cosmology (geometry of the Universe, large-scale structure), reionisation (ionisation state of the IGM at high $z\\ge 6$), to more traditional astronomy, such as the interactions between galaxies and the IGM (metal enrichment, galactic winds and other forms of feedback), and the study of the interstellar medium in high $z$ galaxies through molecules. The requirements on ELTs and their instruments for fulfilling this potential are discussed.  {} ",
        "introduction": "The advent of echelle spectrographs on 8m class telescopes since the early 1990's has revolutionised our understanding of the intergalactic medium (IGM) as observed in quasar spectra. These bright sources have smooth intrinsic spectra with broad emission lines, yet the {\\em observed} spectra contain hundreds of narrow absorption lines due to intervening absorbers. The latter can be studied in great detail from the exquisite, $\\rm S/N> 40$, spectra possible with UVES on VLT and HiRes on Keck.  Most of the absorption in the UV is due to neutral hydrogen left over from the Big Bang, forming a forest of lines (Bahcall \\& Salpter 1965; Gunn \\& Peterson 1965; Lynds 1971). The weaker lines with column density $N$({H~{\\sc i})$\\le 10^{15}{\\rm cm}^{-2}$ are traditionally called the \\lq Lyman-$\\alpha$ forest\\rq, with the strongest lines with column density $\\ge 10^{20.3}{\\rm cm}^{-2}$ that show a measurable damping-wing called \\lq Damped Lyman-$\\alpha$ systems\\rq\\, (DLAs). The number of lines as a function of redshift and column-density, $d^2N/dz/dN$(H~{\\sc i}), is close to a power-law $\\propto~N$(H~{\\sc i})$^{\\beta}$ with $\\beta<0$, as function of column-density, and evolves strongly with redshift as fewer lines are produced as the mean density decreases due to the expansion of the Universe, see Rauch (1998) for a review.  The weaker lines originate in the filaments of the cosmic web which itself is a natural outcome of how structure forms in a dark matter dominated cosmology (Bi \\etal\\ 1992; Cen \\etal\\ 1994; Schaye 2001). The neutral hydrogen fraction is small at redshifts $\\le 6$ (Gunn \\& Peterson 1965) as the gas is photo-ionised and photo-heated by the UV-background, with photo-ionisation rate $\\Gamma$, produced by galaxies and quasars (Haardt \\& Madau 1996). At lower $z\\le 2$, the forest of lines thins-out into a Lyman-$\\alpha$ \\lq savanna\\rq\\,, but the decline is slowed because  $\\Gamma$ also decreases as the emissivity from galaxies and QSOs drops (Theuns \\etal\\ 1998a; Dav\\'e \\etal\\ 1999). Conversely at increasing $z\\ge 6$, the mean density increases, but $\\Gamma$ also decreases as many source have yet to form, turning the forest into a Lyman-$\\alpha$ \\lq jungle\\rq\\, which absorbs (nearly) all light, perhaps signaling the end of reionisation (Becker \\etal\\ 2001; Djorgovski \\etal\\  2001). The forest provides a tremendous probe of how the IGM evolves in the intermediate redshift range $2\\le z\\le 5$, because the absorbers are only mildly non-linear and hence can be simulated reliably (Cen \\etal\\  1994; Hernquist \\etal\\  1996; Theuns \\etal\\ 1998b; Zhang \\etal\\  1998; Bryan \\etal\\  1999). The combination of superb data with reliable models makes it possible to constrain models and determine parameters.  Stronger lines form near galaxies, with the DLAs potential proto-galaxies or proto-galactic lumps (Wolfe 1995; Haehnelt \\etal\\ 1998; Ledoux \\etal\\ 1998).  Since these systems are discovered in absorption, it is a worry that even denser systems might be missed because they make the background QSO too faint to appear in a magnitude-limited survey. For a recent appraisal of this issue see Ellison \\etal\\  (2005 and reference therein).  DLAs shield the UV-background and some fraction of the gas becomes molecular (see, e.g., Srianand this volume). The prospect of studying star formation in small systems at high redshift which are too faint to study in emission, is very exciting.  Quasar spectra also contain \\lq metal\\rq\\, lines from highly ionised species such as \\cfour\\,, \\sifour\\, and \\osix\\, (e.g., Cowie \\etal\\ 1995). These metals were synthesised in stars and managed to diffuse into the lower density surroundings, either as a result of galactic winds, or due to an early generation of population~III~ stars.  The next section gives a short overview of recent results, with emphasis on opportunities for progress with the advent of new observatories. ",
        "conclusions": ""
    },
    "0601/astro-ph0601401_arXiv.txt": {
        "abstract": "We use a combination of high-resolution gasdynamics simulations of high-redshift dwarf galaxies and dissipationless simulations of a Milky Way sized halo to estimate the expected abundance and spatial distribution of the dwarf satellite galaxies that formed most of their stars around $z\\sim 8$ and evolved only little since then. Such galaxies can be considered as {\\it fossils\\/} of the reionization era, and studying their properties could provide a direct window into the early, pre-reionization stages of galaxy formation.  We show that $\\sim 5-15\\%$ of the objects existing at $z\\sim 8$ do indeed survive until the present in the \\MW\\ like environment without significant evolution. This implies that it is plausible that the fossil dwarf galaxies do exist in the Local Group. Because such galaxies form their stellar systems early during the period of active merging and accretion, they should have spheroidal morphology regardless of their current distance from the host galaxy. Their observed counterparts should therefore be identified among the dwarf spheroidal galaxies. We show that both the expected luminosity function and spatial distribution of dark matter halos which are likely to host fossil galaxies agree reasonably well with the observed distributions of the luminous ($L_V\\ga10^6\\Lsun$) Local Group fossil candidates near the host galaxy ($d\\la200\\dim{kpc}$). However, the predicted abundance is substantially larger (by a factor of 2-3) for fainter galaxies ($L_V<10^6\\Lsun$) at larger distances ($d\\ga300\\dim{kpc}$). We discuss several possible explanations for this discrepancy. ",
        "introduction": "\\label{sec:intro}  Abundance and spatial distribution of the faintest dwarf galaxies provide information and constraints on the galaxy formation processes and may give important clues to the nature of dark matter. Cosmological simulations \\citep{kkvp99a,mggl99a} and semi-analytic models of galaxy formation \\citep{kwg93a,bkw00a,s02a,bflb02a} predict the number of gravitationally bound, dark matter clumps (or {\\it subhalos\\/}) around the \\MW\\ type galaxies larger, by up to an order of magnitude, than the observed number of dwarf galaxies in the Local Group.  In addition, the predicted spatial distribution of the dark matter subhalos around their hosts is more extended than the observed distribution of the Local Group satellites \\citep{tsb04b,kgk04a}, indicating that strong spatial bias in galaxy formation should have been present if the CDM scenario is correct.  Formation of galaxies in small-mass, dwarf-sized halos should be affected by heating from the ultra-violet (UV) radiation and internal feedback from stars. UV heating, in particular, is expected to be highly important after the epoch of reionization when the entire volume of the Universe is affected by the cosmic UV background. The UV heating leads to dramatic increase of the characteristic Jeans or {\\it filtering\\/} mass, $M_{\\rm F}$, of intergalactic gas \\citep{gh98a,g00c}. The dark matter halos with masses $M\\lesssim M_{\\rm F}$ (or circular velocities $V_{\\rm c}\\lesssim V_{\\rm F}$) cannot accrete new gas and therefore do not have supply of fresh gas to fuel the continuing star formation \\citep[e.g.,][]{e92a,tw96a,qke96a,ns97a,ki00a,g00c,dhrw04a}. In addition, the UV background can evaporate the existing gas in small halos \\citep{bl99a,sd03a,sir04a}, which would also reduce the star formation rate.  In the presence of these galaxy formation suppression processes, the smallest dwarf galaxies can form: (1) if their host halos have formed sufficiently early, before the reionization epoch and UV heating became significant \\citep{bkw00a} and/or (2) if the host halos become sufficiently massive after reionization such that UV heating effects are not significant \\citep{kgk04a}.  It is of course possible that a halo forms early and forms stars before reionization and then stays above filtering mass for an extended period of time after reionization due to continuing merging and accretion. Such galaxies would exhibit continuous, albeit perhaps episodic \\citep[see][]{kgk04a}, star formation over an extended period of time.  It is also possible that some galaxies form stars before reionization and then not evolve much in terms of their total and stellar mass, with their stellar populations evolving mostly passively.  Most of such galaxies do not survive until $z=0$ but are disrupted via merging and tidal heating, contributing to the formation of the bulge, stellar halo, and halo population of globular clusters of the host galaxies \\citep[e.g.,][]{ws00a,bkw01a,kg05a,bj05a,rea05b,mdmzs05a}.  However, a fraction of them can survive. \\citet{rg05a} found that many of the dwarf galaxies identified in the high-resolution gas dynamic simulation at $z\\approx 8$ do indeed resemble a subset of the observed old dwarf spheroidal galaxies in the Local Group in their kinematic, structural, and chemical properties.  Such galaxies would then represent {\\it fossils\\/} of the reionization era, and studying their properties could provide a direct window into the early, pre-reionization stages of galaxy formation.  As such, it would be extremely interesting to identify and study this population in the Local Group. In this paper, we present theoretical estimates for the abundance and spatial distribution of the fossil galaxies using a combination of sophisticated galaxy formation simulations at high redshifts, which include non-equilibrium chemistry and self-consistent, three-dimensional radiative transfer and high-resolution dissipationless simulations of a \\MW-sized halo, which allows us to track small-mass dwarf dark matter halos from the pre-reionization epoch ($z\\sim 10$) to their present-day \\MW-like environment. ",
        "conclusions": "\\label{sec:discussion}  In this paper we estimate the expected abundance and spatial distribution of the reionization fossils: dwarf galaxies that formed most of their stars around $z\\sim 8$ and evolved only little since then. We show that $\\sim 5-15\\%$ of the objects existing at $z\\sim 8$ do indeed survive until the present in the \\MW\\ like environment without significant evolution. Because such galaxies form their stellar systems early during the period of active merging and accretion, they should have spheroidal morphology and their observed counterparts are most likely to be identified among the dwarf spheroidal galaxies. The spheroidal morphology of the fossil galaxies is thus not related to the tidal heating by the nearby massive galaxies and they can be found even at distances as large as a megaparsec from their host. This could possibly explain existence of distant dwarf spheroidal galaxies in the Local Group, such as Cetus, Tucana, and KKR25.  We show that both the expected luminosity function and spatial distribution of dark matter halos which are likely to host fossil galaxies agrees reasonably well with the observed distribution of the luminous ($L_V\\ga10^6\\Lsun$) Local Group fossil candidates near the host ($d\\la200\\dim{kpc}$). The predicted abundance is substantially larger (by a factor of 2-3) for fainter galaxies ($L_V<10^6\\Lsun$) at larger distances ($d\\ga300\\dim{kpc}$). This discrepancy can have at least three plausible explanations. We list them and the possible tests below.  \\begin{description} \\item[Observations are wrong.] Observational incompleteness for faint, $L_V<10^6\\Lsun$, fossils may potentially account for the difference we find, if it can be as high as a factor of 3 at $d>300\\dim{kpc}$. \\citet{wgdr04a} found that a factor of 3 incompleteness is possible, although that number was on the very edge of their estimates. The same incompleteness factor would have to apply to the distribution of Andromeda satellites, which seems to be less likely. Future observational searches (for example, from SDSS) for dwarf spheroidals between $200\\dim{kpc}$ and $1\\dim{Mpc}$ from the \\MW\\ {\\it or\\/} Andromeda, or observations of nearby galaxy groups with the next generation large telescopes should be able to test this possibility. \\item[Theory is wrong.] The abundance of fossils galaxies may be overestimated by our model. This is possible for example if reionization occurs much earlier than we assumed: for example, at $z>10$, as may be suggested by the first year WMAP data \\citep{ksbb03a}. In this case, the number of small-mass objects that are able to form stars before reionization should be much smaller. However, if most of dwarf galaxies that \\citet{rg05a} identified as reionization fossils are confirmed to be such by future observations, then this work will places a constraint of $z\\la10$ on the reionization history of the region of the universe that collapses into the Local Group. \\item[Spatial bias.] If we take the observational data as complete and at face value, they indicate that lowest luminosity fossils ($L_V\\la10^6\\Lsun$) are located preferentially near the parent galaxy. At first glance, such a conclusion seems counterintuitive, since the feedback (both radiative and kinetic) from the progenitor of the main galaxy is expected to suppress star formation in nearby galaxies. However, galaxies are also biased, more so at high redshift, so dwarf galaxies located closer to the progenitor of a large spiral galaxy will start forming sooner than similar galaxies in more distant regions of space. Which of the two opposing effects (feedback vs bias) wins is unclear, but should be tackled with the future ultra high-resolution cosmological simulations. \\end{description}  The results of this study indicate that existence of the fossil galaxies in the Local Group, in which most of the stars formed during the earliest stages of galaxy formation, is plausible.  The abundance of fossils predicted by our model would constitute as much as $20-40\\%$ of all the luminous Local Group dwarfs.  The improvement in resolution of cosmological simulations should improve our estimates. Observationally, identification and studies of the fossil galaxies can provide an exciting opportunity of studying the properties of star formation in primeval galaxies by direct observations in our own cosmological backyard."
    },
    "0601/astro-ph0601292_arXiv.txt": {
        "abstract": "{}% {By neglecting sideways expansion of gamma-ray burst (GRB) jets and assuming their half-opening angle distribution, we estimate the detectability of orphan optical afterglows.}% {This estimation is carried out by calculating the durations of off-axis optical afterglows whose flux density exceeds a certain observational limit.}% {We show that the former assumption leads to more detectable orphans, while the latter suppresses the detectability strongly compared with the model with half-opening angle $\\theta_j=0.1$. We also considered the effects of other parameters, and find that the effects of the ejecta energy $E_j$ and post-jet-break temporal index $-\\alpha_2$ are important but that the effects of the electron-energy distribution index $p$, electron energy equipartition factor $\\epsilon_e$, and environment density $n$ are insignificant. If $E_j$ and $\\alpha_2$ are determined by other methods, one can constrain the half-opening angle distribution of jets by observing orphan afterglows. Adopting a set of ``standard\" parameters, the detectable rate of orphan afterglows is about $1.3\\times 10^{-2} {\\rm deg}^{-2} {\\rm yr}^{-1}$, if the observed limiting magnitude is 20 in R-band. }% {}% ",
        "introduction": "Orphan afterglows are defined as afterglows whose gamma-ray bursts (GRBs) are not detected, possibly because of the Doppler effect for an off-axis observer. {If the GRB afterglows are modelled perfectly, } the observed rate of orphans and GRBs can be used to constrain the beaming factor of GRBs, as first proposed by \\citet{r97}. However, as many parameters are not determined well for different afterglows, it will be difficult to constrain the beaming factor tightly from optical orphans. Because late radio afterglows behave isotropic emission, the survey of radio afterglows may be helpful for estimating the beaming factor \\citep{lo02, go05}, although one should be careful to rule out radio transients from other sources.  Other than constraining the beaming factor \\citep{r97, dgp02, tp02}, some authors have focused on investigations of the detectability of orphan optical afterglows both theoretically \\citep{npg02, tp02} and experimentally \\citep{h04, bw04, ra05, ma06, rgs06}. \\citet{bw04} give the results of a 5-year (1999-2004) survey of optical transients, but none was identified as an orphan. \\citet{ra05} performed a 1.5-year survey (2003 September to 2005 March) of untriggered GRB afterglows. Although no orphan afterglow has been observed yet, they give the upper limit of the observed rate for a certain limiting magnitude. The surveys are still going on\\citep{ma06,mab06}. On the theoretical side, \\citet{npg02} considered the following afterglow model: all jets propagating in a uniform medium (ISM) have a constant initial half-opening angle $\\theta_j$ and a constant jet energy $E_j$. After a jet break takes place because the hydrodynamics of a sideways-expansion jet enters an exponential regime \\citep{r99,sp99}, the temporal index of light curves becomes $-p$ (where $p$ is the power-law index of shock-accelerated electrons). This decline is too steep for  most of the observed late afterglows\\citep{lz05}. On the other hand, many works \\citep{msb00, hgd00, wl00, s03, kg03, gk03, cgv04} show that the sideways expansion of jets is insignificant at the relativistic stage. Thus, we consider relativistic jets without sideways expansion, and their afterglow light curves for an on-axis observer are shown in Fig. \\ref{f_nu_t_theta}. The light curves are shown in the spherical case and the flux density $f_\\nu \\propto t^{-\\alpha_1}$ {when the Lorentz factor $\\gamma > 1/\\theta_j$}. After the jet break time, the light curves steepen { as $f_\\nu \\propto t^{-\\alpha_2}\\,(\\alpha_2>\\alpha_1)$} because of the edge effect\\citep{mr99}, which is flatter than the sideways-expansion case. This will lead to more detectable orphan afterglows. A relationship between the jet break time and flux density ($f_{\\nu,j}\\propto t_j^{-p}$) was found by \\citet{wdl04} analytically and statistically.  \\begin{figure} \\includegraphics[angle=270,width=0.5\\textwidth]{4853fig1.ps} \\caption{Sketch of afterglow light curves { from jets} without sideways expansion for  an on-axis observer. Dotted, { dot-dashed} and solid lines correspond to three jet opening angles 0.03, 0.1 and 0.3 respectively, with the same total kinetic energy $E_j = 1\\times 10^{51}$ ergs. The dashed line is the connection of jet breaks.} \\label{f_nu_t_theta} \\end{figure}  Fig. \\ref{f_nu_t_theta} is plotted based on the fact that the kinetic energies of all GRB jets have a similar value \\citep{fks01,bf03}. The statistical standard energy of jets has been discussed by several authors, e.g., \\citet{lpp01}, and \\citet{pk02}, who conclude that the collimation-corrected gamma-ray energy has a relatively narrow distribution, around $5\\times 10^{50}$ ergs. \\citet{bkf03} also obtained a standard kinetic energy reservoir of afterglows from the statistics on X-ray luminosity.  Recently, the distribution of the half-opening angle or viewing angle was investigated based on several structured-jet models \\citep{psf03, lwd04, ngg04, gpw05}. Considering a uniform, sharp-edge jet \\citep[favored by][]{ldg05}, whose light curves are similar to those in  the universal structured jet models \\citep{rlr02}; and using the observed distribution of half-opening angles of the jets given by \\citet{ldg05}, one can derive the intrinsic distribution of $\\theta_j$, i.e., $P(\\theta_j) \\propto \\theta_j^{-1}$ \\citep[see also ][]{xwd05}. We use this distribution as a weight for different opening angles as suggested by \\citet{gpw05}.  By considering both the effects of the constant half-opening angle { during} jet propagation and the distribution of {initial jet half-opening angles}, we here estimate the detectability of orphan afterglows and find that our results are different from the ones in earlier works \\citep[e.g.][]{npg02, tp02}. We { present} the theoretical model in \\S \\ref{sec:analysis} and { give the results of} the detectability in \\S \\ref{sec:results}. We summarize our findings and present a brief discussion in \\S \\ref{sec:discuss}. ",
        "conclusions": "\\label{sec:discuss}  We have calculated the number of orphan optical {(  R-band)} afterglows that can be detected in an ideal survey of the whole sky with different limiting magnitudes. We considered jets without sideways expansion during the afterglow phase. This leads to a flatter light curve after {the jet-break time} $t_j$, compared with the sideways-expansion case. Thus, orphan afterglows can persist for a longer time (above a certain flux density), and more expected orphans can be detected. The distribution of half-opening angles of jets was also considered, which suppresses the number of optical orphans { compared with the case that all jets have one single half-opening} angle $\\theta_j=0.1$. When combining { these} two effects, the detectability is less than the canonical results of \\citet{npg02}, { who considered the case in which all jets have an initial half-opening angle $\\theta_j=0.1$ with sideways expansion.}  From Figs. (\\ref{fig:compare})-(\\ref{fig:p-e}), we can conclude that the main factors for the detectability are the total kinetic energy $E_j$, half-opening angle of jet $\\theta_j$ and temporal index $\\alpha_2$.  As $E_j$ is considered to be a standard energy \\citep{fks01}, our results {can be} used to constrain the value of the $E_j$ by detecting orphans. { If} $E_j$ and $\\alpha_2$ are determined accurately by other methods, the distribution function of the jet's { half-opening} angles { can} be determined well by observations of orphan afterglows.  However, our estimation is simplified in several aspects. First, the parameters are undetermined and diverse { in} different bursts. For a variable parameter, we should know its distribution. But this is very difficult. For example, the circum-burst environment seems to be an ISM, a wind, or another density-profile media, but these media cannot be determined clearly in well-observed GRBs. For simplicity, we choose the ISM with number density $n=1 \\rm{cm}^{-3}$. Second, orphan afterglows may have properties similar to other phenomena: e.g., failed gamma-ray burst \\citep{hdl02, r03, hlc05} and ``on-axis orphan afterglow'' \\citep{np03}. Third, we { use} the model with a power-law distribution of half-opening angles of uniform jets. However, there are some other structured jet models that cannot be ruled out \\citep{dg01,rlr02,zm02}. These models also affect the detectability. { Fourth, we assume that the GRB rate is proportional to the SFR, so our results are SFR-dependent.} Fifth, when observing, one should distinguish orphan afterglows from other transients carefully \\citep{lo02, bw04, go05}. {Finally, the dust grains within the jet's opening solid angle may be evaporated by the prompt UV/X-ray photons, and  the dust is possibly opaque  in optical/UV bands outside the jet cone, so an optical orphan afterglow may be generally suppressed.}  {\\it Acknowledgements.} We would like to thank the anonymous referee and Steven N. Shore for valuable suggestions, and E. Nakar for helpful discussions. This work was supported by the National Natural Science Foundation of China (grants 10233010, 10221001, and 10503012). XFW acknowledges the support from the China Postdoctoral Foundation, K.C.Wong Education Foundation (Hong Kong), and Postdoctoral Research Award of Jiangsu Province."
    },
    "0601/astro-ph0601547_arXiv.txt": {
        "abstract": "This article presents a measurement of the proper motion of the Sculptor dwarf spheroidal galaxy determined from images taken with the Hubble Space Telescope using the Space Telescope Imaging Spectrograph in the imaging mode.  Each of two distinct fields contains a quasi-stellar object which serves as the ``reference point.''  The measured proper motion of Sculptor, expressed in the equatorial coordinate system, is ($\\mu_{\\alpha}, \\mu_{\\delta})=(9 \\pm 13,2 \\pm 13)$~~mas~century$^{-1}$.  Removing the contributions from the motion of the Sun and the motion of the Local Standard of Rest produces the proper motion in the Galactic-rest-frame: $(\\mu_{\\alpha}^{\\mbox{\\tiny{Grf}}},\\mu_{\\delta}^{\\mbox{\\tiny{Grf}}}) = (-23 \\pm 13, 45 \\pm 13)$~~mas~century$^{-1}$.  The implied space velocity with respect to the Galactic center has a radial component of $V_{r}=79 \\pm 6$~km~s$^{-1}$ and a tangential component of $V_{t}=198 \\pm 50$~km~s$^{-1}$.  Integrating the motion of Sculptor in a realistic potential for the Milky Way produces orbital elements.  The perigalacticon and apogalacticon are 68~$(31,83)$~kpc and 122~$(97,313)$~kpc, respectively, where the values in the parentheses represent the $95\\%$ confidence interval derived from Monte Carlo experiments.  The eccentricity of the orbit is $0.29~(0.26,0.60)$ and the orbital period is $2.2~(1.5,4.9)$~Gyr.  Sculptor is on a polar orbit around the Milky Way: the angle of inclination is $86~(83,90)$~degrees. ",
        "introduction": "\\label{sec:intro}  Shapley (1938) discovered the Sculptor dwarf spheroidal (dSph) galaxy --- the first example of this type of galaxy in the vicinity of the Milky Way --- on a plate with a 3~hour exposure time taken with the Bruce telescope.  Shapely notes ``... that systems such as the Sculptor cluster may not be uncommon; their luminosity characteristics would enable them to escape easy discovery.''  Since the detection of Sculptor, astronomers have identified eight other dSphs.  Sculptor is at a celestial location of $(\\alpha,\\delta) = (01^{\\mbox{h}}00^{\\mbox{m}}09^{\\mbox{s}}, -33^{\\circ} 42^{\\prime}30^{\\prime\\prime})$ (J2000.0; Mateo 1998), which corresponds to Galactic coordinates of $(\\ell, b)=(287\\fdg 5,-83\\fdg 2)$.  Thus, Sculptor lies nearly at the South Galactic Pole.  Kaluzny \\etal\\ (1995) searched for variable stars in a $15^{\\prime} \\times 15^{\\prime}$ field centered approximately on the dSph by taking $V-$ and $I-$band images with the 1-m Swope telescope at Las Campanas Observatory over a period of more than two months. The search resulted in the identification of 226 RR~Lyr stars.  The average $V-band$ magnitude of the RR~Lyr stars gives a distance modulus of $(m-M)_{V}=19.71$, which corresponds to a heliocentric distance of 87~kpc.  This estimate is practically the same as that obtained by Hodge (1965); it is consistent with the estimate of Baade \\& Hubble (1939), but somewhat larger than the estimate of Kunkel \\& Demers (1977).  This study adopts the estimate of Kaluzny \\etal\\ (1995) for the distance to Sculptor.  Irwin \\& Hatzidimitriou (1995) derives the most comprehensive set of structural parameters for Sculptor --- and seven other dSphs --- using star counts from UK Schmidt telescope plates.  With a luminosity of $(1.4\\pm 0.6) \\times 10^{6}$~L$_{\\odot}$, Sculptor is among the most luminous dSphs.  Its major-axis core and limiting radii are $5.8 \\pm 1.6$~arcmin and $76.5 \\pm 5.0$~arcmin, respectively, which are in good agreement with the values derived by Demers, Kunkel, \\& Krautter (1980).  However, they differ from the values derived by Eskridge (1988a), who also uses star counts from UK Schmidt telescope plates. The discrepancy is likely due to an underestimation of the central density in the latter study for reasons that are discussed in Irwin \\& Hatzidimitriou (1995).  The isopleth map of Sculptor (Panel (f) of Figure~1 in Irwin \\& Hatzidimitriou 1995) shows that the ellipticity of the isodensity contours increases with increasing projected radius from the center of the dSph: the ellipticity is consistent with 0 in the inner 10~arcmin and smoothly increases to a value of 0.32 in the outermost region.  The study observes that Sculptor ``... looks remarkably similar to numerical simulations of dSph galaxies that are tidally distorted.'' The position angle of the major axis is $99 \\pm 1$~degrees.  Eskridge (1988b) finds asymmetric ``structure'' in Sculptor in an isopleth map of the difference between the stellar surface density and a fitted 2-D model.  In contrast, Irwin \\& Hatzidimitriou (1995) finds only the increase of the ellipticity with projected radius in a similar map.  The left panel in Figure~\\ref{fig:fields} shows a 30~arcmin $\\times$ 30~arcmin region of the sky centered on Sculptor.  The dashed ellipse is the boundary of the core.  Walcher \\etal\\ (2003) studies the structure of Sculptor --- together with those of Carina and Fornax --- using $V$-band images taken with the MPG/ESO 2.2-m telescope at La Silla.  The images have an areal coverage of 16.25~square degrees and reach a limiting magnitude of $V \\approx 23.5$.  The study derives major-axis core and limiting radii of $7.56 \\pm 0.7$~arcmin and $40 \\pm 4$~arcmin, respectively, using a King (1962) model, as did Irwin \\& Hatzidimitriou (1995).  The 1-$\\sigma$ disagreement between the two derived core radii is perhaps larger than expected from two data sets that have many stars in common.  The apparently more serious disagreement between the two limiting radii is most likely due only to a larger true uncertainty in the limiting radius caused by uncertainties in the background surface density and the poor fit of the model to the outer part of the surface density profile.  Walcher \\etal\\ (2003) confirms that the ellipticity of the surface-density contours increases with increasing projected radius; the contours also suggest ``extensions'' from the ends of the major axis that are interpreted as tidal tails.  The radial projected density profile shows a ``break'' --- a departure from the fitted King model --- at around 30~arcmin, which the study interprets as evidence for the existence of an extended stellar component.  Walcher \\etal\\ (2003) uses the relation between the King-model tidal radius, the mass, and the perigalacticon of a Galactic satellite developed by Oh, Lin, \\& Aarseth (1992) to deduce that Sculptor has a perigalacticon of 28~kpc.  A recent article by Coleman \\etal\\ (2005) does not confirm the finding in Walcher \\etal\\ (2003) that Sculptor has tidal tails and an extended stellar component.  Instead, the analysis of photometric data in the $V$ and $I$ bands for a $3.1^{\\circ}\\times3.1^{\\circ}$ field shows that a King model with a limiting radius of $72.5\\pm4.0$~arcmin is a satisfactory fit to the radial profile of stars that lie on the giant branch.  The limiting magnitudes of the photometry are $V=20$ and $I=19$, respectively.  The study notes that oversubtracting the field population in its data produces a radial profile with the smaller fitted limiting radius and significant extratidal structure found in Walcher \\etal\\ (2003).  Using additional information from the spectroscopy of 723 stars selected from the red giant branch, Coleman \\etal\\ (2005) derives an upper limit of $2.3\\% \\pm 0.6\\%$ for the contribution from stars beyond the tidal boundary to the total mass of Sculptor.  The study does not find any conclusive evidence for tidal interaction between the Milky Way and Sculptor.  In contrast, Westfall \\etal\\ (2005) does find evidence.  This study uses imaging in the $M$, $T_{2}$, and $DDO51$ bands to separate member giants from foreground dwarfs in a 7.82 degrees$^{2}$ area that covers the eastward side of Sculptor including the central region. Candidate members are also selected from the region of the blue horizontal branch in the color-magnitude diagram.  The selection of members is checked with spectroscopy for 147 candidates.  The study finds members up to 150 arcmin from the center of Sculptor --- the spatial extent of the survey and beyond the tidal boundary if that is identified with the measured King limiting radius of 80~arcmin. Several of these stars are spectroscopically-confirmed members.  The radial surface brightness profile shows a break to a shallower slope at a radius of about 60~arcmin, which resembles the radial profiles seen in simulations of satellites interacting with the Milky Way (Johnston \\etal\\ 1999).  Thus, Westfall \\etal\\ (2005) argues in favor of a significant tidal interaction between the Milky Way and Sculptor. It is beyond the scope of this work to resolve the apparent conflict between Coleman \\etal\\ (2005) and Westfall \\etal\\ (2005) by judging the merits of the analyses presented in both articles.  Needless to say, a disagreement exists about the effect of the Galactic tidal field on the structure of Sculptor; measuring the proper motion of the dSph may allow us to impose constraints on this effect.  Armandroff \\& Da Costa (1986) measured the radial velocities of 16 giants in Sculptor which average to produce a systemic heliocentric velocity of $107.4 \\pm 2.0$~km~s$^{-1}$.  This measurement alleviated the large uncertainty in this quantity, which existed due to mutually contradictory estimates from Hartwick \\& Sargent (1978) and Richter \\& Westerlund (1983).  More recently, Queloz, Dubath, \\& Pasquini (1995) measured the radial velocities of 23 giant stars.  The sample includes 15 stars observed previously by Armandroff \\& Da Costa (1986).  The implied systemic heliocentric velocity of $109.9 \\pm 1.4$~km~s$^{-1}$ (after excluding two stars that are likely binaries) agrees within the quoted uncertainties with the measurement of Armandroff \\& Da Costa (1986).  Our article adopts a mean velocity of $109.9 \\pm 1.4$~km~s$^{-1}$ for calculating the space velocity of Sculptor.  Queloz, Dubath, \\& Pasquini (1995) finds no apparent rotation of the dSph around its minor axis.  Tolstoy \\etal\\ (2004) measured radial velocities for 308 potential members of the dSph and find a systemic velocity of 110~km~s$^{-1}$.  There is no discussion of rotation, but Figure~4 in that article shows that the velocity dispersion does not increase with radius, as would be expected if there were a net rotation larger than the central dispersion.  Interestingly, the study finds that the red horizontal branch stars have a more compact spatial distribution and a smaller velocity dispersion than the older and more metal-poor blue horizontal branch stars.  The two most recent photometric and spectroscopic surveys by Coleman \\etal\\ (2005) and Westfall \\etal\\ (2005) confirm the greater central concentration of the more metal-rich stars.  The two studies also find no evidence for rotation.  The mass-to-light ratio, $(M/L)$, of Sculptor is larger than a typical value for a Galactic globular cluster; it is, however, smaller than the $M/L$s for some other Galactic dSphs.  Armandroff \\& Da Costa (1986) derives a central $M/L_{V}$ of $6.0 \\pm 3.1$ and Queloz, Dubath, \\& Pasquini (1995) determines the somewhat larger value of $13 \\pm 6$, in solar units.  Both studies note that the measured $M/L$ does not imply unequivocal support for dark matter in Sculptor.  The stars of Sculptor are old.  Fitting isochrones to the principal sequences in the color-magnitude diagram, Da Costa (1984) finds that the majority of the stars are younger by ``2--3~Gyr than Galactic globular clusters of similar metal abundance provided the helium abundances and the CNO/Fe ratios are also similar.''  Da Costa (1984) also detects ``blue stragglers'' and estimates their age to be about 5~Gyr under the assumption that they are ``normal'' main-sequence stars, i.e., stars which did not acquire mass from a companion.  No stars younger than 5~Gyr exist in Sculptor, indicating an absence of ongoing or recent star formation.  However, Sculptor contains HI gas.  Carignan \\etal\\ (1998) and later Bouchard \\etal\\ (2003) detect two distinct clouds of HI that are diametrically opposite to each other almost along the minor axis and 20 -- 30~arcmin from the center of the dSph.  These clouds are within the tidal radius.  The HI gas is very likely to be associated with Sculptor because its mean heliocentric velocity is similar to that of the dSph.  Carignan \\etal\\ (1998) discusses mechanisms that might account for the existence of the clouds.  Removing the gas from the dSph by a time-dependent tidal force due the Milky Way is one possibility. Carignan \\etal\\ (1998) suggests that the alignment between the proper motion vector from Schweitzer \\etal\\ (1995) and a line passing through the two clouds supports this hypothesis for the origin of the clouds. However, if the tidal force affects the HI, it should also affect the stars and the possible signatures of tides in the stellar component of Sculptor are either absent or inconsistent with the direction of the Schweitzer \\etal\\ (1995) proper motion vector.  For example, the increasing ellipticity of isodensity contours with increasing projected radius could be due to tides (e.g., Johnston, Spergel, \\& Hernquist 1995), but then the major axis should be along the proper motion vector.  The possible ``tidal extensions'' reported by Walcher \\etal\\ (2003) are at the ends of the major axis.  If these extensions are a continuation of the increasing ellipticity, they also argue for an orbital plane parallel to the major axis.  Walcher \\etal\\ (2003) claims that the eastern extension bends to the south, i.e., parallel to the minor axis and so argues that the orbital plane is aligned in the north-south direction.  However, this alignment is inconsistent with the increasing ellipticity being due to the tidal force since, given the large distance of Sculptor, the Sun is nearly in the orbital plane and so the increasing ellipticity should be aligned with the tidal extensions.  Schweitzer \\etal\\ (1995) reports the first measurement of the proper motion for Sculptor: $(\\mu_{\\alpha},\\mu_{\\delta})=(72 \\pm 22, -6 \\pm 25)$~mas~century$^{-1}$.  This value includes contributions from the motions of the Sun and LSR; this article refers to this quantity as the ``measured proper motion.\"  The measurement derives from 26 photographic plates imaged with a variety of ground-based telescopes using either a ``blue'', B, or V filter.  The earliest epoch is 1938 and the latest is 1991.  The study estimates, among other quantities, the perigalacticon of the implied orbit.  The best estimate ranges from 60~kpc for the ``infinite halo'' potential of the Milky Way to 78~kpc for the ``point mass'' potential.  If the perigalacticon is no smaller than 60~kpc, then the Galactic tidal force has not played a significant role in the evolution of Sculptor. Numerical simulations of Piatek \\& Pryor (1995) or Oh, Lin, \\& Aarseth (1995) show that for a typical dSph, even with a $M/L_V$ as low as 3, a perigalacticon of 60~kpc is too large for tides to have an important effect.  Motivated by the idea that some of the Galactic dSphs and globular clusters may be pieces of a tidally-disrupted progenitor satellite galaxy, several studies propose that they form ``streams'' in the Galactic halo.  Lynden-Bell (1982) hypothesizes that Fornax, Leo I, Leo II, and Sculptor are members of the ``FLS stream.'' Majewski (1994) adds the newly-discovered Sextans to the FLS stream and recalculates its common plane --- naming it the ``FL$^{2}$S$^{2}$ plane.''  The FLS and the FL$^{2}$S$^{2}$ planes differ only slightly. In a more extensive study, Lynden-Bell \\& Lynden-Bell (1995) infers that Sculptor may belong to one of three possible streams (see their Table~2).  Stream \\# 2 contains the LMC, SMC, Draco, Ursa Minor, and, possibly, Sculptor and Carina; stream \\# 4a contains Sextans, Sculptor, Pal 3, and, possibly, Fornax; finally, stream \\# 4b contains Sextans, Sculptor, and, possibly, Fornax.  For each stream, Lynden-Bell \\& Lynden-Bell (1995) calculates the expected proper motion of Sculptor.  Kroupa \\etal\\ (2004) notes that the 11 dwarf galaxies nearest to the Milky Way form a disk with a thickness to radius ratio of $\\leq$~0.15.  The article argues that the distribution expected for such nearby substructure in a cold-dark-matter universe is spherical, that the observed distribution is not, and, thus, that these objects are the tidal debris from the disruption of a larger satellite galaxy. In contrast, Kang \\etal\\ (2005) and Zentner \\etal\\ (2005) find that a planar distribution of nearby Galactic satellites is actually common in numerical simulations of galaxy formation.  A direct comparison between the results of the simulations and the distribution of nearby satellites finds that they are consistent.  A test of the reality of streams or planar alignments is to measure the space motions of the satellites.  Piatek \\etal\\ (2005; P05) reports a proper motion for Ursa Minor.  The implied orbit for Ursa Minor is not in the plane defined by Kroupa \\etal\\ (2004).  The proper motion also rules out membership in the stream proposed by Lynden-Bell \\& Lynden-Bell (1995).  The measured proper motion for Carina (Piatek \\etal\\ 2003; P03) does not agree well with the predictions of Lynden-Bell \\& Lynden-Bell (1995), but is not precise enough to rule out membership in a stream.  Piatek \\etal\\ (2002; P02) finds that a preliminary proper motion for Fornax is inconsistent with the predictions of Lynden-Bell \\& Lynden-Bell (1995) and that its direction is also inconsistent with an orbit in the FL$^{2}$S$^{2}$ plane.  Dinescu \\etal\\ (2004) reports an independent measurement of the proper motion of Fornax.  This motion is consistent, within its uncertainty, with the predictions of Lynden-Bell \\& Lynden-Bell (1995) and the direction of this motion is along the great circle defined by the FL$^{2}$S$^{2}$ plane.  Dinescu \\etal\\ (2004) notes that the proper motion for Sculptor in Schweitzer \\etal\\ (1995) is inconsistent with the FL$^{2}$S$^{2}$ plane.  This article reports a second independent measurement of the proper motion for Sculptor and discusses the implications of the derived space motion on the dSph-Galaxy interaction. Section~\\ref{sec:data} describes observations and the data.  The following section describes the analysis of the data leading to the derivation of the proper motion.  Section~\\ref{sec:pmm} compares the proper motion from Schweitzer \\etal\\ (1995) with the one reported in this article.  The next section, Section~\\ref{sec:orbit}, integrates and describes the orbit of Sculptor.  Section~\\ref{sec:disc} discusses the implications of the orbit for the importance of the Galactic tidal force on the structure and internal kinematics of Sculptor, for the star formation history, and for the membership of Sculptor in the proposed streams of galaxies and globular clusters in the Galactic halo.  The final section is a summary of the main results and conclusions. ",
        "conclusions": "\\label{sec:disc}  Knowing the orbit can help answer several questions about Sculptor, or, at least, increase the level of our understanding of this galaxy.  These questions are: 1. Is Sculptor a member of a stream of galaxies? 2. Is its star formation history correlated with the orbit?  3. What is the origin of the HI clouds detected in close proximity to the dSph?  4. Does Sculptor contain dark matter?  \\subsection{Is Sculptor a Member of a Stream?} \\label{stream}  Lynden-Bell \\& Lynden-Bell (1995) proposes that Sculptor may be a member of one of three possible streams: stream No.~2 (together with the LMC, SMC, Draco, Ursa Minor, and Carina); No.~4a (together with Sextans, Pal 3, and Fornax); or No.~4b (together with Sextans and Fornax).  Columns (2) and (3) in Table~6 give the predicted heliocentric (i.e., ``measured'' in our terminology) proper motion in the equatorial coordinate system for Sculptor if it indeed belongs to any of the three streams.  The magnitude of the proper motion vector, $\\vert \\vec{\\mu} \\vert = \\sqrt {\\mu_{\\alpha}^{2} + \\mu_{\\delta}^{2}}$, and its position angle are in columns (4) and (5).  For easy comparison, the corresponding quantities from our study are in the bottom line of the table.  Comparing the entries shows that the predictions for streams 2 and 4a disagree significantly with our measurement.  However, the prediction for stream No. 4b is closer: the magnitudes differ by 1.6$\\times$ the uncertainty in their difference, while the position angles differ by 1.6$\\times$.  Differences of this size should occur by chance 1\\% of the time.  The measured proper motion based on only the three-epoch data in the SCL~$J0100-3338$ field improves the agreement with the prediction for stream 4b.  Thus, while we rule out the possibility that Sculptor is a member of stream No.~2 or 4a, its membership in stream~4b is possible.  Stream 4b contains both Sculptor and Fornax.  The Dinescu \\etal\\ (2004) proper motion for Fornax is $(\\mu_\\alpha, \\mu_\\delta) = (59 \\pm 16, -15 \\pm 16)$~mas~cent$^{-1}$.  The magnitude and position angle of the proper motion are $61 \\pm 16$~mas~cent$^{-1}$ and $104\\pm 15$~degrees.  The prediction for stream 4b from Lynden-Bell \\& Lynden-Bell (1995) is 20~mas~cent$^{-1}$ and 162~degrees.  The difference between the measured and predicted proper motions would be this large or larger by chance only 0.4\\% of the time.  Thus, the physical reality of stream~4b is doubtful.  Dinescu \\etal\\ (2004) argues that Fornax and Sculptor are members of the same stream that also includes Leo I, Leo II, and Sextans.  Together the galaxies define the FL$^{2}$S$^{2}$ plane.  If they do form a stream, their Galactic-rest-frame proper motion vectors should be aligned with the great circle passing through the galaxies. The position angle of the great circle passing through Sculptor and Fornax is about 99~degrees at the location of Sculptor and 95~degrees at Fornax.  The position angle of the Galactic-rest-frame proper motion for Fornax reported by Dinescu \\etal\\ (2004) is $79 \\pm 25$~degrees, which differs by 0.64$\\times$ its uncertainty from the position angle of the great circle.  If Sculptor and Fornax form a stream, then they should move in the same direction along the great circle connecting them.  Thus, the proper motion of Fornax from Dinescu \\etal\\ (2000) implies that the position angle of the Galactic-rest-frame proper motion of Sculptor should be 99~degrees. The position angle for the proper motion of Sculptor reported here is $333 \\pm 15$~degrees, which differs from the prediction by 8.4$\\times$ its uncertainty.  Discounting the proper motion from this study and instead using the position angle of $40 \\pm 24$~degrees implied by the proper motion measured by Schweitzer \\etal\\ (1995) does not remove the disagreement: the position angle differs by $2.5\\times$ its uncertainty from that of the great circle.  We conclude that Sculptor and Fornax do not belong to the same stream.  Kroupa \\etal\\ (2004) shows that the 11 dwarf galaxies nearest to the Milky Way are nearly on a plane, whose two poles are at $(\\ell,b) = (168,-16)$~degrees and $(348,+16)$~degrees. Adopting the direction of the angular momentum vector as the pole of the orbit, then the location of the pole is \\begin{equation} (\\ell,b) = (\\Omega+90^{\\circ},\\Phi-90^{\\circ}). \\end{equation} Because of the left-handed nature of the Galactic rotation, prograde orbits have $b < 0$ and retrograde orbits have $b > 0$.  Thus, the pole of our orbit for Sculptor is $(\\ell,b) = (5 \\pm 16, -4 \\pm 1.6)$~degrees, where the uncertainties are 1-$\\sigma$ values from the Monte Carlo simulations.  The galactic longitudes of the poles of the plane and orbit agree within the uncertainty, but the galactic latitudes do not.  They differ by 20~degrees, which is more than 12$\\times$ the uncertainty in the location of the pole of the orbit. However, there is also some uncertainty in the orientation of the plane passing through the dwarf galaxies near the Milky Way.  We conclude the plane of the orbit of Sculptor is similar to the plane defined by the nearby dwarf galaxies.   \\subsection{The Effect of the Galactic Tidal Force on the Structure of Sculptor} \\label{sec:tides}  The measured ellipticity of the isodensity contours increases with projected distance from the center of Sculptor (see Figure~1 in Irwin \\& Hatzidimitriou 1995), akin to surface density contours of a model dSph in the numerical simulations of Johnston, Spergel, \\& Hernquist (1995; e.g., see Figure~4).  If the Galactic tidal force deformed Sculptor from an initial nearly-spherical shape to its present elongated shape in the outer regions then, from our vantage point nearly in the orbit plane, the position angle of its projected major axis should be similar to --- or differ by 180 degrees from --- the position angle of the Galactic-rest-frame proper motion vector, as predicted by the numerical simulations of Oh, Lin, \\& Aarseth (1995), Piatek \\& Pryor (1995), or Johnston, Spergel, \\& Hernquist (1995). The position angle of the projected major axis is $99 \\pm 1$~degrees and the position angle of our measured Galactic-rest-frame proper motion vector is $333 \\pm 15$~degrees.  Allowing for the 180-degree degeneracy, the difference between the two position angles is 3.6$\\times$ the uncertainty of their difference, which suggests that the Galactic tidal force has not elongated Sculptor.  \\subsection{Does Star Formation History Correlate with the Orbital Motion of Sculptor?} \\label{sec:sfh}  Da Costa (1984) imaged a 3~arcmin $\\times$ 5~arcmin field located just outside of the core radius of Sculptor in three bands: $B$, $V$, and $R$.  The photometry reaches the main-sequence turn off. Comparing theoretical isochrones with the distribution of stars in the color-magnitude diagram, the study finds that the majority of stars is about 2-3 Gyr younger than the galactic globular clusters of comparable metal abundance.  An earlier study by Kunkel \\& Demers (1977) based on $B$ and $V$ photometry extending to 0.4 magnitudes below the horizontal branch reaches a similar conclusion.  The color-magnitude diagram also shows a population of ``blue stragglers,'' which might be indicative of an extended period of star formation.  Da Costa (1984) concludes, however, that, if an intermediate-age stellar population exists in Sculptor, it is ``infinitesimal compared to that of the Carina system.''  Deep \\textit{HST} imaging by Monkiewicz \\etal\\ (1999) in a single field, reaching 3 magnitudes below the main-sequence turn-off, confirms the basic picture of Sculptor uncovered by Kunkel \\& Demers and Da Costa (1984).  This color-magnitude diagram also reveals the presence of ``blue stragglers'' and implies an age comparable to that of the galactic globular clusters.  The small number of stars in the small field made a search for an intermediate-age stellar population inconclusive.  Majewski \\etal\\ (1999), Hurley-Keller \\etal\\ (1999), and Harbeck \\etal\\ (2001) use wide-field imaging to show the presence of two stellar populations with distinctly different metallicities ([Fe/H] = --2.3 and --1.5; Majewski \\etal\\ 1999).  The more metal rich population is more centrally concentrated in the galaxy.  Most recently, Tolstoy \\etal\\ (2004) confirms the above picture using wide-field imaging and spectroscopy.  Spectroscopically-determined metallicities range from --2.8 to --0.9.  Stars more metal-rich than --1.7 are more centrally concentrated and have a smaller velocity dispersion than the rest of the sample.  However, both stellar populations are older than 10~Gyrs.  The aforementioned studies show that there were at least two episodes of star formation at times more than 10~Gyr ago.  Because 10~Gyrs is much longer than the orbital period of approximately 2.2~Gyr, there is no clear connection between the stimulation of star formation and processes such as the Galaxy-Sculptor tidal interaction or the effects of ram pressure.  The lack of correlation could be due to the loss of all of the gas in Sculptor about 10~Gyr ago.  But, surprisingly, the observations indicate that Sculptor has HI today.  \\subsection{HI Gas in Sculptor} \\label{sec:gas}  Unlike most other Galactic dSphs, Sculptor contains a detectable amount of HI.  Knapp \\etal\\ (1978) detects three clouds of HI in the vicinity of Sculptor and speculates that one of them, with a radial velocity of 120~km~s$^{-1}$, may be associated with Sculptor, whose radial velocity at the time was uncertain.  Carignan \\etal\\ (1998) confirms and refines this detection and puts a lower limit on the mass of HI of $3.0 \\times 10^{4}$~M$_{\\odot}$.  Bouchard \\etal\\ (2003) repeats the observations over a wider field with the Parkes single-dish telescope and, over a smaller region, at higher angular resolution with the Australia Telescope Compact Array.  The better data show that the HI is not associated with a background galaxy and the probability of a chance superposition of a galactic high velocity cloud is less than 2\\%.  These arguments, together with the agreement within 4~km~s$^{-1}$ of the radial velocity of the HI and the radial velocity of Sculptor make a strong case for the physical association of the clouds with the dSph.  The gas is in two clouds.  Figure~\\ref{fig:gas}\\ shows the distribution of HI on the sky in the direction of Sculptor based on the Australia Telescope Compact Array data from Bouchard \\etal\\ (2003).  The asterisk marks the optical center of Sculptor and the two clouds are about 20--30~arcmin from the center, one to the northeast and one to the southwest.  The masses of the clouds are $(4.1 \\pm 0.2) \\times 10^4$~M$_\\odot$ and $(1.93 \\pm 0.02) \\times 10^5$~M$_\\odot$, respectively.  The two clouds lie nearly along the minor axis of the dSph; the orientation of Sculptor is shown in Figure~\\ref{fig:gas}\\ by the ellipse representing the optical boundary.  Although the association between the gas and the dSph seems well-established, why Sculptor still has gas when most of the other galactic dSphs do not and the cause of the observed configuration of the gas are debated.  Mayer \\etal\\ (2005) studies the loss of gas from dwarf galaxies in the Local Group using numerical simulations that include ram pressure stripping.  It quantifies the expected result that a galaxy with a deeper potential well or with a larger perigalacticon is more likely to retain its gas.  The orbit of Sculptor is similar to those of Carina and Ursa Minor, which suggests that differences in the degree of tidal shocking or ram pressure stripping are not the reason for the difference in gas retention.  However, note that the 95\\% confidence intervals for the perigalacticons of all three orbits are still too large to make a conclusive statement.  The larger mass of Sculptor compared to Carina and Ursa Minor is the most likely reason why it was able to retain gas.  Mechanisms that could affect the distribution of HI within the dSph are tidal interaction, ram pressure, and forces from supernovae or winds from young stars.  Tidal interaction and ram pressure tend to spread the gas in the plane of the orbit.  The two arrows in Figure~\\ref{fig:gas}\\ represent the Galactic-rest-frame proper motions as measured by Schweitzer \\etal\\ (1995; dashed) and this study (solid).  The ratio of their lengths is the same as the ratio of the magnitudes of the proper motions.  The dotted line is a section of a great circle that passes through Sculptor and Fornax; its position angle is 99~degrees.  The Schweitzer \\etal\\ (1995) proper motion (dashed arrow) is nearly aligned with the line that connects the centers of the clouds. Carignan \\etal\\ (1998) notes this alignment and suggests that it may indicate a ``tidal'' origin for the clouds: presumably, the Galactic tidal force stretches the gas, initially centered within the dSph, into its observed distribution.  The most serious problem with such a picture is that the tidal force would stretch the distribution of both the stars and the gas, which then aligns the major axes of the gaseous and stellar distributions from our perspective as an observer nearly in the plane of the orbit of Sculptor.  They are not aligned, perhaps because the motion of the gas is governed by both gravitational forces and pressure gradients.  As was noted in Section~\\ref{sec:tides}, if the observed ellipticity of Sculptor is due to the tidal force, then the Galactic-rest-frame proper motion should be aligned with the major axis.  Figure~\\ref{fig:gas}\\ shows that the solid arrow, our measurement, is closer to such an alignment than the dashed arrow. Because the tidal force is zero at the center of a dSph, it cannot by itself separate a single cloud centered on the dSph into two clouds. Thus, within the context of a picture in which tides have had an important effect on Sculptor, our proper motion is more plausible than that of Schweitzer \\etal\\ (1995).  However, is our proper motion consistent with the geometry of the two clouds?  We think yes.  First, the two clouds are elongated in the direction of our proper motion (see particularly Figure~1 in Bouchard \\etal\\ 2003).  The observed elongation could be due to ram pressure from the motion through a gaseous Galactic halo.  Second and more speculatively, the two clouds could be due to the Rayleigh-Taylor instability when the HI in Sculptor moves through the hot and low-density gaseous halo.  Or the gas could have been squeezed out perpendicular to the direction of motion by the compressive tidal shock when Sculptor crosses the Galactic disk.  Expansion combined with infall of the gas forms a ring that looks like two clouds in projection.  \\subsection{Is there dark matter in Sculptor?} \\label{sec:dm}  Estimates of the $M/L_V$ for Sculptor range from $6 - 13$ and estimates of the limiting radius range from about $40 - 80$~arcmin. The estimates of $M/L$ assume that mass follows light.  If this is true, then the implied mass of Sculptor must be large enough, given our orbit, to produce a tidal radius that is at least as large as the observed limiting radius.  Equating the tidal radius and the limiting radius predicts a value for $M/L$, which should agree with the measured value.  Also, the dSph must have a mass and, hence, $M/L$ large enough for it to have survived destruction by the Galactic tidal force on our orbit.  In lieu of numerical simulations, an approximate analytical approach is to calculate the tidal radius, $r_t$, beyond which a star becomes unbound from the dSph.  For a logarithmic Galactic potential, $r_t$ is given by (King 1962; Oh, Lin, \\& Aarseth 1992) \\begin {equation} r_t = \\left(\\frac{(1-e)^2}{[(1+e)^2/2e]\\ln[(1+e)/(1-e)] +1} \\, \\frac{M}{M_G}\\right)^{1/3} a. \\label{eq:rtidal} \\end {equation} Here $e$ is eccentricity of the orbit, $a$ is the semi-major axis ($a \\equiv (R_{a}+R_{p})/2$), $M$ is the mass of the dSph, and $M_G$ is the mass of the Galaxy within $a$.  Equating $r_t$ with the observed limiting radius derived by fitting a King (1966) model, $r_k$, gives a value for $M/L_V$ for a given orbit.  If $r_k = 40$~arcmin, then 28\\% of the orbits in Monte Carlo simulations have $M/L_V > 6$ and 10\\% of the orbits have $M/L_{V} > 13$. If $r_k = 80$~arcmin, then 100\\% of the orbits have $M/L_V > 13$.  These results show that the global $M/L$ of Sculptor is probably larger than the measured $M/L$, if the larger of the measured limiting radii is identified with the tidal radius.  The $M/L$ calculated assuming that mass follows light underestimates the true global $M/L$ if Sculptor contains dark matter that is more spatially extended than the luminous matter (e.g., Pryor \\& Kormendy 1990).  However, equation~\\ref{eq:rtidal}\\ shows that $M \\propto r_t^3$, so the values for $M/L$ derived using this equation are sensitive to the measured value of the limiting radius.  Until kinematic measurements definitively identify the tidal radius, an $M/L$ derived with the above argument should be treated with caution.  The average of the measured values of $M/L_V$ for Galactic globular clusters is 2.3 (Pryor \\& Meylan 1993).  Could the true $M/L_V$ of Sculptor be similar to this average?  Numerical simulations by Oh, Lin, \\& Aarseth (1995) and Piatek \\& Pryor (1995) show that the ratio of the limiting radius derived by fitting a theoretical King model (King 1966), $r_k$, to the tidal radius defined by Equation~(\\ref{eq:rtidal}) is a useful indicator of the importance of the Galactic tidal force on the structure of a dSph.  These simulations show that: if $r_{k}/r_t \\lesssim 1.0$, the Galactic tidal force has little effect on the structure of the dSph; at $r_k/r_t \\approx 2.0$, the effect of the force increases rapidly with increasing $r_k/r_t$; and, for $r_k/r_t \\approx 3.0$, the dSph disintegrates in a few orbits.  Assuming that $M/L_V = 2.3$ and $r_k = 40$~arcmin, $r_k/r_t > 2.0$ for 6\\% of the orbits generated in Monte Carlo simulations.  If $r_k = 80$~arcmin, the fraction is 100\\%.  Thus, it is possible that Sculptor could have survived for a Hubble time on its current orbit if it only contains luminous matter.  We conclude that measured orbit of Sculptor does not require it to contain dark matter.  \\subsection{A Lower Limit for the Mass of the Milky Way} \\label{sec:massofg}  Sculptor is bound gravitationally to the Milky Way.  The Galactocentric space velocity of the dSph imposes a lower limit on the mass of the Milky Way within the present Galactocentric radius of the dSph, $R$.  Assuming a spherically symmetric mass distribution and zero for the total energy of the dSph, the lower limit for the mass of the Milky Way is given by \\begin{equation} \\label{mwmass} M=\\frac{R\\left (V_{r}^{2} + V_{t}^{2} \\right )}{2G}. \\end{equation} Setting $R=87$~kpc and using the values from Table~4 for $V_{r}$ and $V_{t}$, $M = (4.6\\pm 2.0) \\times 10^{11}\\ M_{\\odot}$. This lower limit is consistent with other recent estimates of the mass of the Milky Way, such as the mass of $5.4^{+0.1}_{-0.4} \\times 10^{11}\\ M_{\\odot}$ within $R=50$~kpc found by Sakamoto, Chiba, \\& Beers (2003).  The Milky Way potential adopted in Section~\\ref{sec:orbit} has a mass of $7.8\\times 10^{11}\\ M_{\\odot}$ out to $R=87$~kpc."
    },
    "0601/hep-ph0601134_arXiv.txt": {
        "abstract": "\\noindent The WMAP results on the scalar spectral index $n$ and its running with scale, though preliminary, open a very interesting window to physics at very high energies. We address the problem of finding inflaton potentials well motivated by particle physics which can accomodate WMAP data. We make a model independent analysis of a large class of models: those with flat tree-level potential lifted by radiative corrections, which cause the slow rolling of the inflaton and the running of $n$. This includes typical hybrid inflation models. In the small-coupling regime the predictions for the size and running of $n$  are remarkably neat, {\\it e.g.} $-dn/d\\ln k=(n-1)^2\\ll 1$, and $n$ does not cross $n=1$, contrary to WMAP indications. On the other hand, $n$ can run significantly if the couplings are stronger but at the price of having a small number of e-folds, $N_e$. We also examine the effect of mass thresholds crossed during inflation. Finally, we show that the presence of non-renormalizable operators for the inflaton, suppressed by a mass scale above the inflationary range, is able to give both $dn/d\\ln k\\sim {\\cal O}(-0.05)$~and~$N_e\\sim 50$. ",
        "introduction": "The measurement and analysis of the cosmic microwave background (CMB) by the WMAP collaboration \\cite{WMAP,WMAP1,WMAP2} has provided results of unprecedented accuracy on the spectrum of the CMB anisotropies. The data are basically consistent with the main features of inflationary cosmology \\cite{WMAP2}, implying important constraints for particular models (many of which can be discarded). One of the most intriguing WMAP results, though still uncertain, indicates that the spectral index of scalar perturbations, $n$, may exhibit a significant running with the wave-number, $k$.  More precisely, assuming a linear dependence of $n$ with $\\ln k$, $n(k)=n(k_0)+[d n /d \\ln k]\\ln(k/k_0)$ ($k_0$ is some chosen ``pivotal scale'', $k_0=0.002\\ {\\rm Mpc}^{-1}$), the analysis of ref.~\\cite{WMAP2} gives $n(k_0)=1.13\\pm 0.08$ and $d n /d \\ln k = -0.055^{+0.028}_{-0.029}$ at 68\\% C.L. (These values are obtained from WMAP data combined with different large-scale-structure data.) This means in particular that $n-1$ passes from positive to negative values in the range of $k$-values corresponding to these data ($10^{-4}\\simlt k\\ {\\rm Mpc} \\simlt 4$). It should be said, however, that a fit assuming constant spectral index is also consistent with the data, although somewhat worse, giving $n=0.99\\pm 0.04$ (with WMAP data only \\cite{WMAP}).  As a matter of fact, most inflationary models \\cite{infl} do predict that $n$ runs with $k$, but the details of the ``observed'' running are very difficult to fit \\cite{chung}. The last scales entering the observable universe (which we will denote $k_*\\sim 10^{-4}\\ {\\rm Mpc}^{-1}$) should have crossed the horizon $\\sim 50-60$ e-folds before the end of inflation, at a time we will call $t_{*}$. The range of $k$ probed by the WMAP data corresponds to the first $\\sim 6.5$ e-folds after $t_{*}$ (this range extends to $\\sim 10.5$ e-folds if one adds large-scale-structure data).  During this ``initial'' period of inflation, $n-1$ should change sign (if the WMAP indication is confirmed), and it is not easy to cook up an inflationary potential implementing this feature (see \\cite{chung,other} for an incomplete list of several analyses along these lines).  In the present paper we show, however, that a large and very interesting class of inflationary potentials, including those of many supersymmetric hybrid inflation models, are in principle able to accommodate this behavior quite naturally.  Let us recall a few basic results of inflationary cosmology in the slow-roll approximation. The slow-roll parameters are given by \\bea \\label{SlowRoll} \\epsilon={1\\over 2}M_p^2 \\left({V'\\over V}\\right)^2\\ , \\;\\;\\;\\;\\;\\; \\eta=M_p^2 {V''\\over V} \\ ,\\;\\;\\;\\;\\;\\; \\xi=M_p^4 {V'V'''\\over V^2} \\ , \\eea where $8\\pi M_p^2 = 1/G_{\\rm Newton}$, $V$ is the scalar potential of the (slow-rolling) inflaton field, $\\phi$, and the primes denote derivatives with respect to $\\phi$. The slow-roll approximation requires $\\epsilon, |\\eta|, |\\xi| \\ll 1$, and the failure of this condition marks the end of inflation. The number of e-folds produced during the slow-rolling is given by \\bea \\label{Ne} N_e = \\int_{t}^{t_{\\rm end}}H dt \\simeq {1\\over M_p^2}\\int_{\\phi_{\\rm end}} ^\\phi {V\\over V'} d\\phi = {1\\over M_p}\\int_{\\phi_{\\rm end}} ^\\phi {1\\over \\sqrt{2\\epsilon}} d\\phi \\ , \\eea where $H$ is the Hubble parameter [with $H^2=V/(3M_p^2)$], meaning that the scale parameter of the universe, $a$, grows as $a\\rightarrow e^{N_e}a$. Successful inflation requires\\footnote{It has been argued, however, that the required $50-60$ e-folds could be achieved through several steps of inflation, corresponding to different inflaton fields and/or inflaton potentials \\cite{multinflation}. WMAP would only be sensitive to the first(s) of these steps.} $N_e\\geq 50-60$. At first order in the slow-roll parameters \\cite{makarov}, the power spectrum of scalar perturbations, $P_k$ (called $\\Delta_{\\cal R}(k)=(2.95\\times 10^{-9})A(k)$ in ref.~\\cite{WMAP2}), and  the spectral index, $n$, are given by \\bea \\label{Pk} P_k = {1\\over 24\\pi^2\\epsilon}{V\\over M_p^4}\\ , \\eea \\bea \\label{n} n\\equiv 1 + {d \\ln P_k \\over d \\ln k}\\simeq 1+2\\eta-6\\epsilon\\ . \\eea Let us remind that, in the slow-roll approximation, the derivatives with respect to $\\ln k$ can be related to field derivatives by using \\bea \\label{dlnk} {d \\phi \\over d \\ln k}=-M_p^2{V'\\over V}= -M_p\\sqrt{2\\epsilon}\\ , \\eea as is clear from (\\ref{Ne}) and $a\\propto 1/k$. It is then trivial to obtain \\be \\label{dndlnk} {d n\\over d\\ln k}\\simeq -2\\xi +16\\epsilon\\eta-24\\epsilon^2\\ . \\ee  The WMAP analysis \\cite{WMAP1,WMAP2} gives \\bea \\label{Pk0} P_k=(2.95\\times 10^{-9})\\times(0.70^{+0.10}_{-0.11})\\;\\;\\;\\;  {\\rm at}\\;\\;\\; k=k_0\\equiv 0.002\\ {\\rm Mpc}^{-1} \\ , \\eea and $n$ as mentioned in the first paragraph: \\be n(k_0)=1.13\\pm 0.08\\ ,\\;\\;\\;\\; {d n \\over d \\ln k} = -0.055^{+0.028}_{-0.029}\\ . \\ee Since usually $\\epsilon\\ll \\eta$, then $n\\simeq 1+2\\eta$, and the change of sign in $n-1$ must arise from a change of sign in $\\eta$ (which implies in turn that $V'$ must pass through an extremal point), something which, as mentioned above, is non-trivial to implement (see ref.~\\cite{chung} for a systematic analysis of this and related issues).   On the other hand, the evidence for a running $n$ comes from the first ($l<5$) WMAP multipoles \\cite{bridle}, which corresponds to the first $\\sim 1.5$ e-folds of inflation (after $t_{*}$). Excluding those points, the fit is consistent with a constant index, $n\\simeq 0.94\\pm 0.05$, {\\it i.e.} the most significant running of $n$ seems to be associated to the first few e-folds and then $n$ remains quite stable. In any case, notice that the errors associated to these analyses are big (actually refs.~\\cite{WMAP2} and \\cite{bridle} are only consistent within these errors).  Hence the results concerning the spectral index can be summarized by saying that they indicate \\bea \\label{n1} {d n \\over d \\ln k}= {\\cal O}(-0.05)\\ , \\eea during the first few e-folds of inflation after $t_{*}$, and \\bea \\label{n2} n\\simeq 0.94\\pm 0.05\\ , \\eea during the following e-folds. Of course, if the previous dramatic running is not confirmed by forthcoming analyses, $n$ could still exhibit a milder running, compatible with (\\ref{n2}) in the whole range of observed scales.   In section~2 we study a large class of inflationary models defined by having flat tree-level potential, with radiative corrections being responsible for the slow-roll of the inflaton and the running of $n$. The results depend on whether the couplings are weak (subsection~2.1) or strong (subsection~2.2) but generically one cannot get both a significant running of $n$ and a large $N_e$. We devote section~3 to a numerical analysis of this issue in the well motivated type of $D$-term inflation models.  Things improve in this respect if inflation is affected by some high energy threshold, possibility which is studied in section~4. Subsection~4.1 studies what happens if the inflaton crosses such high-energy threshold shortly after $t_*$ and subsection~4.2 analyzes the most promising case in which the threshold is above the inflationary range but affects inflation through non-renormalizable operators. In this last case we can easily accomodate both $dn/d\\ln k\\sim {\\cal O}(-0.05)$ and $N_e\\sim 50-60$. Finally, Appendix~A gives some technical details relevant to section~3 and subsection 4.1. ",
        "conclusions": "In this paper we have assumed that the WMAP indications on the running of the scalar spectral index are real and then explored the intriguing possibility that this runnig is caused by physics at very high energy scales affecting the inflaton potential. It is tantalizing to imagine that we could be seeing in the sky the imprints of such physics. In more concrete terms we aim for well motivated inflaton potentials (from the point of view of particle physics) that can reproduce the suggested slope of $n$ with scale, $dn/d\\ln k\\sim {\\cal O}(-0.05)$, giving at the same time the right amount of e-folds and the right amplitude of the power spectrum of scalar fluctuations.  This is not an easy task and has atracted some effort recently. We have focused on a large class of inflationary models which we find particularly atractive: those with flat tree-level potential lifted by radiative corrections. These radiative corrections directly control the slow roll of the inflaton and, in turn, the running of the spectral index $n$. These models include some typical hybrid inflation models as a particularly interesting subclass. In these models there is a one-to-one correspondence between the value of the inflaton potential, the related distance scale $k$ and the associated renormalization scale, so that the running of the inflaton down its potential scans different high energy scales and cosmological distances. In this language, WMAP indicates that the spectrum of scalar perturbations is blue ($n>1$) at the largest cosmological scales ($k\\sim 10^{-4}{\\rm Mpc}^{-1}$) and turns to red ($n<1$) below some scale $k\\sim 10^{-2}{\\rm Mpc}^{-1}$. In terms of energy scales, at the start of inflation the inflaton had a very large value (corresponding to very high energy scales) and its slow-roll produced $n>1$ and a negative slope $dn/d\\ln k$. After a few e-folds this negative slope made $n$ cross 1 at some particular value $\\phi_0$ of the inflaton field, corresponding to some particular energy scale $Q_0$. Later on, $n$ gets stabilized around $0.94-0.99$ and the additional number of e-folds accumulates till $N_e\\sim 50-60$, when inflation ends.  In our search for suitable potentials in this type of models we find two different regimes depending on how sizeable the coupling constants of the model parameters are. In the small-coupling regime, these models make sharp predictions for the size of the spectral index $n$ and for its running with scale, namely \\be - {dn\\over d\\ln k}=(n-1)^2\\ll 1\\ . \\ee From this equation, it is clear that the preliminary WMAP indication of $n$ crossing $n=1$ cannot be reproduced.  On the other hand, at larger coupling we find that a much stronger running of $n$ is absolutely natural and even unavoidable. However, reproducing the WMAP indication for $dn/d\\ln k$ {\\em and} a sufficient number of e-folds ($N_e\\sim 50-60$) simultaneously is very difficult. We have illustrated this in a popular $D$-term hybrid inflation model for which we have obtained a `no-go' constraint of the form (\\ref{rule}) \\be (N_e^0)^2 \\left.{dn\\over d\\ln k}\\right|_{Q_0}\\simeq -1.1\\ , \\label{concl2} \\ee where, as mentioned above, $Q_0$ is the scale at which $n$ crosses 1, and $N_e^0$ is the number of e-folds from that scale till the end of inflation. Now, if one requires $dn/d\\ln k\\simeq -0.05$, then $N_e^0\\ll 50$. Therefore, unless a subsequent stage of inflation completes the total number of e-folds, these models are in trouble by themselves to reproduce the behaviour of $n$ mentioned above.  As the details of inflation in this region of scales are sensitive to possible physics thresholds, we considered the natural possibility of the inflaton coupling to heavy new fields with a mass $M$ that falls in the energy range crossed by the inflaton field in the first few e-folds. Somewhat to our surprise we found that, in the presence of sharp thresholds, the results of the previous case are not changed substantially and still one finds that the requirement of a significant running of $n$ leads to a small number of e-folds being produced.  A different way in which high energy physics can leave an imprint in the predictions of inflation is through non-renormalizable operators that modify the inflaton potential in the inflationary energy range . In this case, the new energy  threshold is not crossed during inflation (this is necessary for a reliable treatment of new physics effects in terms of effective operators in the first place!). The key point here is that even operators that are suppressed by a high energy mass scale can become important in the environment of a very flat potential. In this context we are able to write down very simple and well motivated inflationary potentials that can give a running spectral index in agreement with WMAP and a sufficient number of e-folds for reasonable choices of the parameters. The succesful type of potential combines three ingredients: tree-level flatness, smooth dependence on the inflaton through radiative logarithms and a non-renormalizable term: \\be V(\\phi)=\\rho+\\beta \\ln{m(\\phi)\\over Q}+\\phi^4{\\phi^{2N}\\over M^{2N}}\\ . \\label{NROconcl} \\ee We have further motivated this type of potentials by finding supersymmetric realizations through superpotentials that include non-renormalizable operators (NROs) of modest order.  Finally, we summarize the functional dependence of $n(k)$ expected for flat tree-level potentials. In the small-coupling regime, and if no mass thresholds are crossed during inflation, this is simply given by \\bea \\label{smallrunning2concl} n(k) =1 -{1\\over N_e - \\ln (k/ k_*)} \\ , \\eea where $k_*$ is the initial distance scale $k_*\\sim 10^{-4}\\ {\\rm Mpc}^{-1}$. The only parameter in (\\ref{smallrunning2concl}), $N_e$, corresponds typically to the total number of e-folds, i.e. $N_e\\simeq 50-60$, although it could be either smaller [if there are subsequent episodes of inflation], or larger --or even negative-- [if the end of inflation is not marked by $\\eta={\\cal O}(1)$, but by the violation of some hybrid-inflation condition for the tree-level flatness]. This is explained in subsect.~2.1. The case with thresholds and/or not-so-small couplings are addressed in subsects.~2.1,~2.2,~4.1.   On the other hand, when there are effects of high-energy (new) physics through non-renormalizable operators (NROs), $n(k)$ is of the form \\bea \\label{nintegratedNconcl} n (k)& = & n_* + {1\\over N_e} - {1\\over N_e - \\ln (k/ k_*)}\\nonumber\\\\ &-&{N_e\\over (N+1)}\\left(\\left.{d n\\over d\\ln k}\\right|_*+{1\\over N_e^2} \\right)\\left[\\left(1- {1\\over N_e}\\ln{k\\over k_*} \\right)^{N+1}-1\\right]\\ , \\eea where $*$ denotes quantities evaluated at the initial scale $k_*$, and $N$ is the exponent of the dominant NRO, as expressed in eq.~(\\ref{NROconcl}). The quantity $N_e$ has the same meaning as in eq.~(\\ref{smallrunning2concl}). The other parameters in eq.~(\\ref{nintegratedNconcl}) are $n_*$, $d n/d\\ln k |_*$, which can be expressed in terms of the parameters of the inflaton potential (\\ref{NROconcl}), as explained in subsect.~4.2. Finally, if the effect of the NROs is negligible [e.g. if $M\\rightarrow\\infty$ in (\\ref{NROconcl}), which corresponds to $n_*\\rightarrow 1-1/N_e$, $d n/d\\ln k |_*=-1/ N_e^2$ in~(\\ref{nintegratedNconcl})], then eq.~(\\ref{smallrunning2concl}) is recovered. For a more detailed discussion, see subsect.~4.2.  It would be a very nice test for the wide and well-motivated class of inflation models analyzed in this paper, if future analyses and measurements of the spectral index $n(k)$ could be adjusted by one of the previous expressions, which in turn would provide precious information about the inflaton potential and on  physics at very high energy-scales.   \\subsection*{A. The simple hybrid-inflation model with N+1 flavours} \\setcounter{equation}{0} \\renewcommand{\\theequation}{A.\\arabic{equation}}   We consider here the same hybrid-inflation model of sect.~3, with $N+1$ pairs of $H_\\pm$ fields, instead of the unique pair (i.e. $N=0$) of the original model. The results have interest on their own (e.g. they illustrate how the values of the coupling constants and the speed of running are affected by the number of flavours), but they are also useful to illustrate the threshold-crossing procedure discussed in sect.~4.1.  The superpotential reads now \\bea \\label{WN} W=\\Phi(\\sum_{a=1}^{N+1}\\lambda_a H_+^aH_-^a-\\mu^2)\\ , \\eea where $\\lambda_a$ are the $N+1$ Yukawa couplings. As for the $N=0$ case, $H_\\pm^a$ have charges $\\pm 1$ with respect to the $U(1)$ gauge group with gauge coupling $g$ and Fayet-Iliopoulos term $\\xi_D$. The tree-level scalar potential is given by $V_0=V_F +V_D$, with \\bea \\label{VFVDN} V_F&=&\\left|\\sum_{a}\\lambda_a H_+^aH_-^a-\\mu^2\\right|^2 + \\sum_{a}\\lambda_a^2\\left(\\left| H_-\\right|^2 + \\left|H_+\\right|^2\\right)\\left|\\Phi\\right|^2\\ ,\\nonumber\\\\ V_D&=&{g^2\\over 2}\\left[\\sum_{a}\\left(\\left|H_+^a\\right|^2 - \\left|H_-^a\\right|^2 \\right) + \\xi_D\\right]^2\\ . \\eea Again, the global minimum of the potential is supersymmetric ($V_F=0$, $V_D=0$), but for large enough $|\\Phi|$ the potential has a minimum at $H_\\pm^a=0$ and is flat in $\\phi$. The tree-level potential in this region is \\bea \\label{rhoBDN} \\rho=\\mu^4 + {1\\over 2}g^2\\xi_D^2\\ . \\eea The various $\\beta$-functions, defined as derivatives with respect to $\\ln Q$, are given by \\bea \\label{betaeN} \\beta_g&=&{1\\over 8\\pi^2}(N+1) g^3\\ ,\\\\ \\label{betaxiN} \\beta_{\\xi_D}&=& 0\\ ,\\\\ \\label{betagN} \\beta_{\\lambda_a}&=&{1\\over 16\\pi^2}\\lambda_a\\left[3\\lambda_a^2 +\\sum_{b\\neq a} \\lambda_b^2 - 4g^2\\right]\\ ,\\\\ \\label{betagvN} \\beta_{\\mu^2}&=&{1\\over 16\\pi^2} \\sum_{a=1}^{N+1} \\lambda_a^2 \\mu^2 \\ ,\\\\ \\label{betaBDN} \\beta_\\rho&=&{1\\over 8\\pi^2}\\left[(N+1)g^4\\xi_D^2+\\sum_{a=1}^{N+1} \\lambda_a^2 \\mu^4\\right]\\ . \\eea  For simplicity we can take the same Yukawa coupling $\\lambda_a=\\lambda$ for all $\\{H_+^a, H_-^a\\}$ pairs, which is a stable condition under RG evolution."
    },
    "0601/astro-ph0601079_arXiv.txt": {
        "abstract": "{The metallicity in portions of high-redshift galaxies has been successfully measured thanks to the gas observed in absorption in the spectra of quasars, in the Damped Lyman-$\\alpha$ systems (DLAs). Surprisingly, the global mean metallicity derived from DLAs is about 1/10$^{\\rm th}$ solar at 0$\\la$z$\\la$4 leading to the so-called ``missing-metals problem''.  In this paper, we present high-resolution observations of a sub-DLA system at \\zabs=$0.716$ with super-solar metallicity toward SDSS J1323$-$0021. This is the highest metallicity intervening quasar absorber currently known, and is only the second super-solar absorber known to date. We provide a detailed study of this unique object from VLT/UVES spectroscopy. We derive [Zn/H]=$+$0.61, [Fe/H]=$-$0.51, [Cr/H]=$<-$0.53, [Mn/H] = $-$0.37, and [Ti/H] = $-$0.61.  Observations and photoionisation models using the CLOUDY software confirm that the gas in this sub-DLA is predominantly neutral and that the abundance pattern is probably significantly different from a Solar pattern. Fe/Zn and Ti/Zn vary among the main velocity components by factors of $\\sim 3$ and $\\sim 35$, respectively, indicating non-uniform dust depletion. Mn/Fe is super-solar in almost all components, and varies by a factor of $\\sim 3$ among the dominant components. It would be interesting to observe more sub-DLA systems and determine whether they might contribute significantly toward the cosmic budget of metals. ",
        "introduction": "Damped Lyman-$\\alpha$ systems (DLAs) seen in absorption in the spectra of background quasars are selected over all redshifts independent of the intrinsic luminosities of the underlying galaxies. They have hydrogen column densities, \\loghi\\ $\\ga$ 20.3 and are the major contributors to the neutral gas in the Universe at high redshifts (Storrie-Lombardi \\& Wolfe 2000; P\\'eroux \\e\\ 2003b). But it has been suggested that at least some of the \\hi\\ lies in systems with \\hi\\ column density below that required by the traditional DLA definition, in the ``sub-Damped Lyman-$\\alpha$ Systems (sub-DLAs)'' with $19.0$ $<$ \\loghi\\ $<$ 20.3. The DLAs and sub-DLAs offer direct probes of element abundances over $ > 90 \\%$ of the age of the Universe. Zn is a good probe of the total (gas and solid phase) metallicity, in DLAs because Zn tracks Fe in most Galactic stars with [Fe/H]$> -$3, it is undepleted on interstellar dust grains, and the lines of the dominant ionisation species Zn II are often unsaturated (e.g., Pettini et al. 1999). Abundances of depleted elements such as Cr or Fe relative to Zn probe the dust content and the relative abundances can also yield information about the nucleosynthetic processes (e.g., Pettini et al. 1997; Kulkarni, Fall, \\& Truran, 1997; P\\'eroux et al., 2002; Khare et al. 2004). A study of the cosmological evolution of the \\hi\\ column density-weighted mean metallicity in DLAs (e.g., Kulkarni \\& Fall, 2002) shows surprising results. Contrary to most models of cosmic chemical evolution (e.g., Malaney \\& Chaboyer, 1996; Pei, Fall \\& Hauser, 1999), recent observations indicate at most a mild evolution in DLA global metallicity with redshift for 0$\\la$z$\\la$ 4 (Prochaska \\e\\ 2003; Khare et~al. 2004; Kulkarni, et~al., 2005; and references therein). Even theoretical models such as Smoothed-particle-hydrodynamics simulations (Nagamine, Springel, \\& Hernquist, 2004) predict that the true DLA metallicities could be 1/3 solar at z=2.5 and higher at lower redshifts.  Even at z=2.5, making a census of the predicted and observed neutral comoving densities of gas, $\\Omega$, and metals, $\\Omega_{\\rm Z}$, one finds that most of the baryons are in the Lyman-$\\alpha$ forest but its metal content is extremely low. The measured value of $\\Omega_{\\rm HI}$(DLA) is only a small fraction of $\\Omega_{\\rm baryons}$ and the DLA global mean metallicity is about 1/10$^{\\rm th}$ solar. The metallicity of Lyman break galaxies is still poorly constrained; but, in any case, these objects are known to be star-forming galaxies and may not be representative of the normal galaxy population. In total, these three components account for no more than $\\approx 10-15$\\% of what we expect to have been produced by z=2.5 (Pettini \\e\\ 2003; Bouch\\'e \\e\\ 2005). The missing metals problem in low-redshift DLAs is even more surprising since the high global star formation rate estimates at z$>$1.5 (e.g. Madau \\e\\ 1998) imply that higher metallicities should be expected at low redshift.   It is possible that $\\Omega_Z$(DLA) has been underestimated. Metal-rich DLAs could obscure quasars due to their possible high dust content (Fall \\& Pei 1993). This may be the reason for the apparent low metallicity in the DLAs observed in optically selected quasars (e.g., Fall \\& Pei 1993; Boiss\\'e et al. 1998). Recently Vladilo and P\\'eroux (2005) have shown that the fraction of high-redshift DLAs missed due to dust obscuration, could be up to 50\\%, which is consistent with the results of surveys of radio selected quasars (Ellison \\e\\ 2001). They have estimated that at z$\\sim$2.3, the real mean metallicity of DLAs could be 5 to 6 times higher than what is observed, which may help alleviate the missing metals problems. Indeed, systems at lower redshift may have significantly more dust at any given metallicity simply because the dust in these objects has had more time to process the elements.   On the other hand, new lines of evidence are pointing toward lower \\nhi\\ quasar absorbers like Lyman Limit Systems (LLS) and sub-DLAs being more metal-rich (P\\'eroux \\e\\ 2003a; Jenkins \\e\\ 2005). The dust bias, if real, is also likely to be less severe for metal-rich sub-DLAs as compared to the metal-rich DLAs due to the lower gas and therefore dust content in the former, for a constant dust to gas ratio. Thus, the obscuration bias will affect the DLAs at a lower dust-to-gas ratio as compared to the sub-DLAs. This scenario is consistent with the recent radio surveys (e.g., Vladilo \\& P\\'eroux 2005), but still needs to be further quantified observationally. Indeed, Zn measurements exist for only two sub-DLAs at low $z$: the marginally super-solar sub-DLA toward Q0058$+$019 with $z_{\\rm abs}$=0.61, \\loghi=20.08, and [Zn/H]=$+$0.08 (Pettini \\e\\ 2000), and the supersolar sub-DLA toward SDSS J1323$-$0021 with $z_{\\rm abs}$=0.72, \\loghi=20.21, and [Zn/H]=$+$0.40 (Khare et al. 2004). Although Khare et al. reported [Zn/H]=$+$0.40 for this latter absorber, the modest resolution of their MMT data could not resolve the Mg I+Zn II $\\lambda$ 2026 and Cr II+Zn II $\\lambda$ 2062 blends.The extent of line saturation on the derived column densities was also unclear.  With the goal of addressing these issues with high-resolution data, we obtained VLT/UVES spectra of this quasar, which are presented here. These new data are essential to confidently determine the metallicity by minimising the problem of line saturation.  Section 2 presents the observational set-up and data reduction process, while section 3 presents the analysis. Section 4 provides a discussion of the results. ",
        "conclusions": "Table~\\ref{t:ab} lists the abundances, using $\\loghi=20.21^{+0.21}_{-0.18}$ (Khare \\e\\ 2004) that we derived from Voigt profile fitting of the damped Ly-$\\alpha$ line in the publicly available HST/STIS spectrum of SDSS J1323$-$0021 (program GO 9382; PI: Rao). Rao et al. (2005) obtained $\\loghi=20.54^{+0.16}_{-0.15}$ from the same data set. We use the former value since it gives a smaller residual with respect to the data and therefore regard this absorber as a sub-DLA. For either \\nhi, the strength of the Zn II lines detected in our UVES spectrum implies a super-solar metallicity. Using the standard definition: $[X/H] = \\log [N(X)/N(H)]_{DLA}- \\log [N(X)/N(H)]_{\\odot}$, we find [Zn/H]=$+$0.61.  In principle, if Mg I $\\lambda$ 2852 were substantially saturated, the contribution of Mg I $\\lambda$ 2026 could be higher than our best-fit estimate. However, based on our apparent optical depth measurements and profile-fitting results, we estimate that it would take $\\sim$8 times more total Mg~{\\sc i} than the value derived from the $\\lambda 2852$ line to contribute the entirely of the $\\lambda 2026$ line. Such a high value of Mg~{\\sc i} can be ruled out by the observed profile of the $\\lambda 2852$ line. (Of course, such a scenario would also be inconsistent with the observed strength and profile of the $\\lambda 2062$ line.) To understand this issue in more detail, we estimated the maximum Mg~{\\sc i} in the dominant components that would still give the shape of the Mg I $\\lambda$2852 profile consistent with the observed profile within the noise level in the continuum. This maximum total Mg~{\\sc i} $= 2.7e13$ is about $50 \\%$ larger than the best-fit value listed in Table 2. Putting this maximum Mg I model in the $\\lambda 2026$ line, the corresponding total Zn~{\\sc ii} needed to fit the remaining part of the $\\lambda$ 2026 line would be 2.5e13, lower by $< 10 \\%$ from our best-fit value. Thus [Zn/H] is at least $>+$0.59, indicating that our result would not be affected much by saturation of Mg I $\\lambda$ 2852. Finally, as an additional check, we also estimated the maximum contribution of Cr II 2062 by rebinning our spectrum by factors of 10 or 20, measuring the upper limit for Cr II$\\lambda$ 2056 in the rebinned spectrum. We then spread this upper limit for N(Cr II) over the 2062 profile, using the velocity model derived from the combination of the lines. We assume the same Fe/Cr ratio in all components, taking the percentage of Cr II in each components relative to total N(CrII) summed over all components to be the same as the corresponding percentage of Fe II in that component. Fitting the remaining part of the $\\lambda 2062$ line with a 15-component model of Zn II, we estimated $N_{\\rm Zn II} > 2.05e13$, i.e. $[Zn/H] > 0.50$.  The abundances of Fe, Cr, Mn, and Ti lie in the range of $-$0.4 to $-$0.6 dex, and indicate that this absorber is not only metal-rich, but also very dusty. Using the model from Vladilo (2004), we find that 95\\% of the Fe is in dust phase and the total metallicity is even slightly higher than 0.6 dex. This sightline also shows substantial reddening compared to the SDSS quasar composite ($\\Delta (g-i) = 0.47$; Khare et al. 2004). Considering only the better-determined components between $-$13 and 100 km s$^{-1}$, Fe/Zn varies by a factor of $\\sim3$ and Ti/Zn varies by $\\sim$35. [Mn/Fe] varies by $\\sim 3$, but indicates a super-solar Mn abundance with respect to Fe in all components. The relative abundance of Mn with respect to Fe is not expected to exceed the solar value as Mn, unlike Fe, has an odd atomic number. However, [Mn/Fe]$>$0 is often seen in the Galactic interstellar gas due to the stronger depletion of Fe on to dust grains.  To summarise, our high resolution VLT/UVES data have allowed us to alleviate the saturation issue in Zn II and Cr II lines and therefore unambiguously prove the super-solar metallicity of the sub-DLA at $z_{\\rm abs}$=0.716 toward SDSS J1323$-$0021. If the dust obscuration bias for DLAs is indeed significant, as proposed by Vladilo and P\\'eroux (2005), sub-DLA systems such as the one reported here could be better probes of dusty regions with significant past star formation (e.g. Lauroesch \\e\\ 1996; York \\e\\ 2006), as similar DLA systems will be missed due to dust obscuration.  On the other hand, it is also, possible to envisage a scenario where the dust obscuration is not very significant in quasar absorbers (as is indicated by the rising spectrum of gamma-ray burst afterglows). In this scenario it is possible that the sub-DLA systems may indeed be more metal-rich as compared to DLAs as indicated by observations of P\\'eroux \\e\\ (2003a) and by the observations of super-solar metallicity in such systems presented here and in Pettini et al (2003). With the large-scale spectroscopic surveys of quasars currently underway (e.g. the SDSS, York \\e\\ 2000, York \\e\\ 2001), such metal-rich sub-DLA systems may be found in large numbers. If future observations indeed find such systems, sub-DLAs may contribute significantly to the overall global metallicity."
    },
    "0601/astro-ph0601553_arXiv.txt": {
        "abstract": "The correlation between the cosmological rest--frame \\nufnu spectrum peak energy, \\epi , and the isotropic equivalent radiated energy, \\eiso, discovered by Amati et al. in 2002 and confirmed/extended by subsequent osbervations, is one of the most intriguing and debated observational evidences in Gamma--Ray Bursts (GRB) astrophysics. In this paper I provide an update and a re--analysis of the \\epeiso correlation basing on an updated sample consisting of 41 long GRBs/XRFs with firm estimates of $z$ and observed peak energy, \\epo , 12 GRBs with uncertain valeus of $z$ and/or \\epo, 2 short GRBs with firm estimates of $z$ and \\epo and the peculiar sub--energetic events GRB980425/SN1998bw and GRB031203/SN2003lw. In addition to standard correlation analysis and power--law fitting, the data analysis here reported includes a modelization which accounts for sample variance. All 53 classical long GRBs and XRFs, including 11 \\swift events with published spectral parameters and fluences, have \\epi and \\eiso values, or upper/lower limits, consistent with the correlation, which shows a chance probability as low as $\\sim$7$\\times$10$^{-15}$, a slope of $\\sim$0.57 ($\\sim$0.5 when fitting by accounting for sample variance) and an extra--Poissonian logarithmic dispersion of $\\sim$0.15, it extends over $\\sim$5 orders of magnitude in \\eiso and $\\sim$3 orders of magnitude in \\epi and holds from the closer to the higher $z$ GRBs. Sub--energetic GRBs (980425 and possibly 031203) and short GRBs are found to be inconsistent with the \\epeiso correlation, showing that it can be a powerful tool for discriminating different classes of GRBs and understanding their nature and differences. I also discuss the main implications of the updated \\epeiso correlation for the models of the physics and geometry of GRB emission, its use as a pseudo--redshift estimator and the tests of possible selection effects with GRBs of unknown redshift. ",
        "introduction": "Since 1997, with the first discoveries of optical counterparts and host galaxies, redshift estimates for Gamma--Ray Bursts (GRBs) have become available, allowing the study of the intrinsic properties of this challenging astrophysical phenomena. Among these, the correlation between the photon energy (commonly called {\\it peak energy}) at which the cosmological rest frame \\nufnu spectrum peaks, \\epi, and the total isotropic--equivalent radiated energy, \\eiso, is one of the most intriguing and discussed. This correlation was discovered by \\cite{Amati02} based on BeppoSAX data and subsequently confirmed and extended to X--Ray Rich GRBs (XRRs) and X--Ray Flashes (XRFs) based on HETE--2 data \\citep{Amati03,Lamb04,Sakamoto04,Sakamoto05a}. It can be used to constrain the parameters of the various scenarios for the physics of GRB prompt emission, it is a challenging test for jet and GRB/XRF (X--Ray Flashes) unification models and it can provide hints on the nature of different sub--classes of GRBs (sub--energetic GRBs, short GRBs). Also, the \\epeiso correlation has been used for building up redshift estimators and is frequently assumed as an imput or as a required output for GRB population synthesis models. In this paper, I provide an update (up to December 2005) and a re--analysis of the \\epeiso correlation based on a sample of a total of 56 events which includes  \\swift GRBs with known redhsift and published spectral parameters and two very recent short GRBs with firm estimates of redshift and \\epi. The analysis here reported includes also fitting the data with a model which accounts for sample variance, given that this correlation is highly significant but also shows a dispersion which cannot be accounted for only by statistical fluctuations and is an important source of information. I also discuss the various explanations and implications of the existence and properties of the \\epeiso correlation, its possible use for the estimate of pseudo--redshifts, and the tests based on GRBs with unknown redshift aimed to the evaluation of the impact of possible selection effects. For this last purpouse, I also make use of published spectral parameters and fluences of a sample of 46 HETE--2 GRBs.\\\\ In order to introduce some basic information for the discussion reported in Section 5 and to allow a comparison between the results here reported and those reported in previous works, the description and properties of the updated sample (Section 3) the description and results of the data analysis results (Section 4) the discussion of the main implications and explanations of the \\epeiso correlation (Section 5) and the discussion of pseudo--redshift estimates and tests based on GRBs with unknown redshift (Section 6) are preceeded (Section 2) by a brief review of spectral and energetics properties of GRBs and of previous studies of the \\epeiso correlation.  \\begin{table*} \\begin{minipage}{155mm} \\caption{\\epi and \\eiso values for long GRBs and XRFs with firm estimates of both $z$ and \\epo (41 events), the peculiar sub--energetic event GRB980425 and the only two short GRBs with firm estimates of $z$ and \\epo , GRB050709 and GRB051221. The uncertainties are at 1$\\sigma$ significance. The \"Type\" column indicates wether the event is a normal long GRB (LONG), an X--Ray Flash (XRF) or is sub--energetic (SUB--EN). The \"Instruments\" column reports the name of the experiment(s), or of the satellite(s), that provided the estimates of spectral parameters and fluence (BAT = BATSE, SAX = \\sax, HET = HETE--2, KON = Konus, SWI = \\swift). The last two columns report the references for the spectral parameters and the references for the values and uncertainties of \\eiso, respectively. The \"--\" sign indicates that \\eiso was computed specifically for this work (see text). GRBs detected by \\swift are marked with an asterisk.}  \\begin{tabular}{llllllcc} \\hline GRB  & Type & $z$  & \\epi  & \\eiso & Instruments &  Refs. for $^{\\rm (a)}$ & Refs. for $^{\\rm (a)}$ \\\\ &   &   &  (keV)     & (10$^{52}$ erg)  &  &  spectrum &  \\eiso \\\\ \\hline 970228 & LONG  &   0.695 &  195$\\pm$64 &  1.86$\\pm$0.14 &       SAX &    (1) & (1)  \\\\ 970508 & LONG  &   0.835 &  145$\\pm$43 &  0.71$\\pm$0.15 &       SAX &  (1) & (1)  \\\\ 970828 & LONG &    0.958 &  586$\\pm$117 &  34$\\pm$4 & BAT &   (2) & (3)  \\\\ 971214 & LONG &     3.42 &  685$\\pm$133 &  24$\\pm$3 & SAX &   (1) & (1)  \\\\ 980425  & SUB--EN & 0.0085 & 55$\\pm$21 & 0.00010$\\pm$0.00002 & BAT & (4)  & (4) \\\\ 980613 & LONG &      1.096 &  194$\\pm$89 &  0.68$\\pm$0.11 &       SAX &   (1) & (1) \\\\ 980703 &  LONG &    0.966 &  503$\\pm$64 &  8.3$\\pm$0.8 &       BAT &   (2) & (3) \\\\ 990123 &  LONG &     1.60 &  1724$\\pm$466$^{\\rm (b)}$ &  266$\\pm$43 & SAX/BAT/KON &       (1,2,5) &  (1,5) \\\\ 990506 &  LONG &     1.30 &  677$\\pm$156$^{\\rm (b)}$ &  109$\\pm$11 & BAT/KON &       (2,5) & (3,5)  \\\\ 990510 &  LONG &     1.619 &  423$\\pm$42 &  20$\\pm$3 & SAX &       (1) & (1)  \\\\ 990705 &  LONG &    0.842 &  459$\\pm$139$^{\\rm (b)}$ &  21$\\pm$3 & SAX/KON &       (1) &  (1) \\\\ 990712 &  LONG &    0.434 &  93$\\pm$15 &  0.78$\\pm$0.15 &       SAX &   (1) &  (1) \\\\ 991208 &   LONG &   0.706 &  313$\\pm$31 &  25.9$\\pm$2.1 &       KON &   (5) &  (5)  \\\\ 991216 &   LONG &    1.02 &  648$\\pm$134$^{\\rm (b)}$ &  78$\\pm$8 & BAT/KON &     (2) &  (3) \\\\ 000131  &   LONG &    4.50 &  987$\\pm$ 416$^{\\rm (b)}$ &  199$\\pm$35 & BAT/KON &    (5) & (5) \\\\ 000210 &  LONG &    0.846 &  753$\\pm$26 &  17.3$\\pm$1.9 & KON &    (5) & (5)  \\\\ 000418 &  LONG &     1.12 &  284$\\pm$21 &  10.6$\\pm$2.0 & KON &     (5) & (5)  \\\\ 000911 &   LONG &    1.06 &  1856$\\pm$371$^{\\rm (c)}$ &  78$\\pm$16 & KON &    (6)  & (6) \\\\ 000926 &   LONG &    2.07 &  310$\\pm$20 &  31.4$\\pm$6.8 & KON &     (5)  & (5) \\\\ 010222 &  LONG &     1.48 &  766$\\pm$30 &  94$\\pm$10 & KON &     (5) & (5)  \\\\ 010921 &   LONG &   0.450 &  129$\\pm$26 &  1.10$\\pm$0.11 &       HET &  (8) & (3) \\\\ 011211 &   LONG &    2.14 &  186$\\pm$24 &  6.3$\\pm$0.7 &       SAX &   (3) & (3)  \\\\ 020124 &  LONG &     3.20 &  448$\\pm$148$^{\\rm (b)}$ &  31$\\pm$3 & HET/KON &    (8,5) & (3)  \\\\ 020813 &  LONG &     1.25 &  590$\\pm$151$^{\\rm (b)}$ &  76$\\pm$19$^{\\rm (b)}$  & HET/KON &   (5,7) & (3,5)  \\\\ 020819b &  LONG &    0.410 &  70$\\pm$21 &  0.79$\\pm$0.20 &       HET &   (8) & -- \\\\ 020903 &  XRF &    0.250 &  3.37$\\pm$1.79 & 0.0028$\\pm$0.0007 &       HET &  (9) & (9)  \\\\ 021004 &   LONG &    2.30 &  266$\\pm$117 &  3.8$\\pm$0.5 &       HET &    (8) & (26) \\\\ 021211  &   LONG &   1.01 &  127$\\pm$52$^{\\rm (b)}$ & 1.3$\\pm$0.15 & HET/KON    &    (5,10) & (5,13)  \\\\ 030226 &   LONG &    1.98 &  289$\\pm$66 &  14$\\pm$1.5 & HET &  (8) &  (13) \\\\ 030328 &  LONG &     1.52 &  328$\\pm$55 &  43$\\pm$4 & HET/KON &  (5,8) & -- \\\\ 030329 &  LONG &     0.17 &  100$\\pm$23$^{\\rm (b)}$ & 1.7$\\pm$0.3 $^{\\rm (b)}$ & HET/KON &   (5,8) & -- \\\\ 030429 &  LONG &     2.65 &  128$\\pm$26 &  2.50$\\pm$0.30 &       HET &  (8) & (26) \\\\ 040924 &   LONG &   0.859 &  102$\\pm$35$^{\\rm (b)}$ &  1.1$\\pm$0.12 &       HET/KON &   (14,28) & (27)  \\\\ 041006 &  LONG &    0.716 &  98$\\pm$20 &  3.5$\\pm$1.0 &       HET &   (29) & --  \\\\ 050318* &  LONG &     1.44 &  115$\\pm$25 &  2.55$\\pm$0.18 &       SWI &   (16) & (16) \\\\ 050401* &  LONG &     2.90 &  467$\\pm$110 &  41$\\pm$8 & KON &  (17) &   (17)   \\\\ 050416a* & XRF &     0.650 &  25.1$\\pm$4.2 &  0.12$\\pm$0.02 &       SWI &  (19) & (19)  \\\\ 050525* &   LONG &   0.606 &  127$\\pm$10 &  3.39$\\pm$0.17 &       SWI & (20) &   --     \\\\ 050603* &   LONG &    2.821 &  1333$\\pm$107 &  70$\\pm$5 & KON & (21) &  --  \\\\ 050709  & SHORT & 0.16 & 100$\\pm$16 & 0.0103$\\pm$0.0021 & HET & (18) &  (18) \\\\ 050922c* &  LONG &     2.198 &  415$\\pm$111 &  6.1$\\pm$2.0 & HET &  (22) &  (22)  \\\\ 051022* & LONG &  0.80 &  754$\\pm$258$^{\\rm (b)}$ &  63$\\pm$6 & HET/KON &  (23,30) &  --  \\\\ 051109* & LONG &      2.346 &  539$\\pm$200 &  7.5$\\pm$0.8 &       KON &  (24) &  --   \\\\ 051221* &  SHORT &   0.5465 & 622$\\pm$35  &  0.29$\\pm$0.06 &       KON &  (25) &  (25) \\\\ \\hline \\end{tabular} \\begin{list}{}{} \\item[]Notes. $^{\\rm (a)}$References for the spectral parameters and for the values and uncertainties of \\epi and \\eiso: (1) \\cite{Amati02}, (2) \\cite{Jimenez01}, (3) \\cite{Amati03}, (4) \\cite{Yamazaki03b}, (5) \\cite{Ulanov05}, (6) \\cite{Price02}, (7) \\cite{Barraud03}, (8) \\cite{Sakamoto05b}, (9) \\cite{Sakamoto04}, (10) \\cite{Crew03}, (11) \\cite{Atteia03}, (12) \\cite{Atteia05}, (13) \\cite{Ghirlanda04a}, (14) \\cite{Golenetskii04}, (15) \\cite{Friedman05}, (16) \\cite{Perri05}, (17) \\cite{Golenetskii05a}, (18) \\cite{Villasenor05}, (19) \\cite{Sakamoto05a}, (20) \\cite{Cummings05}, (21) \\cite{Golenetskii05b}, (22) \\cite{Crew05a}, (23) \\cite{Golenetskii05c}, (24) \\cite{Golenetskii05d}, (25) \\cite{Golenetskii05e}, (26) \\cite{Friedman05}, (27) \\cite{Ghirlanda05b}, (28) \\cite{Fenimore04}, (29) Official HETE GRB page (http://space.mit.edu/HETE/Bursts/), (30) \\cite{Doty05} $^{\\rm (b)}$ Uncertainties enlarged in order to account for significantly different values measured by different instruments (see text). $^{\\rm (c)}$ Value derived from time resolved spectral analysis by weighting each time interval with the ratio between its fluence and the total GRB fluence. A 20\\% error is conservatively assumed. \\end{list} \\end{minipage} \\end{table*} ",
        "conclusions": "The analysis presented in previous Sections, based on an updated sample containing about twice events with respect to previous works and including the recent \\swift GRBs, confirms and strenghtens the \\epeiso correlation for long GRBs/XRFs and gives its best characterization up to now in terms of index and normalization of the best--fit power--law and of its dispersion. Remarkably, the correlation now extends over $\\sim$5 orders of magnitude in \\eiso,  $\\sim$3 orders of magnitude in \\epi and over a redshift range $\\sim$0.15$<$$z$$<$4.5 (but also the highest $z$ event, GRB050904 at $z$=6.29, has \\epi and \\eiso values consistent with it). Since its discovery in 2002 \\citep{Amati02} and in particular its confirmation and extension to XRFs \\citep{Amati03, Lamb04, Sakamoto04}, the origin of the \\epeiso correlation and its implications for GRB models have been investigated by several works. The impact of this robust observational evidence on prompt emission models concerns mainly the physics, the geometry (i.e. shape and properties of jets), viewing angle effects and GRB/XRF unification. Indeed, the existence of the \\epeiso correlation and its properties are also often used as an ingredient or a test output for GRB synthesis models, as in the case of, e.g., the GRB/XRF model by \\cite{Barraud05}, the multi--subjets model by \\cite{Toma05}, the uniform jet model by \\cite{Lamb05}, the study of the impact of off--jet relativistic kinematics by \\cite{Donaghy05a}, the dissipative photosphere models by \\cite{Rees05}, the supercritical pile model by \\cite{Mastichiadis06}. Another important outcome of the analysis presented in this paper is the clear evidence that, in addition to the peculiar sub--energetic GRB980425 (and possibly GRB031203), short GRBs do not follow the \\epeiso correlation, as suggested by the different location between long and short BATSE GRBs in the hardness--intensity plane. Below I summarize these topics and discuss also the possible origin of the extra--Poissonian dispersion of the correlation.  \\subsection{Physics of prompt emission}  The physics of the prompt emission of GRBs is still far to be settled and a variety of scenarios, within the standard fireball picture, have been proposed, based on different emission mechanisms (e.g. SSM internal shocks, Inverse Compton dominated internal shocks, SSM external shocks, photospheric emission dominated models) and different kinds of fireball (e.g. kinetic energy dominated or Poynting flux dominated), see e.g. \\cite{Zhang02} for a review. In general, both \\epi and \\eiso are linked to the fireball bulk Lorentz factor, $\\Gamma$, in a way that varies in each scenario, and the existence and properties of the \\epeiso correlation allow to constrain the range of values of the parameters, see, e.g., \\cite{Zhang02} and \\cite{Schaefer03}. For instance, as shown, e.g., by \\cite{Zhang02}, \\cite{Ryde05} and \\cite{Rees05}, for a power--law electron distribution generated in an internal shock within a fireball with bulk Lorentz factor $\\Gamma$, it is possible to derive the relation \\epi $\\propto$ $\\Gamma$$^{-2} L^{1/2} t_\\nu^{-1}$ , where $L$ is the GRB luminosity and $t_\\nu$ is the typical variability time scale. Clearly, in order to produce the observed \\epeiso correlation the above formula would require that $\\Gamma$ and $t_\\nu$ are approximately the same for all GRBs, an assumption which is difficult to justify. Things get even more complicated if one takes into account that the models generally assume $L$ $\\propto$ $\\Gamma^{\\beta}$, with the value of $\\beta$ varying in each scenario and is typically $\\sim$2--3 \\citep{Zhang02,Schaefer03,Ramirez05a,Ryde05}. More specific examples of the constraints put by the \\epeiso correlation on the parameters of SSM and IC based emission models, both in internal and external shocks, can be found, e.g.,in \\cite{Zhang02,Schaefer03}. An interesting possibility, which is currently the subject of many theoretical works, is that a substantial contribution to prompt radiation of GRBs comes from direct or Comptonized thermal emission from the photosphere of the fireball \\citep{Zhang02,Ryde05,Rees05,Ramirez05a}. This could explain the very hard spectra observed at the beginning of several events \\citep{Preece00,Frontera00b,Ghirlanda03}, inconsistent with SSM models, and the smooth curvature characterizing GRBs average spectra. In this scenarios, \\epi is mainly determined by the peak temperature $T_{pk}$ of black--body distributed photons and thus naturally linked to the luminosity or radiated energy. For instance, for Comptonized emission from the photosphere one can derive the relations \\epi $\\propto$ $\\Gamma$T$_{pk}$ $\\propto$ $\\Gamma$$^{2} L^{-1/4}$ or \\epi $\\propto$ $\\Gamma$T$_{pk}$ $\\propto$ r$_0$$^{-1/2} L^{1/4}$ (where $r_0$ is a particular distance between the central engine and the emitting region), depending on the assumptions made \\citep{Rees05}. Also in this case, off course, the $L$ $\\propto$ $\\Gamma^{\\beta}$ relation plays a decisive role. As shown by \\cite{Rees05}, in this scenario the correct \\epeiso relation can be obtained for some specific physical conditions just below the photosphere.  Finally, also, the fact that the \\epi distribution is broader than inferred previously basing mainly on the observed \\epo values of bright BATSE GRBs, as shown in Figure 1, can put important constraints on models for the physics of GRB prompt emission \\citep{Zhang02}.   \\begin{figure} \\includegraphics[width=9cm]{disp} \\caption{Logarithmic dispersion of the \\epi values of the 41 GRBs with firm redshift and \\epi estimates around the power--law best fitting the \\epeiso correlation without accounting for sample variance. I also show the best fit Gaussian, which has a dispersion of $\\sigma_{log(E_p,i)}$ $\\sim$ 0.21 .} \\end{figure}  \\subsection{Jets, viewing angle effects and GRB/XRF unification models}  The validity of the \\epeiso correlation from the most energetic GRBs to XRFs (see Figure 2) confirms that these two phenomena have the same origin and is a very challenging observable for GRB jet models. Indeed, these models have to explain not only how \\eiso and \\epi are linked to the jet opening angle, $\\theta$$_{jet}$, and/or to the viewing angle with respect to the jet axis, $\\theta$$_v$, but also how \\eiso can span over several orders of magnitudes. In the most simple scenario, the uniform jet model \\citep{Frail01,Lamb05}, jet opening angles are variable and the observer measures the same value of \\eiso independently of $\\theta$$_v$. In the other popular scenario, the universal structured jet model (e.g. Rossi, Lazzati \\& Rees \\nocite{Rossi02}), \\eiso depends on $\\theta$$_v$. As discussed in Section 2, in the hypothesis that achromatic breaks found in the afterglow light curves of some GRBs with known redshift are due to collimated emission, it was originally found \\citep{Frail01,Berger03} that the collimation corrected radiated energy, \\ega, is of the same order ($\\sim$10$^{51}$ erg) for most GRBs and that \\eiso $\\propto$ $\\theta$$_{jet}^{-2}$, assuming a uniform jet. In the case of structured jet models, which assume that $\\theta$$_{jet}$ is similar for all GRBs (hence this scenario is also called universal jet model) the same observations imply that \\eiso $\\propto$ $\\theta$$_{v}^{-2}$. Thus, always under the assumption of a nearly constant \\ega, the found \\epeiso correlation implies \\epi $\\propto$ $\\theta$$_{jet}^{-1}$ and \\epi $\\propto$ $\\theta$$_{v}^{-1}$ for the uniform and structured jet models, respectively. \\cite{Lamb05} argue that the structured universal jet model, in order to explain the validity of the \\epeiso correlation from XRFs to energetic GRBs, predicts a number of detected XRFs several orders of magnitude higher than the observed one ($\\sim$1/3 than that of GRBs). In their view, the uniform jet model can overcome these problems by assuming a distribution of jet opening angles N($\\theta$$_{jet}$) $\\propto$ $\\theta$$_{jet}^{-2}$. This implies that the great majority of GRBs have opening angles smaller than $\\sim$1$^{\\circ}$ and that the true rate of GRBs is several orders of magnitude higher than observed and comparable to that of SN Ic. On the other hand, \\cite{Zhang04} show that the requirement that most GRBs have jet opening angles less than 1 degree, needed in the uniform jet scenario in order to explain the \\epeiso correlation, as discussed above, implies values of the fireball kinetic energy and/or of the interstellar medium density much higher than those inferred from the afterglow decay light curves. Together with other authors, e.g. \\cite{Lloyd04,Dai05}, they propose a modification of the universal structured jet model, the quasi--universal Gaussian structured jet.  In this model, the measured \\eiso undergoes a mild variation for values of $\\theta$$_{v}$ inside a typical angle, which has a quasi--universal value for all GRBs/XRFs, whereas it decreases very rapidly (e.g. exponentially) for values outside the typical angle. In this way, the universal structured jet scenario can reproduce the \\epeiso correlation and predict the observed ratio between the number of XRFs and that of GRBs. Recently, a Fisher--shape has been proposed, for both the variable angle and universal angle scenarios \\citep{Donaghy05b,Dai05}, as a very promising alternative, in particular for the explanation of the validity of the \\epeiso correlation from the brightest GRBs to XRFs. Other jet models proposed very recently that can reproduce the \\epeiso correlation include the ring--shaped jet model, see \\cite{Eichler04}, and the multi--component (subjets) model, see \\cite{Toma05}.  Of particular interest are the off--axis scenarios, in which the jet is typically assumed to be uniform but, due to relativistic beaming and Doppler effects, for $\\theta$$_{v}$ $>$ $\\theta$$_{jet}$ the measured emissivity does not sharply go to zero and the event is detected by the observer with \\eiso and \\epi dropping rapidly as $\\theta$$_{v}$ increases \\citep{Yamazaki03a,Granot02,Eichler04,Donaghy05a}. In these models, XRFs are those events seen very off--axis and the XRFs rate with respect to GRBs and the \\epeiso correlation can be correctly predicted. As shown e.g. in \\cite{Yamazaki04} for a simple model of GRB jet, if the Doppler shift factor is $\\delta = [\\Gamma(1 - \\beta cos(\\theta_{v} - \\theta_{jet})]^{-1}$ (where $\\beta$ is the velocity of the outflow in units of speed of light), \\epi and \\eiso scale, with respect to their values observable at the edge of the jet, as \\epi $\\propto$ $\\delta$ and \\eiso $\\propto$ $\\delta^{1-\\alpha}$ , where $\\alpha$ is the spectral index of the prompt emission photon spectrum in the hard X--ray energy band. By combining these relations one can obtain the \\epeiso correlation with index 0.5 for classical GRBs ($\\alpha$$\\sim$$-$1) and 0.3 for XRFs ($\\alpha$$\\sim$$-$2). A detailed study of off--jet relativistic kinematics effects has been recently performed by \\cite{Donaghy05a} for a set uniform (i.e. top hat shaped -- variable opening angle) jet models, finding that these scenarios predict a significant population of bursts away from the \\epeiso correlation, unless $\\Gamma$ $>$ 300 for all bursts or there is a strong anti--correlation between $\\Gamma$ and the jet solid angle. Finally, the off--axis effects for very weak and soft events can be applied in a similar way as described above in the context of the cannon ball (CB) model for GRBs in order to reproduce the \\epeiso correlation \\citep{Dar04}.  \\subsection{The dispersion of the correlation}  In addition to its existence and slope, also the extra--Poissonian dispersion of the \\epeiso correlation is an important source of information. As discussed in Section 3 and shown in Table 3, while the correlation is very highly significant, the scatter of the data around the best fit power--law exceed that expected by statistical fluctuations alone and produces high values of \\cqr. By fitting with a Gaussian the dispersion of the central values of log(\\epi) around the best fit model, I obtain $\\sigma$$_{logE_{p,i}}$$\\sim$0.2, while by fitting the whole data with the method by \\cite{Dagostini05}, which includes sample variance directly in the model, I obtain $\\sigma$$_{logE_{p,i}}$ = 0.15$_{-0.4}^{+0.4}$ . A similar scatter, even if computed with only the first of the two methods reported above, is found for the \\epi -- $L_{\\rm iso}$ and \\epi -- $L_{\\rm peak,iso}$ correlations, see, e.g., \\cite{Ghirlanda05a}. Intriguingly, the \\epega correlation shows instead a lower dispersion, of the order of $\\sigma$$_{logE_{p,i}}$$\\sim$0.1 \\citep{Ghirlanda04a,Nava06}, even if this correlation is based on a still low number of events, it requires an estimate of $t_b$ in addition to \\epo and $z$, it depends on jet model and circum--burst environment properties (density, distribution) and there are possible outliers, as discussed e.g. by \\cite{Friedman05} and may be indicated by the lack of break in the X--ray afterglow light curve of some \\swift GRBs with known redshift. A low 3--D dispersion characterizes also the \\epeisotb correlation, which is a kind of model--independent version of the \\epega correlation \\citep{Liang05,Nava06}. The existence of both the \\epeiso and \\epega correlations is due to the fact that the collimation angles of GRBs are distributed over a relatively narrow range of values; the lower dispersion of the \\epega correlation indicates that at least part of the scatter of the \\epeiso correlation is due to the dispersion of jet opening angles. And indeed, the comparison of the properties of the two correlations has been used, in addition to the study of the relation between jet opening angle and radiated energy, to infer the distribution of jet opening angles, as done, e.g., by \\cite{Ghirlanda05c,Bosnjak06a,Donaghy05a}. Very recently, it has been found evidence that a relevant contribution to the dispersion of the correlation is due to temporal properties of the prompt emission, like the half--width of the auto--correlation function \\citep{Borgonovo06} and the \"high--signal\" time scale \\citep{Firmani06}. Other contribution to the scatter of the \\epeiso correlation may come from viewing angle effects, e.g. \\cite{Levinson05}, the dispersion of the parameters of the fireball and/or of the time scales (as discussed in Section 5.1 concerning synchrotron emission in internal shocks), the inhomogeneous structure of the jet, e.g. \\cite{Toma05}), the possible presence of significant amount of material in the circum--burst region, that would affect the estimates of both \\eiso and \\epi with a global qualitative effect of steepening the correlation and increasing its dispersion \\citep{Longo05}. In general, several GRB population synthesis models, like those mentioned at the beginning of this Section, predict a scatter of the \\epeiso correlation which depend on the parameters values.  When investigating the above physical interpretations and implications, it is important to take into account that at least part of the dispersion of the \\epeiso correlation could arise from instrumental and other systematic effects in the estimates of \\epi and \\eiso and their uncertainties. As discussed e.g. by \\cite{Lloyd00,Lloyd02}, data truncation effects, i.e. the systematics introduced by the limited energy band of the detector, may affect significantly the estimate of \\epo. Indeed, as discussed in Section 2, typical GRB spectra are characterized by a very smooth curvature and cover $\\sim$3 orders of magnitude or more in photon energy. Thus, unless the energy band extends from few keV to few MeV, only a portion of the spectrum can be detected by a single instrument, which may cause a bias in the estimate of the spectral parameters, especially when \\epo is not far from the low or high thresholds. Also, both \\sax ($\\sim$2--700 keV) and HETE--2 (2--400 keV), in addition to be capable to detect X--ray rich events and XRFs which could not be triggered by BATSE ($\\sim$25--2000 keV), showed that X--ray emission below $\\sim$30 keV of normal GRBs can last up to several tens of seconds more than in the hard X--ray energy band. Thus, a GRB detector working at energies higher than few tens of keV, like BATSE, may have lost, for a fraction of GRBs, a substantial portion of the soft X--ray emission, with a consequent overestimate of $\\alpha$ and \\epo . These effects can indeed be seen when comparing the X-- and hard X--rays duration and light curves of \\sax \\citep{Frontera00b,Amati02} and HETE--2 \\citep{Sakamoto05b} GRBs and the average spectral parameters estimated by \\sax/WFC+GRBM and BATSE \\citep{Amati02,Jimenez01} for those events revealed by both satellites. This is true, even if to a minor extent, when comparing the best fit spectral models obtained with HETE--2/FREGATE (7--400 keV) alone \\citep{Barraud03} with those obtained by jointly fitting HETE--2/WXM (2--30 keV) and HETE--2/FREGATE data \\citep{Sakamoto05b}.  Concerning \\eiso, the main source of possible systematics is the extrapolation to the 1--10000 keV cosmological rest--frame energy band of the spectral model obtained by fitting data in the instrument energy band (see Section 2). Indeed, the (known) statistical uncertainties and (unknown) biases in the estimates of spectral parameters may affect significantly the estimate of \\eiso . This is particularly true for estimates based on spectra from instruments with high energy bound at a few hundreds of keV, like HETE--2 or \\swift/BAT, which in several cases cannot provide a reliable estimate of the high energy spectral index $\\beta$ . In addition, the typical choice of computing \\eiso in the 1--10000 keV rest--frame energy band may not be optimal for very soft events with values of \\epi below a few tens of keV, for which this method can likely lead to an underestimate of \\eiso. Off course, also the choice of the cosmological parameters for the computation of the luminosity distance, usually made by assuming values in the ranges given by the so called \"concordance cosmology\" based on type Ia SNe and CMB measurements, affects the values of \\eiso .  Finally, very recently \\swift/XRT found evidence of X--ray flares following the end of the prompt emission as detected by \\swift/BAT, e.g. \\cite{Burrows05}. One possible interpretation of these phenomena is that they are due to continued activity of the engine and/or late internal shocks \\citep{Fan05,Burrows05,Wu05,Perna06,Proga06} and thus can be considered part of the prompt emission. Given that these events are typically soft and that their fluence can be a significant fraction of that of the GRB, their non--detection (because of sensitivity) by past and current GRB detectors may also have biased the estimates of \\eiso (under--estimate) and  of \\epi (over--estimate).   \\subsection{Short GRBs}  In the past, the \\epeiso correlation has been studied basing on data of long GRBs, given that no redshift information was available for short GRBs. Nevertheless, thanks to the measurements performed by HETE--2 and \\swift, in the last year it has been possible to detect afterglow emission and possible optical counterparts and/or host galaxies for a few short GRBs. As discussed in Section 3, I have included in the analysis reported in this paper the two short GRBs with firm estimates of $z$ and \\epo: GRB050709 \\citep{Villasenor05,Hjorth05} and GRB051221 \\citep{Golenetskii05e,Berger05}. As can be clearly seen in Figure 3, these events are outliers to the \\epeiso correlation, even when taking into account its extra--Poissonian dispersion. In addition, the spectral data of other recently localized short GRBs with possible redshifts, GRB050509b \\citep{Gehrels05,Pedersen05}, GRB050724 \\citep{Krimm05,Berger05a} and GRB050813 \\citep{Sato05b,Berger05b}, indicate a likely inconsistence with the \\epeiso correlation, even though an estimate of \\epo was not possible for these events. These evidences confirm the expectations based on the fact that short GRBs tend to form a separate class with respect to long GRBs in the hardness--intensity plane, i.e. they tend to be weaker and harder, and clearly shows the potentiality of the use of the \\epeiso plane for distinguishing different classes of GRBs and understanding their nature. In particular, given that short and long GRBs partially overlap in the hardness--intensity and hardness--duration planes, that a fraction of short GRBs show a softer extended emission which can last tens of seconds, e.g. \\cite{Norris06}, the possible relevant impact of spectral/temporal trigger selection effects, it is sometimes difficult to establish to which class a GRB belongs. Thus the \\epeiso plane may be a powerful tool under this respect, especially when combined e.g with the lack of spectral lag and spectral evolution which seems to characterize short GRBs \\citep{Norris06}.  While the commonly accepted hypothesis for the progenitors of long GRBs is their association with the core collapse of massive fastly rotating stars (so called \"hypernova\" or \"collapsar\" models), based on their duration (from few seconds up to hundreds of seconds), their huge \\eiso (Figure 1), their typical location inside blue galaxies with high star formation rate and the evidence of a metal--rich circum--burst environment (as inferred from absorption / emission features in prompt and afterglow emission spectra), short GRBs are thought to originate from the coalescence of neutron star -- neutron star or neutron star -- black hole binaries (the \"merger\" scenarios). It has also been proposed that a fraction of them may be giant flare from extra--galactic soft gamma repeaters, see, e.g., \\cite{Meszaros02,Piran05} for reviews. The very recent observations mentioned above, show that, as long GRBs, also short GRBs lie at cosmological distances (0.1$<$z$<$1) but they are less energetic (see Table 1 and Figures 1 and 3). Also, their host galaxies show different morphologies and star forming activity: for instance, short GRB050724 (as possibly short GRB050509b) came from an elliptical galaxy with low star formation, whereas short GRB050709 was associated with an irregular late type star forming galaxy. Both events were located towards the outskirts of their host galaxies. All these properties match the predictions of the merger scenarios (e.g. Belczynski et al. 2006 \\nocite{Belczynski06}). Off course, these evidences concern still a very low number of short GRBs and thus it is premature to draw any definitive conclusion. Anyway, the different nature of the progenitors between short and long GRBs can help in understanding their different behavior in the \\epeiso plane. \\cite{Ghirlanda04b} found, basing on BATSE data, that the emission properties of short GRBs are similar to those of the first $\\sim$1--2 s of long GRBs. This could indicate that the central engine is the same for the two classes, but works for a longer time in long GRBs. This would explain the low radiated energy by short GRBs and, given that the emission would stop before the typical hard to soft evolution observed for long GRBs, also their high \\epi (with respect to their \\eiso). The merger scenarios naturally explain the short life of the central engine, and thus the low radiated energy and the quitting of the emission before hard to soft spectral evolution. In addition, they also predict a weak afterglow emission, as recently observed \\citep{Fox05}, because of the cleaner and lower density circum--burst medium with respect to that predicted by the hypernova scenarios for long GRBs, which would cause a very inefficient external shock. In the GRB scenarios where most of the prompt emission of long GRBs is due to the external shock too, this naturally explains the lack of long lasting and softening emission in short GRBs. Again, this would produce an high average \\epi value with respect to the radiate energy, and thus the inconsistency with the \\epeiso correlation holding for long GRBs. Finally, in order to explain the very low or 0 spectral lag observed in short GRBs with respect to long GRBs, \\cite{Norris06} consider the hypothesis that the typical $\\Gamma$ of short events is several times that of long ones, and thus of the order of $\\sim$500--1000, as predicted e.g by the merger model of \\cite{Aloy05}. With such a high Lorentz factor, internal shocks are expected to have a low efficiency, and indeed this is one possible explanation for the weakness and softness of XRFs (e.g Barraud et al. 2005 \\nocite{Barraud05}). If we assume the scenario proposed by e.g. \\cite{Ghirlanda03}, in which the spectrally hard emission observed in the first seconds of long GRBs is due to (possibly Compton dragged) thermal emission from the photosphere and the later emission to synchrotron processes occurring in internal shocks, the high $\\Gamma$ would thus produce a short--hard emission possibly followed by a weak soft component, as observed for several short GRBs \\citep{Villasenor05,Norris06}.  \\subsection{Sub--energetic GRBs and the GRB/SN connection}  As can be seen in Figure 3, in addition to the two short GRBs also the prototype event for the GRB/SN connection, GRB980425/SN1998bw, is characterized by values of \\epi and \\eiso completely inconsistent with the \\epeiso correlation holding for the other events. From an observational point of view, this is a direct consequence of the fact that the event is characterized by a fluence and a measured peak energy in the range of classical long GRBs but, based on the commonly accepted association with SN1998bw, it lies at a much lower distance (z = 0.0085). Figure 3 shows that also another event associated with a SN event, GRB031203, is characterized by a value of \\epi which, combined with its low value of \\eiso, makes it completely inconsistent with the correlation. Given that GRB031203 is the most similar event to GRB980425 under several points of view (although lying at a larger distance, z $\\sim$ 0.1), in particular the strong evidence of association with a SN event (SN2003lw) and the low afterglow energy inferred from radio observations \\citep{Soderberg04,Sazonov04}, this inconsistency has been invoked as a further evidence of the existence of a class of close sub--energetic GRBs. However, the lower limit on \\epi based on ISGRI data \\citep{Sazonov04} and the \\epi estimate based on Konus data (Table 2, Ulanov et al. 2005) are currently debated based on the dust echo observed with XMM--Newton, which could indicate a much softer prompt emission spectrum \\citep{Watson06}. The sample of GRBs with most evidence of association with a SN include also GRB030329 (SN2003dh) and GRB021211 (SN2002lt), which, in converse, are not sub--energetic and are characterized by \\epi and \\eiso values fully consistent with the \\epeiso correlation. The fact the two closest and sub--energetic among those GRBs most clearly associated with a SN are outliers to the \\epeiso correlation is intriguing. As in the case of short GRBs, these evidences show the potential use of the \\epeiso plane to distinguish among different sub--classes of GRBs, and have important implications for GRB/XRF/SN unification models. The most common explanations assume that the peculiarity of these events is due to particular and uncommon viewing angles, as proposed e.g. by \\cite{Yamazaki03b} for GRB980425 and \\cite{Ramirez05b} for GRB031203. Based on relativistic beaming and Doppler effects and the assumption of a uniform jet, they find that \\epi $\\propto$ $\\delta$ and \\eiso $\\propto$ $\\delta$$^{3}$, where $\\delta$ is the relativistic Doppler factor (see Section 5.2). For large off--axis viewing angles the different dependence of \\epi and \\eiso on $\\delta$ would cause significant deviations form the \\epeiso correlation. An alternative explanation has been suggested by \\cite{Dado05} in the framework of the CB model \\citep{Dar04}. In this scenario, the \\nufnu spectra of GRBs are characterized by two peaks, one at sub--MeV energies, the normal peak following the \\epeiso correlation, and one in the GeV--TeV range. When a GRB is seen very off--axis, the same relativistic Doppler and beaming effects discussed above would shift the high energy peak at low energies, making it to be confused with the normal low energy GRB peak. Very recently, off--axis scenarios have been seriously challanged by the detection of a very close ($z = 0.0331$) and sub--energetic event, GRB060218/SN2006aj, which, contrary to GRB980425, is characterized by \\epi and \\eiso values fully consistent with the \\epeiso correlation \\citep{Amati06,Ghisellini06}. \\\\ Finally, \\cite{Bosnjak06b}, based on BATSE data of GRBs with unknown redshift found evidence that a part of GRB with indications of association with a SN are inconsistent with the \\epeiso correlation, as also the lag--luminosity relation \\citep{Norris00}, for any value of $z$, which would confirm the existence of a peculiar class of sub--energetic, SN--associated events. See, however, next Section for a discussion of this method."
    },
    "0601/astro-ph0601415_arXiv.txt": {
        "abstract": "We investigate energetic type Ic supernovae as production sites for $^6$Li and Be in the early stages of the Milky Way. Recent observations have revealed that some very metal-poor stars with [Fe/H]$<-2.5$ possess unexpectedly high abundances of $^6$Li. Some also exbihit enhanced abundances of Be as well as N. From a theoretical point of view, recent studies of the evolution of metal-poor massive stars show that rotation-induced mixing can enrich the outer H and He layers with C, N, and O (CNO) elements, particularly N, and at the same time cause intense mass loss of these layers. Here we consider energetic supernova explosions occurring after the progeniter star has lost all but a small fraction of the He layer. The fastest portion of the supernova ejecta can interact directly with the circumstellar matter (CSM), both composed of He and CNO, and induce light element production through spallation and He-He fusion reactions. The CSM should be sufficiently thick to energetic particles so that the interactions terminate within its innermost regions. We calculate the resulting $^6$Li/O and $^9$Be/O ratios in the ejecta$+$CSM material out of which the very metal-poor stars may form. We find that they are consistent with the observed values if the mass of the He layer remaining on the pre-explosion core is $\\sim 0.01-0.1 \\Msun$, and the mass fraction of N mixed in the He layer is $\\sim 0.01$. Further observations of $^6$Li, Be and N at low metallicity should provide critical tests of this production scenario. ",
        "introduction": "Among the light elements Li, Be and B (LiBeB), $^7$Li is thought to arise from a variety of processes, including big bang nucleosynthesis \\citep{Spite82}, asymptotic giant branch stars, novae \\citep{DAntona91}, and the $\\nu$-process in type II supernovae \\citep{Woosley90}; the latter may also contribute to $^{11}$B. On the other hand, the main production channel for the rest, in particular for $^6$Li and Be, is believed to be cosmic-ray induced nuclear reactions. The most widely discussed models of LiBeB production are based on cosmic rays accelerated in supernova shocks \\citep{Meneguzzi71,Vangioni00}. Observations of metal-poor stars in the Galactic halo show a primary relation between [Fe/H] and [Be/H] or [B/H], which is consistent with spallation by cosmic rays enriched with C, N, and O (CNO) from fresh SN ejecta impinging on interstellar H or He \\citep{Duncan92}.  An alternative possibility was proposed by \\citet{Fields96, Fields02}, who considered explosions of Type Ic supernovae (SNe Ic) as a site for primary LiBeB production. Since it is expected that a fraction of C and O in the surface layers of the ejecta are accelerated to energies above the threshold for spallation reactions when the shock passes through the stellar surface, LiBeB production can occur through the direct interaction of the ejecta with the ambient material, without the need for shock acceleration of cosmic rays. This was explored in greater depth by \\citet{Nakamura04}, who used more realistic stellar models and equations of state, together with a 1-dimensional relativistic hydrodynamic code to investigate in detail how much of the ejecta mass acquires sufficient energies for the LiBeB production. All of these studies used stellar models which completely lost their H and He envelopes and assumed the target to be interstellar matter (ISM) consisting of 90\\% H and 10\\% He, ignoring any circumstellar matter (CSM). Therefore, the only reactions under consideration were ${\\rm C, O} + {\\rm H, He} \\rightarrow {\\rm LiBeB}$.  Recent observations by VLT/UVES \\citep{Asplund05} and Subaru/HDS (Inoue  et al. 2005, Aoki et al, in prep.) have revealed that some very metal-poor stars possess surprisingly high abundances of $^6$Li, much higher than expected from standard supernova cosmic ray models \\citep{Ramaty00,Suzuki01,Prantzos06}. The measured values are also higher than can be accommodated in the SN Ic production scenario discussed above, even in the case of an energetic explosion similar to SN 1998bw, as can be seen from the results of \\citet{Nakamura04}. However, besides spallation of C and O, the fusion reaction ${\\rm He + He} \\rightarrow {\\rm ^6Li}$ may also be potentially important. This reaction will be significant in the conceivable case that a small fraction of He is still remaining in the surface layer of the core at the time of the explosion, while most of the He has been transported to the CSM through mass loss. \\citet{Meynet02} and \\citet{Meynet06} have recently calculated the evolution of metal-poor massive stars taking into account the effects of rotation-induced mixing, and shown that extensive mass loss occurs even in extremely metal-poor cases, in addition to significant enrichment of the N abundance near the surface. If this N is accelerated at the shock breakout of the SN explosion, a significant amount of Be can be produced through the reaction N + He $\\rightarrow$ $^9$Be, because of its low threshold $E_{\\rm th}$ ($\\sim 6$ MeV/A) and high cross section at peak $\\sigma_{\\rm p}$ ($\\sim 24$ mb) compared with other $^9$Be-producing reactions (eg. $E_{\\rm th} \\sim 8$ MeV/A and $\\sigma_{\\rm p} \\sim 8$ mb for O + He $\\rightarrow$ $^9$Be, see Fig. \\ref{fig-Ecs}). Indeed, recent observations indicate that the abundances of both Be and N may be enhanced in some metal-poor stars in which $^6$Li is detected \\citep{Primas00a, Primas00b,Israelian04}. In this Letter, we focus on the interactions ${\\rm He,CNO} + {\\rm He, N} \\rightarrow {\\rm LiBeB}$ occurring between the ejecta accelerated by the energetic SN explosion and the CSM, in order to account for the abundances in such very metal-poor stars.  We summarize recent observations of the relevant elements in very metal-poor stars in \\S \\ref{sec-obs}. In \\S \\ref{sec-SNIC} we discuss SN explosions of stars with stripped He layers embedded in the CSM. Our calculations and results are shown in \\S 4 with an appropriate parameter range. Conclusions are presented in \\S \\ref{sec-disc}. ",
        "conclusions": "\\label{sec-disc} In this Letter we have proposed a new mechanism to produce the light elements, especially $^6$Li and Be recently observed in metal-poor stars. This is based on recent theoretical findings that rotating, metal-poor massive stars can lose substantial amounts of their envelopes, and end up with SNe Ic with a small amount of He and N in the outermost layers. The outer layers composed of He and CNO are accelerated at the shock breakout and undergo spallation and He-He fusion reactions to produce mainly $^6$Li and $^9$Be. Our calculations show that if $\\sim 0.01-0.1\\,M_\\odot$ of the He layer is remaining on the core in a SN explosion of a massive star with $M_{\\rm ms}>30\\,M_\\odot$, then the observed abundance ratios $^6$Li/O can be reproduced. If the He layer contains a small amount ($\\sim 0.5 - 1$\\%) of N, we can also reproduce $^9$Be/O.  Compared with the standard picture of cosmic ray shock acceleration by normal supernovae, our high yield of $^6$Li results from three crucial differences. First, we consider a large explosion energy, of order $10^{52}$ erg, as observed in some SNe such as SN 1998bw. Second, our mechanism considers the direct interaction of fast ejecta with the CSM, and does not involve an efficiency factor for cosmic ray acceleration and is not affected by losses or escape during ISM propagation. Third, the energy distribution of the accelerated particles is a very steeply falling power-law, with spectral index $\\sim 4.6$ as opposed to $2$ for shock accelerated particles. This means that most of the energy is contained in the lowest energy portions, where the cross sections for both ${\\rm He + He} \\rightarrow {\\rm ^6Li}$ and  ${\\rm N + He} \\rightarrow {\\rm ^9Be}$ peak (see Fig. \\ref{fig-Ecs}). In fact, the total energy of particles above 10 MeV/nucleon in our fiducial model is $8\\times10^{50}$ erg, much higher than is achievable in the standard picture, 10-30\\% of $10^{51}$ erg.  One may claim that $^6$Li, and probably also $^9$Be \\citep{GPerez05}, in metal-poor stars may have been depleted from their initial values, since the observed abundances of $^7$Li/H are a few times smaller than that predicted by standard big-bang nucleosynthesis, and $^6$Li is more fragile than $^7$Li \\citep{Asplund05}. The actual survival fraction of $^6$Li is difficult to evaluate reliably because the pre-main-sequence Li destruction is sensitive to convection, which cannot be modeled without free parameters. However, depletion factors as  large as 10 may still be compatible with our picture. The estimates in the preceding section is based on an explosion energy of the order of $10^{52}$ ergs. The mass of ejecta with energy per nucleon above a certain value scales very strongly with the explosion energy, $E_{\\rm ex}$ as $E_{\\rm ex}^{3.4}$ \\citep{Nakamura04}. Thus a SN with the explosion energy 2 times higher than that considered in the preceding section will be able to produce $\\sim 10$ times larger amounts of light elements depending on the distribution of He and N. Nevertheless, if such very energetic explosions turn out to be rare, the mechanism proposed here may play a rather limited role compared to other potential LiBeB production processes.  Both the degree of mass loss and the amount of N enrichment in the progenitor star are expected to be sensitive to its initial rotation speed \\citep{Meynet02, Meynet06}. The actual rotation speed is presumably distributed over a wide range, as is the explosion energy and the extent of mass loss at the time of the explosion. Therefore, dispersions in the $^6$Li and Be abundances are expected, and current observations suggest that this may indeed be the case for $^6$Li \\citep{Aoki04,Asplund05,Inoue05} as well as Be \\citep{Boesgaard05}. Our scenario predicts a close relation between Be and N, since Be arises directly as a consequence of N spallation. A looser correlation between $^6$Li and N is also expected, as effective $^6$Li production requires sufficient mass loss, which in turn implies significant N enrichment. Further observations of $^6$Li, Be and N for a large sample of metal-poor stars should provide definitive tests. Note that such correlations are not expected for other scenarios which involve mainly $^6$Li production only, such as structure formation cosmic rays \\citep{Suzuki02,Rollinde05}, active galactic nuclei outflows \\citep{Prantzos06,Nath06}, and production processes in the early universe \\citep{Jedamzik05}. It is also mentioned that mass loss onto companion stars in binary systems may represent an additional pathway to our scenario, provided that sufficiently thick CSM is remaining at the time of the explosion. We reiterate that more observational data is necessary to elucidate what fraction of the LiBeB abundances seen in halo stars of different metallicity can be explained by the mechanism proposed here.  We are grateful to Georges Meynet, Raphael Hirschi, Takeru K. Suzuki and Sean Ryan for valuable discussions. We also acknowledge the contributions of an anonymous referee, which led to clarification of a number of points in our original manuscript. This work has been partially supported by the grant in aid (16540213, 17740108) of the Ministry of Education, Science, Culture, and Sports in Japan."
    },
    "0601/astro-ph0601309_arXiv.txt": {
        "abstract": "We have measured accurate near-infrared magnitudes in the J and K bands of 39 Cepheid variables in the irregular Local Group galaxy IC 1613 with well-determined periods and optical VI light curves. Using the template light curve approach of Soszy{\\'n}ski, Gieren and Pietrzy{\\'n}ski, accurate mean magnitudes were obtained from these data which allowed to determine the distance to IC 1613 relative to the LMC from a multiwavelength period-luminosity solution in the optical VI and near-IR JK bands, with an unprecedented accuracy. Our result for the IC 1613 distance is $(m-M)_{0} = 24.291 \\pm 0.014$ (random error) mag, with an additional systematic uncertainty smaller than 2\\%. From our multiwavelength approach, we find for the total (average) reddening to the IC 1613 Cepheids $E(B-V) = 0.090 \\pm 0.007$ mag, which is significantly higher than the foreground reddening of about 0.03 mag, showing the presence of appreciable dust extinction inside the galaxy. Our data suggest that the extinction law in IC 1613 is very similar to the galactic one. Our distance result agrees, within the uncertainties, with two earlier infrared Cepheid studies in this galaxy of Macri et al. (from HST data on 4 Cepheids), and McAlary et al. (from ground-based H-band photometry of 10 Cepheids), but our result has reduced the total uncertainty on the distance to IC 1613 (relative to the LMC) to less than 3\\%. With distances to nearby galaxies from Cepheid infrared photometry at this level of accuracy, which are currently being obtained in our Araucaria Project, it seems possible to significantly reduce the systematic uncertainty of the Hubble constant as derived from the HST Key Project approach, by improving the calibration of the metallicity effect on PL relation zero points, and by improving the distance determination to the LMC. ",
        "introduction": "Cepheid variables are the most important standard candles to calibrate the first rungs of the extragalactic distance ladder, out to some 30 Mpc. As young stars, Cepheids tend to lie in dusty regions in their spiral or irregular parent galaxies. As a consequence, Cepheid distances derived from the period-luminosity (PL) relation in optical photometric bands are quite sensitive to a precise knowledge of the total reddening, foreground and intrinsic, of their parent galaxy. While the galactic foreground reddening towards any direction in the sky is usually well established, particularly in directions far away from the galactic equator, the correct assessment of the reddening produced by dust extinction {\\it intrinsic to the host galaxy} is usually a difficult task, and in most work on Cepheid distances based on optical data such an intrinsic contribution to the reddening has simply been ignored. Just for this one particular reason, it is clear that more accurate Cepheid distances to galaxies can be derived in near-infrared passbands, where dust absorption is small as compared to visual wavelengths, and the distance results become increasingly independent of errors in the assumed total reddenings. Efforts along these lines have started in the early eighties with the pioneering work of McGonegal et al. (1982), and Welch et al. (1985). Yet, an important obstacle to carry out accurate Cepheid distance work in the infrared has been, until very recently, the lack of well-calibrated fiducial PL relations in the near-infrared JHK bands. This problem has now been solved by the work of Persson et al. (2004) who provided such well-calibrated relations for the LMC. Very recently, Gieren et al. (2005a) have also provided well-calibrated PL relations in the JHK bands for Milky Way Cepheids, which agree with the corresponding LMC relations when an improved version of their infrared surface brightness technique (Gieren, Fouqu{\\'e} \\& G{\\'o}mez, 1997, 1998) is used.  In the {\\it Araucaria Project}, started by our group some time ago (Gieren et al. 2005c), we have conducted surveys for Cepheid variables in a number of galaxies in the Local Group, and in the more distant Sculptor Group in order to investigate the effect of environmental properties on the PL relation, and to improve the accuracy of Cepheids as distance indicators. While we are discovering Cepheids in optical photometric bands, where these stars are rather easy to detect due to their relatively large amplitudes and typical light curve shapes (e.g. Pietrzynski et al. 2002, 2004), the main goal of the program is to undertake near-infrared follow-up imaging of selected subsamples of Cepheids in our target galaxies to obtain accurate reddening information, and thus to obtain more accurate distances than what is possible from optical (VI) data alone. Near-infrared PL relations from such Cepheids with existing information on their periods, and V and/or I light curves can be obtained very economically because it is possible to obtain accurate mean JHK magnitudes for these stars from just one single-phase observation from the template light curve approach of Soszynski, Gieren and Pietrzynski (2005). The success of this approach was recently demonstrated in the case of the Sculptor galaxy NGC 300 (Gieren et al. 2005b). For this galaxy, a combination of the PL relations obtained in the optical VI and infrared JK bands has allowed to determine a distance which is practically unaffected by any remaining uncertainty on reddening. It has also been shown in that paper that from the combined optical/near-infrared approach a total uncertainty as small as 3 \\% can be obtained for the Cepheid distance  (as measured relative to the LMC) for such a relatively nearby (2 Mpc) galaxy.  In the present paper, we are applying the same approach to the Local Group dwarf irregular galaxy IC 1613. IC 1613 is a very important galaxy in the Araucaria Project because of the very low metallicity of its young stellar population close to -1.0 dex (Skillman et al. 2003), making it the lowest-metallicity galaxy in our sample. It is therefore a key object in our effort to determine the effect of metallicity on the Cepheid PL relation, and on other stellar distance indicators, like blue supergiant stars (Kudritzki et al. 2003). A first survey for Cepheid variables in IC 1613 was carried out by Sandage (1971) who used photographic images previously obtained by Baade. Modern work on the Cepheid PL relation in IC 1613, in the optical V and I bands, has been carried out by the OGLE Project (Udalski et al. 2001) who has discovered many new, previously unknown Cepheids in this irregular galaxy. More recently, Antonello et al. (2006) have extended this work to the B and R bands. From the work of the OGLE group, it could be established that the slope of the PL relation in optical bands is identical to the slope observed for the more metal-rich LMC Cepheids, arguing for a metallicity- independent slope of the PL relation. In the near-infrared, a pioneering paper on the distance of IC 1613 from H-band photometry of 10 Cepheids was published by McAlary et al. (1984) already 20 years ago; however, the uncertainty on this distance result was rather large due to the technical difficulties to obtain accurate IR photometry at those times, and the lack of an accurate calibrating PL relation. Much more recently, Macri et al. (2001) determined a near-infrared Cepheid distance to IC 1613 from H-band photometry of four variables obtained with HST/NICMOS. The accuracy of this determination suffers, however, from the very small number of stars used in the PL solution. A main goal of the present study was to derive {\\it truly accurate near-infrared PL relations for IC 1613}, based on a large number of well-observed and well-selected Cepheids (see section 3.2.), and this way reduce the current uncertainty on the distance to IC 1613 to the very small level of 3-5\\% we have achieved in our previous study of NGC 300.  We have organized this paper in the following way: in section 2, we describe the observations, reductions and calibration of our data; in section 3, we present the calibrated infrared mean magnitudes of the Cepheids in our selected fields in IC 1613, and determine the distance and reddening; in section 4, we discuss our results; and in section 5, we summarize the main results of this work, and present some conclusions. ",
        "conclusions": "We have measured accurate NIR magnitudes in the J and K bands for 39 Cepheid variables in the Local Group galaxy IC 1613 with well-determined periods and optical (VI) light curves. Mean magnitudes in J and K were derived for these variables using the single-phase approach of Soszy{\\'n}ski et al. (2005). After carefully cleaning the Cepheid list from blended objects, Population II variables and overtone pulsators, we have determined accurate PL relations. Fits to these observed relations were made using the slopes of the LMC relations determined by Persson et al. (2004), which gave an excellent representation of the IC 1613 data, providing for the first time solid evidence that the slope of the Cepheid PL relation is independent of metallicity down to the low metallicity of -1.0 dex of the IC 1613 young population in the near-infrared domain, too. This is in agreement with the theoretical predictions of Bono et al. (1999). By combining the zero points of the J and K band PL relations in our study with the ones derived by Persson et al. for the LMC, we derive relative distance moduli of IC 1613 with respect to the LMC in both bands. Combining these infrared moduli with the distance moduli previously derived by Udalski et al. (2001) in V and I, we determine the total (average) reddening of the Cepheids in IC 1613, and the true distance modulus of this galaxy with an unprecedented accuracy. For the reddening, we find E(B-V) = 0.090 $\\pm$ 0.007 mag, and for the true distance modulus of IC 1613 from our multiwavelength approach we obtain 24.291 $\\pm$ 0.014 mag (random error). As in the case of our study of NGC 300 with the same method, we find evidence that there is a significant contribution to the total reddening from dust absorption {\\it intrinsic} to IC 1613, which had been neglected in the previous Cepheid distance work on this galaxy. The excellent fit of the distance moduli to the assumed galactic extinction law suggests that the interstellar extinction in this small irregular galaxy follows closely the galactic law.  We show that our derived Cepheid distance is very insensitive to systematic uncertainties caused by the fiducial PL relations used in our fits, possible inhomogeneous filling of the instability strip by our Cepheid sample, and problems with blending of the variables. Any remaining influence of the uncertainty of reddening on our distance result is negligible. All these possible sources of error contribute less to the total systematic uncertainty of our result than the two dominant sources of error, which are the zero points of our JK photometry ($\\pm$ 0.03 mag), and the distance to the LMC, which we have {\\it adopted} as 18.50 mag, and whose current uncertainty seems in the order of $\\pm$ 0.10 mag.  Our distance determination for IC 1613 is in reasonable agreement with the previous determination of Macri et al. (2001) from HST H-band photometry of four Cepheids, 24.43 $\\pm$ 0.08 mag. We attribute the 0.14 mag difference mainly to the small number of stars available to Macri et al. in their study. Our new distance determination is also in very good agreement with the very early infrared work of McAlary et al. (1984); these authors had obtained a distance modulus of 24.31 $\\pm$ 0.12 from H-band data of 10 Cepheids in IC 1613. A change of their assumed reddening of 0.03 mag to our larger value found in this study still produces excellent agreement of their result with ours.  The distance to IC 1613 was also determined  by Dolphin et al. (2001) using HST data, and employing a number of distance  indicators (TRGB, red clump stars, RR Lyrae stars, Cepheids). We note that their Cepheid sample was very small, which can clearly lead to spurious results. Udalski et al. 2001 have observed an order of magnitude larger sample of Cepheids in this galaxy and  showed that all these different distance indicators yield consistent distances to this galaxy. Those distance measurements are all in very good agreement with the distance of IC 1613 obtained in this study if the revised reddening of 0.09 mag found in this study, and a LMC true distance modulus of 18.5 mag are assumed.  As a final conclusion, we have produced in this work a determination of the Cepheid distance to IC 1613 whose random error is of the order of 1\\%, and estimated systematic error (excluding the uncertainty of the adopted LMC distance) is in the order of 2\\%. This accuracy, when combined with distance determinations of similar accuracy we pretend to obtain for most of the other galaxies of the Araucaria Project, should enable us to pin down the metallicity dependence of the PL relation zero points in the different optical and near-infrared photometric bands with the 1-2\\% accuracy needed to produce a significant improvement in the determination of the Hubble constant from distance determinations to galaxies in the nearby field from their Cepheid populations, which is the approach used in the HST Key Project of Freedman et al. (2001). Our work in the Araucaria Project should therefore strongly contribute, in the very near future, to make best use of the past work of the Key Project team."
    },
    "0601/astro-ph0601623_arXiv.txt": {
        "abstract": "{} {In this paper we confirm the presence of a globally-ordered, kG-strength magnetic field in the photosphere of the young O star {$\\theta^1$~Orionis~C}, and examine the properties of its optical line profile variations.} {A new series of high-resolution MuSiCoS Stokes $V$ and $I$ spectra has been acquired which samples approximately uniformly the rotational cycle of {$\\theta^1$~Orionis~C}. Using the Least-Squares Deconvolution (LSD) multiline technique, we have succeeded in detecting variable Stokes $V$ Zeeman signatures associated with the LSD mean line profile. These signatures have been modeled to determine the magnetic field geometry. We have furthermore examined the profile variations of lines formed in both the wind and photosphere using dynamic spectra.} {Based on spectrum synthesis fitting of the LSD profiles, we determine that the polar strength of the magnetic dipole component is $1150 \\la B_{\\rm d}\\la 1800$~G and that the magnetic obliquity is $27\\degr \\la \\beta \\la 68\\degr$, assuming $i=45\\pm 20\\degr$. The best-fit values for $i=45\\degr$ are $B_{\\rm d} = 1300 \\pm 150\\ (1\\sigma)$~G and $\\beta = 50\\degr \\pm 6\\degr\\ (1\\sigma)$. Our data confirm the previous detection of a magnetic field in this star, and furthermore demonstrate the sinusoidal variability of the longitudinal field and accurately determine the phases and intensities of the magnetic extrema. The analysis of ``photospheric'' and ``wind'' line profile variations supports previous reports of the optical spectroscopic characteristics, and provides evidence for infall of material within the magnetic equatorial plane.}{} ",
        "introduction": "The detection of magnetic fields in O-type stars has proven to be a remarkably challenging observational problem \\citep[e.g.,][]{Donati01}. The apparent absence of magnetic signatures has often been interpreted as a selection effect, since { about 5\\% of} B- and A-type stars \\citep[see, e.g.,][]{Johnson04} do display organized magnetic fields with disk-averaged strengths ranging from a few hundred G to several tens of kG \\citep{Mathys01}. These fields are  believed to be the fossil remnants of either interstellar fields swept up during the star formation process or fields produced by a pre-main sequence envelope dynamo that has since turned off\\footnote{Credible quantitative models invoking {\\em contemporaneous} dynamos operating in the convective core have also been proposed (e.g. Charbonneau \\& MacGregor 2001), but these models generally have significant difficulty explaining the intensities and topologies of the observed fields, their diversity, lack of correlation of field with angular velocity, as well as the young ages of many magnetic stars.}. As there is no particular reason to suspect that similar processes should not occur during the formation of more massive stars, it seems reasonable {\\it a priori} to expect that magnetic fields of similar structure and intensity should also exist in O-type stars.  In the absence of direct magnetic measurements, this view has been supported by the wide-spread occurrence of variability in the winds of O-type stars, which may provide indirect evidence for the presence of dynamically important magnetic fields in their atmospheres. During the past 20 years, sustained optical and UV spectroscopic observations have shown that the winds of O-type stars are highly structured, { exhibiting both coherent and stochastic behaviour,} as well as cyclical variability on timescales ranging from hours to days. \\footnote{See, e.g., \\citet{Fullerton03} for a recent review of the characteristics of variability in hot-stars winds.} A key result of this work has been the conclusion that a major component of this variability results from rotational modulation of structures imposed on the wind by some deep-seated process. Magnetic fields are presently considered a likely source of such structures (e.g. \\citep{Cranmer96})\\footnote{Magnetic fields are also seen by some investigators as an ingredient necessary to explain intrinsic X-ray fluxes and non-thermal radio emission from massive stars.}.  Although by no means prototypical, $\\theta^1$~Orionis~C (HD~37022; HR~1895) is perhaps the best example of an O-type star with distinctive, periodic variability of its spectroscopic stellar wind features. It is a { very} young, peculiar $\\sim$O7 star, and the brightest and hottest member of the Orion Nebula Cluster. It exhibits strictly periodic spectroscopic variability which is strongly suggestive of a magnetic rotator: H$\\alpha$ and {\\ion{He}{ii} $\\lambda$4686} emission, peculiar ultraviolet {\\ion{C}{iv} $\\lambda\\lambda$1548, 1550} wind lines, photospheric absorption lines, and ROSAT X-ray emission all appear to vary with a single  well-defined period of $15.422\\pm0.002$ d \\citep{Stahl93,Stahl96,Walborn94,Gagne97}. The similarity of this behaviour to that of  some magnetic A- and B-type stars \\citep[see, e.g.,][]{Shore90} has motivated the suggestion first made by \\citet{Stahl96} that {$\\theta^1$~Ori~C} also hosts a fossil magnetic field which confines the stellar wind and produces the observed variability by rotational modulation. { Such a phenomenon appears to have been first suggested (for the O star $\\zeta$~Pup) by Moffat \\& Michaud (1981) (although no field has ever been detected in that star).}  In order to test this suggestion, \\citet{Donati99b} obtained longitudinal magnetic field measurements of {$\\theta^1$~Ori~C}, but failed to detect the presence of a photospheric field. However, in a seminal paper, \\citet{Donati02} reported the detection of a fossil magnetic field in the photosphere of {$\\theta^1$~Ori~C} based on 5 measurements of the Stokes $V$ profiles of selected photospheric absorption lines. This detection provided the first empirical support that dynamically important magnetic fields exist in O-type stars, and that their behaviour is directly linked to spectroscopic variability.\\footnote{Recently, \\citet{Donati05} reported the detection of a 1.5~kG dipolar magnetic field in the O-type spectrum variable HD~191612 [Of?p], which exhibits stellar wind variations with a period of 538 days. They suggested that HD~191612 represents an evolved version of {$\\theta^1$~Ori~C}.} It also represents the youngest main sequence star in which a fossil-type field has been detected \\citep[comparable in age to NGC 2244-334;][]{Bagnulo04}. Since {$\\theta^1$~Ori~C} is clearly a pivotal object, it is important that this detection be corroborated and the determination of its magnetic properties be refined.  The primary goal of the present paper is to report confirmation of the detection of a magnetic field in the photospheric lines of {$\\theta^1$~Ori~C}. Our new spectropolarimetric observations, which provide a substantially larger data set that samples the $15\\fd 422$ rotational period more fully, are described in \\S2. The analysis of the Least-Squares Deconvolved circular polarisation spectra, both in terms of the mean longitudinal magnetic field, {$\\langle B_z\\rangle$}, and the LSD mean Stokes $V$ profiles, is described in \\S3. In \\S4 we discuss the variability characteristics of features in the Stokes $I$ spectrum, and comment on consistency with both the magnetically-confined wind shock model of \\citet{Donati02} as well as very recent MHD simulations of magnetically-channelled winds by \\citet{udDoula02} and \\citet{Gagne05}. Finally, in \\S5 we summarise our results, and discuss implications for our understanding of the origin and evolution of magnetic fields in intermediate- and high-mass stars, and for our understanding of wind modulation in massive stars. ",
        "conclusions": "Using a new series of 45 Stokes $I$ and $V$ spectra obtained with the MuSiCoS spectropolarimeter at Pic du Midi observatory, we have detected the photospheric magnetic field of the young O7 Trapezium member $\\theta^1$~Ori C. We confirm and extend the conclusions of \\citet{Donati02}: that the Stokes $V$ variations are consistent with a dipolar magnetic field with a polar strength between 1150 and 1800 G, and an obliquity $27\\degr\\leq \\beta\\leq 68\\degr$ if $i=45\\pm 20\\degr$. Moreover, we demonstrate that the variation of the longitudinal magnetic field is sinusoidal to within the errors; the phase of maximum longitudinal field is $0.05\\pm 0.05$ according to the ephemeris of \\citet{Stahl96}; and the longitudinal field varies between $-35\\pm 85$~G and $+550\\pm 85$~G. We have also exploited our high-resolution Stokes $I$ spectra to study the cyclical variations of spectral absorption and emission lines formed in the photosphere and wind of $\\theta^1$~Ori C. We confirm the variability properties reported previously by, e.g., \\citet{Stahl96}, and highlight evidence that suggests the presence of infalling material in the magnetic equatorial plane, which is consistent with earlier suggestions of such phenomena by \\citet{Donati01} and the theoretical predictions of \\citet{Gagne05}.  $\\theta^1$~Ori C holds a special place in our understanding of magnetism in intermediate- and high-mass stars, because it is one of only two O-type stars in which a magnetic field has been detected and characterized at multiple epochs. It is also one of the youngest stars (with an age of about $10^6$ years) in which a magnetic field has been detected, demonstrating once again (e.g. Bagnulo et al. 2004) that magnetic fields are apparent at the surfaces of some intermediate and high mass stars at very early main sequence evolutionary stages. The confirmation of a field in {$\\theta^1$~Ori~C} extends the range of known stars hosting fossil magnetic fields by a factor of about 2 in effective temperature, and by { about 4} in stellar mass."
    },
    "0601/astro-ph0601145_arXiv.txt": {
        "abstract": "We calculate photometric redshifts from the Sloan Digital Sky Survey Main Galaxy Sample, The Galaxy Evolution Explorer All Sky Survey, and The Two Micron All Sky Survey using two new training--set methods. We utilize the broadband photometry from the three surveys alongside Sloan Digital Sky Survey measures of photometric quality and galaxy morphology. Our first training--set method draws from the theory of ensemble learning while the second employs Gaussian process regression both of which allow for the estimation of redshift along with a measure of uncertainty in the estimation. The Gaussian process models the data very effectively with small training samples of approximately 1000 points or less.  These two methods are compared to a well known Artificial Neural Network training--set method and to simple linear and quadratic regression. Our results show that robust photometric redshift errors as low as 0.02 RMS can regularly be obtained.  We also demonstrate the need to provide confidence bands on the error estimation made by both classes of models.  Our results indicate that variations due to the optimization procedure used for almost all neural networks, combined with the variations due to the data sample, can produce models with variations in accuracy that span an order of magnitude. A key contribution of this paper is to quantify the variability in the quality of results as a function of model and training sample.  We show how simply choosing the ``best\" model given a data set and model class can produce misleading results. ",
        "introduction": "Using broadband photometry in multiple filters to estimate redshifts of galaxies was likely first attempted by \\citet{Baum62} on 25 galaxies in nine broadband imaging filters in the visible and near--infrared range. Given the low throughput of spectrographs much is to be gained by attempting to estimate galaxy redshifts from broadband colors rather than from measurement of individual spectra. In the Sloan Digital Sky Survey \\citep[SDSS,][]{York2000} 100 million galaxies will have accurate broadband u,g,r,i,z photometry, but only 1 million galaxy redshifts from this sample will be measured. If a method can be found to obtain an accurate estimate of the redshift for the larger SDSS photometric catalog, rather than the smaller spectroscopic one, much better constraints on the formation and evolution of large--scale structural elements such as galaxy clusters, filaments, and walls and cosmological models in general \\citep[e.g.][]{BlakeBridle2005} may be achieved.  Two approaches, spectral energy distribution fitting (SED fitting: also known as ``template--fitting\") and the training--set method (TS method), have been used to obtain photometric redshifts over the past 30 years. In order to use TS methods galaxies with a similar range in magnitude and color over the same possible redshift range must be used to estimate the redshifts from the broadband colors measured. Since this type of data has not always been available SED fitting has historically been the preferred method \\citep[e.g.][]{Koo85,Loh86,Lanzetta96,Kodama99,Benitez00,Massarotti01,Babbedge04,Padmanabhan05} given the historically low numbers of galaxies with spectroscopically confirmed redshifts in deep photometric surveys of the universe. This is due to the fact that photometric surveys have always gone, and continue to go, deeper than is possible with spectroscopy.  Another alternative has been to use training sets consisting of a combination of both observed galaxy templates and those from galaxy evolution models \\citep[e.g. hyperz,][]{Bolzonella00}.  There are many approaches to SED fitting. For example, \\citet{Kodama99} use four--filter ({/it BVRI}) photometry and a Bayesian classifier using SED fitting which they have tested out to z=1 and claim is valid beyond this redshift.  The approach of \\citet{Benitez00} makes use of additional information such as the shape of the redshift distributions and fractions of different galaxy types. This may be helpful in instances where one has a limited sample size at large redshifts. However, all estimators, Bayesian or otherwise, can be biased due to small sample size effects.  TS methods rely on having a complete sample of galaxies in magnitude, color and redshift. Hence these methods have been restricted to relatively nearby z$<$1 surveys, such as the SDSS, rather than much deeper surveys such as the Hubble Deep Field \\citep{Williams96}. In fact, for redshifts above 1 there have not been sufficiently large and complete enough measured samples of galaxy redshifts, magnitudes and colors to use TS methods with much accuracy \\citep[e.g.][]{Wang98}. As well, the colors of galaxies change nearly monotonically up to z$=$1, but beyond this the color--redshift space becomes much more complex and simple linear and quadratic regression will fail. Hence SED fitting has been used almost exclusively for surveys of z$>$1.  See \\citet{Benitez00} for an excellent detailed discussion of the differences and similarities between these two commonly used approaches.  In the past 10 years a large number of empirical fitting techniques for TS methods have come into use and new techniques continue to be developed. Some examples of linear and non--linear methods include: 2nd and 3rd order polynomial fitting \\citep{Brunner97,Wang98,Budavari05}; quadratic polynomial fitting \\citep{Hsieh05,Connolly95}; support vector machines \\citep{Wadadekar05}; nearest neighbor and kd--trees \\citep{Csabai03}, and artificial neural networks \\citep{Firth03,Tagliaferri03,Ball04,Collister04,Vanzella04}.  We explore the problem of estimating redshifts from broadband photometric measurements using the idea of a virtual sensor \\citep{SrivastavaOza2005,SrivastavaStroeve2003}.  These methods allow for the estimation of unmeasured spectral phenomena based on learning the potentially nonlinear correlations between observed sets of spectral measurements. In the case of estimating redshifts, we can learn the nonlinear correlation between spectroscopically measured redshifts and broadband colors.  Statistically speaking, this amounts to building a regression model to estimate the photometric redshift.  However, the procedure is much more complex than a simple regression due to the  significant effort required for model building and validation.  The concept of virtual sensors applies to the entire chain of analytical steps leading up to the prediction of the redshift. Figure~\\ref{fig:figure1} shows a schematic of the assumptions behind a Virtual Sensor with a cartoon on the left and the real--world case with the five SDSS bandpasses and a sample galaxy spectrum overlaid on the right.  As a baseline comparison, results from a TS--based neural network package called ANNz \\citep{Collister04} are presented.  Linear and quadratic fits along the lines discussed in \\cite{Connolly95} are also presented.  Unlike all other previous work, we also discuss the application of bootstrap resampling \\citep{Efron79,ET93} for the linear, quadratic, and ANNz models.  We apply the TS methods discussed above to the SDSS five--color ({\\it ugriz}) imaging survey known as the Main Galaxy Sample \\citep[MGS,][]{Strauss02} which has a large calibration set of spectroscopic redshifts for the SDSS Data Release 2 \\citep[DR2,][]{Abazajian04} and SDSS Data Release 3 \\citep[DR3,][]{Abazajian05}. The Two Micron All Sky Survey \\citep[2MASS,][]{Skrutskie06}\\footnote{http://www.ipac.caltech.edu/2mass/} extended source catalog along with Galaxy Evolution Explorer \\citep[GALEX,][]{Martin05}\\footnote{http://www.galex.caltech.edu/} data are also used in conjunction with the SDSS where all three overlap to create a combined catalog for use with our TS methods.  The data sets used in our analysis are discussed in {\\S} 2, discussion of the photometry and spectroscopic quality of the data sets along with other photometric pipeline output properties of interest is given in {\\S} 3, the classification schemes used to obtain photometric redshifts are in {\\S} 4, comparison of the results takes place in {\\S} 5, and we summarize in {\\S} 6.  \\begin{table} \\begin{center} \\caption{Survey filters and characteristics \\label{tbl-1}} \\begin{tabular}{ccrrc} \\hline \\hline Bandpass& Survey& $\\lambda_{eff}$ & $\\Delta\\lambda$ & FWHM \\tablenotemark{1}\\\\ &       &    ({\\AA})      &     ({\\AA})     &   ($\\arcsec$)         \\\\ \\hline FUV     & GALEX & 1528       & 442       & 4.5 \\\\ NUV     & GALEX & 2271       & 1060      & 6.0 \\\\ u       & SDSS  & 3551       & 600       & 1-2 \\\\ g       & SDSS  & 4686       & 1400      & 1-2 \\\\ r       & SDSS  & 6165       & 1400      & 1-2 \\\\ i       & SDSS  & 7481       & 1500      & 1-2 \\\\ z       & SDSS  & 8931       & 1200      & 1-2 \\\\ j       & 2MASS & 12500      & 1620      & 2-3 \\\\ h       & 2MASS & 16500      & 2510      & 2-3 \\\\ k$_{s}$ & 2MASS & 21700      & 2620      & 2-3 \\\\ \\hline \\hline \\end{tabular} \\tablenotetext{1}{The Full Width at Half Maximum is dependent on the seeing at the time of the observation for ground based data.} \\end{center} \\end{table} ",
        "conclusions": ""
    },
    "0601/astro-ph0601373_arXiv.txt": {
        "abstract": "A simple 3D-reconstruction method for gamma-ray induced air showers is presented, which takes full advantage of the assets of a system of Atmospheric Cherenkov Telescopes combining stereoscopy and fine-grain imaging like the High Energy Stereoscopic System (\\hess). The rich information collected by the cameras allows to select electromagnetic showers on the basis of their rotational symmetry with respect to the incident direction, as well as of their relatively small lateral spread. In the framework of a 3D-model of the shower, its main parameters --- incident direction, shower core position on the ground, slant depth of shower maximum, average lateral spread of Cherenkov photon origins (or ``photosphere 3D-width'') and primary energy --- are fitted to the pixel contents of the different images. For gamma-ray showers, the photosphere 3D-width is found to scale with the slant depth of shower maximum, an effect related to the variation of the Cherenkov threshold with the altitude; this property allows to define a dimensionless quantity $\\omega$ (the ``reduced 3D-width''), which turns out to be an efficient and robust variable to discriminate gamma-rays from primary hadrons. In addition, the $\\omega$ distribution varies only slowly with the gamma-ray energy and is practically independent of the zenith angle. The performance of the method as applied to \\hess~is presented. Depending on the requirements imposed to reconstructed showers, the angular resolution at zenith varies from 0.04$^{\\circ}$ to 0.1$^{\\circ}$ and the spectral resolution in the same conditions from 15\\% to 20\\%.\\newline   Keywords: gamma-ray astronomy, H.E.S.S., stereoscopy, Cherenkov telescopes, 3D-reconstruction, analysis method\\\\  PACS-2003: 95.55.Ka, 95.75.-z\\\\ ",
        "introduction": "Important progress has recently been achieved in gamma-ray astronomy above 100~GeV due to the performance of new stereoscopic systems of Atmospheric Cherenkov Telescopes (ACT's) equipped with high-definition imaging cameras, such as CANGAROO-3 \\cite{cangaroo}, \\hess \\cite{hofmann} and VERITAS \\cite{veritas}. These setups bring significant improvements in flux sensitivity, angular resolution and energy resolution, with respect to previous experiments in ground-based gamma-ray astronomy. The increase in sensitivity is directly related to the capability of rejecting hadron-induced showers on the basis of the image shapes and, at least for known point-like sources, of shower directions. Most presently available results from ACT's were obtained with this last constraint, the signal being extracted from the distribution of the pointing angle (the so-called ``$\\alpha$ or $\\theta^2$ plots''). However, stereoscopic systems now provide images of extended sources --- e.g. supernova remnants \\cite{nature} --- and discover unexpected sources \\cite{hotspot}; in such studies, hadronic showers must be rejected {\\it without using the shower direction}. In the present article, we refer to this context as well as to that of point-like sources.  The stereoscopic observation directly provides a simple geometrical reconstruction method based on the Hillas parameters \\cite{hillas} from the different images; this method has been applied in the HEGRA experiment \\cite{Daum97} and is presently adapted to \\hess \\cite{wystan}. High-definition imaging provides additional constraints which have already proved very useful, even with a single telescope, as in the CAT experiment~\\cite{LeBohec98}; in particular, the longitudinal light profile in the image can be modeled, providing a likelihood parameter discriminating gamma-rays from hadrons as well as a simultaneous determination of the direction of the primary $\\gamma$-ray, of the impact parameter and of the shower energy. This method was recently extended to \\hess, both in single and multi-telescope modes \\cite{de Naurois03}.  The method described in this article is based on a simple 3D-modeling of the Cherenkov light emitted by an electromagnetic air shower. The rich information contained in several fine-grained images of such a shower provides enough constraints to allow an accurate reconstruction, even by means of a simple model incorporating the rotational symmetry of the electromagnetic cascade with respect to its incident direction. This allows to select $\\gamma$-ray-induced showers on the basis of only two criteria with a direct physical meaning: rotational symmetry and small lateral spread. In section \\ref{sec-recmeth}, the simple assumptions used in the shower model are described and justified. In section \\ref{sec-discri}, on the basis of simulated gamma-rays and of real \\hess~data from the blazar PKS~2155-304, it is shown that a single variable $\\omega$, directly related to the lateral spread of Cherenkov photon origins, is an ideal parameter for discriminating primary $\\gamma$-rays from hadrons. In section \\ref{sec-perfo}, the performance of the method in terms of gamma-ray reconstruction efficiency, of angular resolution and of hadronic rejection is evaluated. Finally, the spectral resolution is discussed. ",
        "conclusions": "The reconstruction method described above differs from the more traditional ones (ref. \\cite{wystan} and \\cite{de Naurois03}) in several aspects: \\begin{itemize} \\item The analysis is based on shower parameters in 3 dimensions, not on image parameters; in this way correlations between different stereoscopic views of the same shower are taken into account. \\item Gamma-ray candidates are selected on the basis of a few criteria based on physical properties: rotational symmetry of electromagnetic shower, depth of shower maximum, lateral spread of the Cherenkov photosphere at shower maximum. \\item The distribution of the ratio of the last two variables, namely the ``reduced 3D-width'' $\\omega$, is found to be almost independent of the $\\gamma$-ray energy and zenith angle and $\\gamma$-rays can be efficiently selected by requiring $0.8 \\times 10^{-3}< \\omega < 2 \\times 10^{-3}$ for all observing conditions. This criterium provides a $\\gamma$-ray/hadron discrimination based on the shower shape (thus relevant for the study of extended sources) and completely independent of simulations. \\end{itemize} With a minimal set of natural cuts, the selection efficiency for $\\gamma$-rays as applied to \\hess~data is rather uniform between 100~GeV and 10~TeV and of the order of 80\\% (shape criteria) and of 50\\% (shape and direction criteria); the angular resolution (from 0.04$^{\\circ}$ to 0.1$^{\\circ}$ at zenith) and the energy resolution (from 15\\% to 20\\% in the same conditions) are comparable to those obtained with the standard \\hess~analysis \\cite{wystan}. This is illustrated in a recent article on the blazar H2356-309~\\cite{martin} where both analyses were used: using the 3D-model, an excess of 715 $\\gamma$-ray candidates is obtained with a significance of $10.9\\,\\sigma$, to be compared with 591 $\\gamma$-ray candidates and a significance of $9.7\\,\\sigma$ with the H.E.S.S. standard analysis. Furthermore, requiring at least 3 triggering telescopes, the 3D-model yields a higher significance ($11.6\\,\\sigma$) while keeping 453 $\\gamma$-ray candidates. Thus, the reconstruction method explained above illustrates how the quality of the stereoscopy, characterized by the telescope multiplicity, improves the angular resolution and the hadronic rejection of the array, thus its sensitivity.  \\clearpage"
    },
    "0601/astro-ph0601690_arXiv.txt": {
        "abstract": "{} {To confirm, or rule out, the possible hot spot nature of two previously detected radio sources in the vicinity of the Cygnus X-3 microquasar.} {We present the results of a radio and near infrared exploration of the several arc-minute field around the well known galactic relativistic jet source Cygnus X-3 using the Very Large Array and the Calar Alto 3.5~m telescope.} {The data this paper is based on do not presently support the hot spot hypothesis.  Instead, our new observations suggest that these sources are most likely background or foreground objects. Actually, none of them appears to be even barely extended as would be expected if they were part of a bow shock structure. Our near infrared observations also include a search for extended emission in the Bracket $\\gamma$ (2.166 $\\mu$m) and $H_{2}$ (2.122 $\\mu$m) lines as possible tracers of shocked gas in the Cygnus X-3 surroundings. The results were similarly negative and the corresponding upper limits are reported.} {} ",
        "introduction": "Cygnus X-3 is among the most intensively studied microquasars in the Galaxy. The system is a high-mass X-ray binary with a WN Wolf-Rayet companion star (see e.g. \\cite{f99} and references therein) seen through a high interstellar absorption ($A_V$\\gtsima 10 mag) that renders the optical counterpart undetectable in the visual domain. Cygnus X-3 has been observed to undergo strong radio flares with flux density increments by almost three orders or magnitude above the normal quiescent level of $\\sim0.1$ Jy at cm wavelengths.  The first of them was the historic flare extensively described by \\cite{g72} and subsequent papers.  Relativistic jets from this microquasar have been reported soon after some of these events outflowing collimately in the North-South direction (see e.g. \\cite{marti01}, \\cite{miller04}). The reader is referred to these papers for a more detailed account on the flaring and sub-arcsecond properties of the source.  In recent times, concerns have arised about the possible effects of continuate energy and momentum injection into the interstellar medium (ISM) during the flaring lifetime of microquasar systems. These effects appear to be far more important than previously thought. The recent report of a `dark jet' in Cygnus X-1 by \\cite{ga2005} illustrates a likely case of a black hole microquasar silently evacuating a significant fraction of its accretion power into its surroundings and affecting them on a several pc scale.  In a previous paper, \\cite{marti05} have reported the existence of two possible hot spot candidates (HSCs) associated with Cygnus X-3, thus suggesting an analogy with Fanaroff-Riley type II radio galaxies. Hereafter and following their original designation, we will refer here to these objects as the hot spot candidate North and South (HSCN and HSCS), respectively.  They appeared as faint radio sources with non-thermal spectrum at angular distances of 7\\prp 07 and 4\\prp 36 from Cygnus X-3. Moreover, the line joining them was within one degre of the almost North-South position angle of the inner arc-second radio jets as measured by \\cite{marti01}.  The \\cite{marti05} results, however, were hampered by the fact that the HSCs were far from the phase centre of the interferometric array and consequently suffered from significant bandwidth smearing. Searches for near infrared counterparts provided also negative results within the sensitivity of Two Micron All Sky Survey (2MASS).  In this paper, we report new radio and near infrared observations of both the HSCs and the Cygnus X-3 nearby environments. The improved observational data do not confirm the proposed hot spot nature and indicate that they are most likely background or foreground objects. This fact leaves open again for Cygnus X-3 the issue of searching for signatures of energy deposition from its relativistic jets into the ISM. In the following sections we describe and discuss our recent observational work that finally leads us to such different conclusion. ",
        "conclusions": "After conducting intensive radio and near infrared observations of Cygnus X-3 and its surroundings, our conclusions can be summarized as follows:  \\begin{enumerate}  \\item The previously proposed hot spot nature of two nearby radio sources in almost perfect alignment with the relativistic jets of this microquasar has not been possible to be confirmed.  The absence of clear indications for these objects being hot spots, such as extended emission and no stellar-like counterpart for both of them, leads us to be very cautious about our former hypothesis.  \\item We have also presented deep line narrow band images of the Cygnus X-3 field and provided strong upper limits for the brightness of extended emission due to possible shocked gas.  \\item The observational data collected and reported in the present paper clearly show that searching for signatures of jet-ISM interaction is a very difficult task which would require a more sensitive sampling on a wide range of angular scales.  \\item Finally, a comparison of the Cygnus X-3 stellar profile to that of other stars in the field does not indicate evidences of additional components within \\gtsima1$^{\\prime\\prime}$ at near infrared wavelegths.   \\end{enumerate}"
    },
    "0601/astro-ph0601529_arXiv.txt": {
        "abstract": "Determining temperatures in molecular clouds from ratios of CO rotational lines or from ratios of continuum emission in different wavelength bands suffers from reduced temperature sensitivity in the high-temperature limit.  In theory, the ratio of far-IR, submillimeter, or millimeter continuum to that of a $\\cO$ (or $\\Co$) rotational line can place reliable upper limits on the temperature of the dust and molecular gas.  Consequently,  far-infrared continuum data from the {\\it COBE}/{\\it DIRBE} instrument and Nagoya 4-m $\\cOone$ spectral line data were used to plot 240$\\um$/$\\cOone$ intensity ratios against 140$\\um$/240$\\um$ dust color temperatures, allowing us to constrain the multiparsec-scale physical conditions in the Orion$\\,$A and B molecular clouds.  The best-fitting models to the Orion clouds consist of two components: a component near the surface of the clouds that is heated primarily by a very large-scale (i.e. $\\sim 1\\,$kpc) interstellar radiation field and a component deeper within the clouds.  The former has a fixed temperature and the latter has a range of temperatures that varies from one sightline to another. The models require a dust-gas temperature difference of 0$\\pm 2\\,$K and suggest that 40-50\\% of the Orion clouds are in the form of dust and gas with temperatures between 3 and 10$\\,$K.  These results have a number implications that are discussed in detail in later papers.   These include stronger dust-gas thermal coupling and higher Galactic-scale molecular gas temperatures than are usually accepted, an improved explanation for the N(H$_2$)/I(CO) conversion factor, and ruling out one dust grain alignment mechanism. ",
        "introduction": "}  Development of a complete theory of star formation requires understanding the physical conditions of the molecular gas in which stars form, both on sub-parsec and multi-parsec scales.  On the latter size scales, observations of molecular clouds can provide insights into large-scale patterns of star formation such as in spiral arms, rings, or in starburst regions of external galaxies \\citep[e.g., see][] {Rand92, Vogel88, Aalto99, Kuno95, Wong00, Sheth00, Seaq98, Kiku98, Pag97, Gus93, Rieu92, Wild92, Harris91, Tilanus91, Turner91, Verter89, Lo87, Sorai00, GB00, Harr99, W91, Sage91, Scov85, Nish01, Nish01a, Kennicutt89, Martin01, Rownd99, Bryant99, Mulchaey97, Aalto95, W93, Puxley90, Verter87, Verter88}.  The diagnostics of molecular cloud physical, or excitation, conditions are usually obtained from observed molecular spectral lines.  Despite being composed almost entirely of molecular hydrogen, only very warm ($\\gsim few\\times 100\\unit K$) gas in interstellar molecular clouds is sampled {\\it in emission\\/} with the purely rotational lines of $\\rm H_2$ \\citep{Rodriguez01, Rosenthal00, Wright99, Parmar91} and the ro-vibrational lines sample even warmer ($\\gsim 2000\\unit K$) or fluorescently excited $\\rm H_2$ \\citep[e.g.,][]{Takami00, Fuente99, Moorwood90, Moorwood88, Black87}.  This is because the purely rotational transitions are quadrupolar and only occur between widely spaced ($\\rm h\\nu/k\\geq 510\\,K$) levels \\citep[for a more detailed discussion, see] [and references therein]{W99}.   Consequently, estimating the properties of the bulk of the molecular gas, which is cold \\citep[kinetic temperatures of $\\sim 5$--20$\\,$K, see, e.g.,][]{Sanders85}, requires observing lines of trace molecules with dipole transitions and closely spaced levels.  The molecule CO and its isotopologues, like $\\cO$ and $\\Co$, are the usual molecules of choice for estimating the large-scale properties of molecular clouds, because the most common isotopologue of CO --- $\\COex$ --- is the most abundant molecule after $\\rm H_2$, the levels are closely spaced ($\\rm h\\nu/k = 5.5\\,K$ for $\\Jone$ of $\\COex$), and the critical density of the lowest transition is low (i.e. $\\nH\\simeq 3\\times 10^3\\,cm^{-3}$ for $\\Jone$).  These characteristics of CO ensure that the lines are often strong --- with radiation temperatures or $\\Tr$ (i.e., Rayleigh-Jeans brightness temperatures as opposed to Planck brightness temperatures) of a few Kelvins --- for $\\COone$ observations of clouds in our Galaxy, \\citep[e.g.,][]{Sanders85} and occur throughout most of the molecular ISM.  Observations of lines of other molecules are also important for coming to a holistic picture of molecular clouds.  For example, the rotational lines of CS sample the dense gas ($\\nH\\gsim 10^5\\, cm^{-3}$) most directly associated with star formation \\citep{ELada92}.  Nevertheless, because the rotational lines of CO sample the excitation conditions in the bulk of the molecular gas, these lines provide insights into the environment in which the star-forming dense gas exists and, consequently, providing insights into star formation on the scales of many parsecs. Multi-parsec-scale observations of the lines of CO and its isotopologues, $\\cO$ and $\\Co$, are carried out on external galaxies (e.g., see the references listed near the beginning of this paragraph) or by low-resolution (few arcminutes), wide-coverage (many degrees) mapping of Galactic molecular clouds \\citep[e.g., see][]{Lee01, Dame01, Dame87, Dame94, Grabelsky88, Bronfman88, Cohen86, Sakamoto95, Robinson88, Sanders84, Tachihara00, Mizuno98, Kawamura98, Heyer98, Boulanger98, Onishi96, Oka98, Nagahama98, Sakamoto97a, Sakamoto94, Bally87, Maddalena86}.  Besides the usually observed main isotopologue $\\CO$ (written this way to explicitly distinguish it from the rarer isotopologues), the lines of the rarer isotopologues $\\cO$ (short for $\\cOex$) and $\\Co$ (short for $\\Coex$) furnish very important additional information on molecular cloud structure.  The lower-J lines of $\\CO$ are often optically thick and cannot probe deeply into the substructures, such as clumps or filaments, of molecular clouds.  Also, the opacity of the $\\CO$ lines reduces their sensitivity to the effects of changing density and changing radiative trapping.  Therefore, the lines of $\\cO$ or $\\Co$ are crucial for realistically constraining the molecular cloud excitation conditions, because these lines, being from lower abundance isotopologues, are often optically thin on multi-parsec scales.  This greater sensitivity of the $\\cO$ and $\\Co$ lines to cloud structure and physical properties makes them capable of identifying interesting regions within molecular clouds.  Of particular interest is locating regions of warm (e.g., $\\gsim 50$--$100\\unit K$) molecular gas, because energetic events and massive stars associated with recent star formation heat the surrounding molecular gas.  Such gas is sometimes identified from ratios of strengths of CO rotational lines \\citep[e.g.,][]{Wilson01, Plume00, Howe93, Graf93, Graf90, Boreiko89, Fixsen99, Harris85, Harr99, W91, Gus93, Wild92, Harris91}.  Examples of how two of the line ratios involving the three lowest rotational lines of CO vary with gas kinetic temperature, $\\Tk$, in local thermodynamic equilibrium or LTE are shown in Figure~\\ref{fig1}.  All the curves in Figure~\\ref{fig1} show the same trend: the line ratios lose their sensitivity to $\\Tk$ in the high-$\\Tk$, or Rayleigh-Jeans, limit. This is a consequence of the line ratios depending on ratios of two Planck functions at different frequencies.  The maximum $\\Tk$ to which a given line ratio is sensitive can be quantified by adopting a calibration uncertainty for the line ratios.  If we adopt the optimistic value of 20\\% for the uncertainties in the ratios, then a ratio that has an observed value $\\simeq 80\\%$ of its asymptotic value in the Rayleigh-Jeans limit is observationally indistinguishable from this asymptotic value. Similarly, the value of the kinetic temperature at this 80\\% level is the maximum kinetic temperature for which the given ratio is sensitive.  Two important conclusions can be drawn from the plots of Figure~\\ref{fig1}: \\begin{enumerate} \\item Optically thin lines, like those of $\\cO$, have a greater range of sensitivity to $\\Tk$ than do optically thick lines, like those of $\\CO$. Specifically, $\\Tr(\\Jtwo)/\\Tr(\\Jone)$ loses $\\Tk$ sensitivity above 46$\\,$K in the optically thin case and above 10$\\,$K in the optically thick case. \\item Ratios involving higher-J lines have a greater range of temperature sensitivity than those with lower-J lines. For $\\Tr(\\Jthree)/\\Tr(\\Jtwo)$ the maximum $\\Tk$ values are 73$\\,$K and 13$\\,$K for the optically thin and thick cases respectively.  These values are higher than the corresponding values for the $\\Tr(\\Jtwo)/\\Tr(\\Jone)$ ratio, especially for the optically thin case. \\end{enumerate} It is clear then that unambiguous identification of warm molecular gas requires observations of high-J lines of $\\CO$ and, sometimes, the optically thin lines of $\\cO$ (see references cited at the end of the previous paragraph).  (It should be mentioned that simply observing the optically thick lines of $\\CO$ does not automatically give the kinetic temperature, because the observed lines are not necessarily thermalized and the gas does not necessarily fill the beam in the velocity interval about the line peak.)  However, identifying warm gas in this manner can be complicated.  One complication is that observing higher and higher lines up the rotational ladder becomes increasingly difficult with ground-based observations, because the atmosphere becomes increasingly opaque, on average, and because sub-millimeter receivers become increasingly noisier with decreasing observational wavelength.  Another complication is that non-LTE effects become more important higher in the rotational ladder: distinguishing between the density effects on the line ratios and kinetic temperature effects becomes difficult.  One way to separate such effects is to observe many CO rotational lines, but this can be observationally intensive --- especially if high-J $\\cO$ lines are also observed.  Another approach to finding warm gas is using the $\\CO/\\cO$ intensity ratio. This ratio is very sensitive to the molecular gas kinetic temperature in the LTE, high-$\\Tk$ limit.  For example, the $\\COone/\\cOone$ intensity ratio in LTE goes like $\\Tk^2$ to within 25\\% for $\\Tk\\geq 20\\,K$.  While the advantage of this ratio is that it only requires two lines with very similar frequencies, the considerable disadvantage is that the $\\CO$ and $\\cO$ lines occur in very different optical depth regimes.  This results in three problems.  One is that the $\\CO$ line will be effectively probing a different volume of gas from that of the $\\cO$ line.  Another is that the levels of radiative trapping will be very different between the $\\CO$ and $\\cO$ lines, resulting in some sensitivity to density for the $\\CO/\\cO$ intensity ratio.  The third problem caused by the very different opacities in the $\\CO$ and $\\cO$ lines is that the $\\CO/\\cO$ intensity ratio will have some dependence on gas column density; indeed, in LTE this ratio is the inverse of the optical depth of the $\\cO$ line (for the $\\CO$ line in the optically thick limit and $\\cO$ in the optically thin limit).  While the problems above are not necessarily prohibitively difficult, it would be helpful nonetheless to have two optically thin tracers with very different dependences on temperature, with relative insensitivity to gas density, and that are ubiquitous in the molecular ISM.  The use of such tracers would complement the use of lines of CO and of other molecules in diagnosing the presence of warm molecular gas simply and unambiguously.  Another way to trace molecular gas, and interstellar gas in general, is to observe the continuum emission from the dust grains found throughout most of the ISM.  This continuum emission is usually observed at mid-IR, far-IR, submillimeter, and millimeter wavelengths.  On parsec scales, dust continuum emission is usually optically thin for wavelengths $\\gsim 15\\um$.  All-sky surveys like IRAS and {\\it COBE}/{\\it DIRBE} cover specific IR wavelengths from the mid-IR to far-IR and near-IR to far-IR, respectively \\citep{irasx, dirbex}. Surveys at such wavelengths in the continuum are valuable for probing the structure and excitation of the ISM \\citep[see, for example,][]{Dupac01, W96, Bally91, Zhang89, Werner76, Heiles00, Reach98, Boulanger98, Lagache98, Goldsmith97, Sodroski94, Boulanger90, Sellgren90, Scoville89, Sodroski89, Leisawitz88}. Ratios of continuum intensities at different wavelengths, for instance, can constrain estimates of dust grain temperatures.  A physical model that includes dust-gas thermal coupling could then, in theory, constrain gas temperatures \\citep[e.g., ][]{Weingartner01, Mochizuki00, Hollenbach89, Tielens85, Burke83}. Optically thin dust continuum intensity is $\\tau_\\nu B_\\nu(T_d)$, where $\\tau_\\nu$ is the optical depth at frequency $\\nu$, and $\\rm B_\\nu(\\Td)$ is the Planck function evaluated at frequency $\\nu$ and dust temperature, $\\Td$. In the far-infrared, the optical depth is often approximated as a power-law $\\tau_\\nu\\propto\\nu^\\beta$, where the emissivity index $\\beta$ depends on the size and optical properties of the grains. Consequently, as with CO rotational lines, the ratio of continuum intensities at different wavelengths loses sensitivity to dust temperature in the Rayleigh-Jeans limit.  As an example, we will consider the ratio of continuum intensities (per unit frequency) at the wavelength of 450$\\um$ to that at 850$\\um$, i.e. $\\rm I_\\nu(450\\um)/I_\\nu(850\\um)$, and adopt a 20\\% calibration uncertainty for this ratio.  We will also adopt $\\beta=2$, which is a reasonable number for wavelengths longer than about 100$\\um$ \\citep{Andriesse74, Seki80, W96}.  These adopted values imply that the $\\rm I_\\nu(450\\um)/I_\\nu(850\\um)$ ratio loses sensitivity to dust temperature for $\\Td>37\\, K$ (again from adopting 80\\% of the asymptotic value in the Rayleigh-Jeans limit).  This maximum $\\Td$ can be improved by roughly a factor of 2 by combining the 450$\\um$ observations with observations at wavelengths longer than 850$\\um$, but real improvement over this 37$\\,$K limit requires observations at wavelengths even shorter than 450$\\um$.   Observations at such short wavelengths must deal with increased atmospheric opacity and sometimes must be space-based (e.g., IRAS and {\\it COBE}).  A more direct ground-based approach is suggested by the work of \\citet{Schloerb87}, who show that a direct comparison of optically thin CO line emission with submillimeter continuum can provide temperature estimates of the dust and molecular gas in molecular clouds \\citep[see also][]{Swartz89}.  \\citet{Schloerb87} derive an expression for the millimeter continuum to integrated $\\Cotwo$ line strength ratio, $\\rm I_\\nu(submm)/I(\\Cotwo)$, which goes like $\\Tk\\Td$ in the high-temperature limit in LTE.  For similar gas and dust temperatures, this means that the $\\rm I_\\nu(submm)/I(\\Cotwo)$ ratio and, similarly, the $\\rm I_\\nu(submm)/I(\\cOone)$ ratio are actually {\\it more sensitive to temperature as that temperature increases.\\/} This is in stark contrast to ratios of rotational lines of a given isotopologue of CO and to ratios of continuum intensities at different frequencies, which, as discussed above, {\\it lose\\/} sensitivity to temperature in the high-temperature limit.  In addition, comparing the continuum intensity with that of a low-J line like the $\\Jone$ line of $\\cO$ (or of $\\Co$) reduces the sensitivity of the continuum-to-line ratio to gas density.  In short, the intensity ratio of the submillimeter continuum to an optically thin CO line may be the sought-for alternative diagnostic of molecular gas temperature that was mentioned above.  Despite the potential diagnostic advantages of the continuum-to-line ratio \\hfil\\break [e.g., $\\rm I_\\nu(submm)/I(\\cOone)$], interpretation of this ratio and its spatial variations assumes the following: \\begin{itemize} \\item {\\it The ratio of the mass of $\\cO$ to that of dust does not vary spatially.\\/} Even though the gas-to-dust mass ratio may vary, the $\\cO$-to-dust mass ratio may not vary significantly within a single source, but may vary by up to an order of magnitude from source to source \\citep{Swartz89}. \\item {\\it The gas along a given line of sight is predominantly molecular.\\/} Appreciable quantities of atomic or ionized gas and their associated dust would complicate the interpretation of $\\rm I_\\nu(submm)/I(\\cOone)$.  This is because the submillimeter continuum emission would {\\it not\\/} originate in dust {\\it only\\/} associated with molecular gas, but the $\\cO$ emission {\\it would\\/} originate only from the molecular gas.  Consulting H$\\,$I 21-cm and radio continuum data can check this assumption. \\item {\\it The molecular gas density does not vary spatially.\\/}  If this assumption is false, then spatial gradients in the $\\rm I_\\nu(submm)/I(\\cOone)$ ratio may be probing spatial gradients in the molecular gas density rather than in the molecular gas temperature.  Again, observing the lower-J rotational lines of $\\cO$, especially the $\\Jone$, line can minimize this problem. \\item {\\it The gas and dust temperatures have spatial gradients in the same direction.\\/}  Even if $\\Tk\\ne\\Td$, only assuming that the spatial gradients are in the same direction for both $\\Tk$ and $\\Td$ is necessary.  This is reasonable since both gas and dust are cooler with greater distance from heating sources, especially on multi-parsec scales. \\item {\\it Dust grain properties, like the dust mass absorption coefficient ($\\rm\\kappa_{d\\nu}$) and emissivity index ($\\beta$), do not vary spatially.\\/} Observations of Orion, for example, show that such variations do occur for $\\beta$ on sub-parsec scales \\citep[e.g.,][]{Johnstone99}, but, to first order, adopting a single value on multi-parsec scales is possible. \\item {\\it The $\\cO$ line used in the comparison is truly optically thin.\\/} If the $\\cO$ line opacity varies from optically thin to optically thick, the $\\rm I_\\nu(submm)/\\Ic$ ratio would again show gradients. The cause of the opacity change could still be due to a change in temperature, but it could also be due to a change in column density (and even volume density).  Observing the corresponding $\\Co$ line can test this. \\item {\\it There is no temperature gradient along the line of sight.\\/}  Such a gradient would result in the submillimeter emission being biased toward the warmer regions along the line of sight and the $\\cO$ emission much less so biased (and even biased toward the colder regions if LTE applies). Even in the hypothetical case of $\\Tk=\\Td$ at every point along the line of sight, these biases could lead the observer, armed with observations at additional frequencies, to believe that $\\Tk\\ne\\Td$.  This could also lead to the erroneous result that comparisons from one line of sight to another show no coupling between gas and dust temperatures.  The effects of gradients along the line of sight can be minimized by observing clouds on multi-parsec scales, where such gradients would be less extreme. \\item {\\it Thermal and non-thermal gas emission make a negligible contribution to the submillimeter continuum\\/}.  This should indeed be the case. In M$\\,$82, for example, \\citet{Thronson89} estimate that $\\sim 30\\%$ of the 1300$\\um$ flux in the central 30-45$''$ is due to gas emission. This would be even less at shorter wavelengths and in regions with no dominant supernova remnants nor HII regions. \\end{itemize} These assumptions may be reasonable first approximations for observations on multi-parsec scales or, depending on the assumption, can be tested for each source observed.  Testing the overall validity of the above assumptions and, as a result, the behavior and effectiveness of the $\\rm I_\\nu(submm)/\\Ic$ ratio as a diagnostic of molecular cloud physical conditions requires maps of a molecular cloud in the submillimeter or far-infrared (FIR) continuum, in the $\\cOone$ or $\\Coone$ spectral lines, and of dust or gas temperatures, all with resolutions at multi-parsec scales.  The most reliable maps of temperatures associated with molecular clouds would be from ratios of FIR maps.  Ignoring certain complications for the moment, such as line-of-sight gradients in dust properties (e.g., $\\Td$), uncertainties in those properties (e.g., $\\beta$), and dust grains {\\it not\\/} in thermal equilibrium, the ratio of the intensities of FIR maps at two different wavelengths gives reliable dust temperatures because such intensities are in the Wien limit of the Planck function.  In this limit, the dust temperature derived from the observed intensity ratio is insensitive to that ratio.  For example, the observed continuum intensity ratio of $\\Ia/\\Ib=1.66$ would imply $\\Td=20\\,K$.  A 10\\% uncertainty in that intensity ratio would imply only a 5\\% uncertainty in the derived dust temperature (or $\\pm 1\\unit K$).  One complication, that of stochastically-heated dust grains (i.e., grains {\\it not\\/} in thermal equilibrium), can be effectively avoided by using two FIR maps at wavelengths of $\\gsim 100\\um$ \\citep[for example, see][]{Desert90, W96}.  This, therefore, rules out using IRAS maps, whose nomimal wavelengths do not go beyond 100$\\um$. On the other hand, maps from the {\\it DIRBE\\/} instrument aboard the {\\it COBE\\/} spacecraft cover the whole sky in ten IR bands from 1.25$\\um$ out to 240$\\um$ \\citep{dirbex}.  Two maps are at nominal wavelengths of longer than 100$\\um$ (i.e., 140$\\um$ and 240$\\um$), making them virtually free of emission from stochastically-heated grains.  Thus the {\\it DIRBE\\/} maps at 140$\\um$ and 240$\\um$ and their ratio can be compared with a map of $\\cOone$ (or $\\Coone$ if available) of a given molecular cloud.  Maps of entire clouds in our Galaxy in $\\cOone$ have recently become available \\citep[e.g., see] []{Lee01, Mizuno98, Kawamura98, Nagahama98}.  Therefore, the observational data exist to test the sensitivity and reliability to molecular cloud physical conditions of the ratio of the submillimeter or FIR continuum to $\\cO$ line intensities.  The appropriate molecular cloud, or cloud{\\it s\\/}, for which this testing should be applied must satisfy a number of criteria: \\begin{itemize} \\item[---] The test clouds should be bright in the desired tracers (i.e., FIR and $\\cO$) and completely or nearly completely mapped in these tracers, \\item[---] largely mapped in tracers of atomic and ionized gas, such as HI 21-cm and radio continuum, \\item[---] out of the Galactic plane to reduce confusion with foreground and background material, \\item[---] several degrees in extent so as to accommodate many beams of the large {\\it DIRBE\\/} beam ($0\\degree.7$), \\item[---] and should have as large a range of dust temperatures as possible at the {\\it DIRBE\\/} resolution. \\end{itemize} The clouds of Orion most closely satisfy these criteria \\citep[see, e.g.,][]{Bally91, W96, Nagahama98, Zhang91, Green91, Chromey89, Heiles74, Reich78, Berkhuijsen72, Haslam70}.  While there are other clouds that {\\it may seem\\/} to satisfy these criteria as well \\citep[e.g., Chamaeleon-Musca,][]{Mizuno98}, the Orion clouds have the largest range of dust temperatures at the {\\it DIRBE\\/} resolution \\citep[see Figure~2 of][]{Lagache98}.  The Orion$\\,$A and B molecular clouds are bright in $\\cOone$ \\citep{Nagahama98}, and in the FIR continuum at 140$\\um$ and 240$\\um$ \\citep{W96}, and has been completely or largely mapped in these tracers. Much of the Orion constellation has been mapped in H$\\,$I \\citep{Zhang91, Green91, Chromey89, Heiles74} and radio continuum \\citep{Reich78, Berkhuijsen72, Haslam70}, thereby permitting tests of whether or not molecular gas dominates the cloud surface, or column, density for each line of sight.  Much of Orion has also been mapped in $\\CO$ lines \\citep{Wilson05, Lang00, Ikeda99, Sakamoto97a, Sakamoto94, Maddalena86}, which can also provide crude estimates of the molecular gas column densities.  Note that only \\citet{Maddalena86} have mapped the Orion all three major clouds: Orion$\\,$A, Orion$\\,$B and the $\\lambda\\,$Orionis ring.  The Orion clouds are about 15-20$^\\circ$ out of the Galactic plane and about $6^\\circ\\times 3^\\circ$ (or about $50\\unit pc\\times 25\\, pc$ for an adopted distance of 450$\\,$pc) in size, down to some ill-defined zero level.  An additional advantage to the clouds of Orion is that these clouds have been surveyed at many other wavelengths on the scales of many parsecs, thereby providing further checks on models.  These other wavelengths include, for example, the optical surveys of the Orion OB1 and $\\lambda\\,$Ori OB associations \\citep{Brown95, Brown94, Warren78, Warren77, Murdin77, Blaauw64}, the near-IR to far-IR all-sky survey of {\\it COBE}/{\\it DIRBE\\/} \\citep{W96}, the mid-IR to far-IR all-sky survey of IRAS \\citep{Bally91}, and submillimeter maps of small sections of Ori$\\,$A and B \\citep[e.g.,][]{Johnstone01, Mitchell01, Johnstone99}.  Therefore, the Orion$\\,$A and B molecular clouds were chosen for testing the behavior of the $\\rm I_\\nu(submm)/\\Ic$ ratio, using the {\\it DIRBE\\/} maps at 140$\\um$ and 240$\\um$ and the $\\cOone$ data of \\citet{Nagahama98}.  In addition to testing the diagnostic abilities of the continuum-to-line ratio, the current data may also permit us to test theoretical models describing interactions between dust and gas.   Specifically, models of dust-gas thermal coupling use very simple assumptions, such as large, single-size, spherical dust grains with smooth surfaces \\citep[e.g., see][]{Goldsmith01, Burke83}.   Are these simplistic assumptions sufficient for explaining the observations?   This will be examined.  The current paper, or Paper~I, describes the one- and two-component modeling of the data. Paper~II \\citep{W05} in this series confirms and, in some cases, modifies the model results using simulated data.  Paper~III \\citep{W05a} examines systematic effects not explicitly considered in the models and also discusses the scientific implications of the results. ",
        "conclusions": "}  We used the 240$\\um$ continuum to the $\\cOone$ intensity ratio plotted against the 140$\\um$/240$\\um$ dust color as a diagnostic of the dust and gas physical conditions in the Orion$\\,$A and B molecular clouds.  Model curves were fitted to the data with orthogonal regression.  Two sets of one-component models were applied to the data: Local Thermodynamic Equilibrium (LTE) and Large Velcocity Gradient (LVG).  Two sets of two-component models were also applied to the data: two-component, LVG models and two-component, two-subsample, LVG models.  In each case, the model fit gave the conversion of the observed intensities into gas/dust column densities and beam-area filling factors.  The results of the fitting are summarized below: \\begin{itemize} \\item {\\it One-Component, LTE Models:\\/} These models fit the data only poorly with a reduced chi-square of $\\chi_\\nu^2=16.5$.  When considering systematic uncertainties, the only parameter to be fitted, $\\DT$, was very poorly determined: $\\DT$ lies anywhere between $-24\\,$K and $+6\\,$K.  Comparison of the continuum-derived gas column densities with those derived from the $\\cOone$ spectral line showed considerable scatter with $\\chi_\\nu^2=16.6$. The filling factors exhibit an understandable increase with increasing column density and are all less than unity.  At best, the one-component, LTE models only crudely reproduce the overall trend in the data. \\item {\\it One-Component, LVG Models:\\/} These models also fit the data only poorly with a reduced chi-square of $\\chi_\\nu^2=16.9$.  When considering systematic uncertainties, the $\\DT$ was again very poorly determined: $\\DT$ lies anywhere between $-14\\,$K and $-1\\,$K.  Comparison of the continuum-derived versus spectral line-derived gas column densities showed comparable scatter to the LTE model case, with $\\chi_\\nu^2=15.5$.  The filling factors behaved similarl to those in the case of the LTE modeling.  The LVG model results overall are, at best, only marginally better than the single-component, LTE models. \\item {\\it Two-Component, LVG Models:\\/} In contrast to the one-component models, these models explain the overall trend in the data with $\\chi_\\nu^2=5.7$; this is better than the one-component models at a confidence level of better than 99.9\\% according to the F-test.  Also in contrast to the one-component model results, $\\DT$ is known to be within 1$\\,$K of 0$\\,$K even when the systematic uncertainties are considered.  The component-0 dust temperature {\\it seems\\/} to be accurately known and is consistent with the large-scale interstellar radiation field, i.e. 18$\\,$K.  Other parameters, such as gas densities and column densities per velocity interval, are known from factors of a few to more than an order of magnitude about some geometric mean.  The continuum-derived versus spectral line-derived gas column densities showed much better agreement for these models than for the one-component component models, the former having $\\chi_\\nu^2=7.4$.  Despite their advantages, these models suffer from two problems.  The first is that many of the filling factors have unphysical values --- i.e., greater than unity.  The second problem is that the model curves do not fit well to the data with 140$\\um$/240$\\um$ color temperatures greater than 20$\\,$K. \\item {\\it Two-Component, Two-Subsample, LVG Models:\\/} Here the sample of data points was split into two subsamples: one with dust color temperatures $<20\\,$K and the other with such temperatures $\\ge 20\\,$K.  These models have all the advantages of the simpler two-component models described previously, and also fix the problems of the unphysical filling factors and the poor fit to the data points with 140$\\um$/240$\\um$ color temperatures $\\ge 20\\,$K.  The only disadvantage is that this approach is ad hoc.  Nevertheless, the derived parameter values agree with those of the simpler two-component models to within the systematic uncertainties. \\item {\\it Cloud Masses:\\/} After correcting for systematic effects, such as the lack of coverage of the $\\cOone$ map, the total {\\it hydrogen\\/} gas mass derived in the Orion fields is $4.2\\times 10^5\\unit M_\\odot$, as found in the more realistic two-component case (and is $2.5\\times 10^5\\unit M_\\odot$ for the less realistic one-component case).  Due to the systematics, the uncertainty in this mass is $\\pm 40\\%$.  The extra mass in the two-component case is due to the apparent presence of cold dust/gas with temperatures of 3-10$\\,$K.  Attempting to confirm the existence of this cold dust/gas with the long-wavelength FIRAS data was inconclusive due to the insufficient spatial resolution. \\end{itemize}  The above results have a number of astronomical implications, which are mentioned here only briefly: \\begin{enumerate} \\item {\\it Dust Temperatures Equal to Gas Temperatures:\\/} If this surprising result applies to the gas and dust on large scales in our Galaxy, then models of gas-to-dust thermal coupling must be revised.  Typically, these models assume spherical grains with a single large grain size and a smooth grain surface \\citep[e.g., see][]{Goldsmith01, Burke83}.  Clearly, relaxing these simple assumptions will increase the surface area-to-volume ratio of the grains and improve the dust-gas thermal coupling, reducing $\\DT$ (see Papers~III and IIIa for more details). Having dust and gas temperatures equal on galactic scales implies that large-scale molecular gas temperatures are nearly double the temperatures previously believed (i.e. closer to 20$\\,$K than to 10$\\,$K).  Also, one dust grain alignment mechanism would be ruled out. \\item {\\it At least two temperatures are needed along each sightline to give reasonable column densities:\\/}  The two-component models show good agreement between spectral line-derived and continuum-derived column densities.   The main difference between the two components is that one component has unvarying temperature from sightline to sightline and the other component has varying temperature.  Therefore, except for those few sightlines where the two components have the same temperature, each sightline has two temperatures (i.e. one temperature --- gas and dust --- for each component).   Requiring more than one temperature along each sightline is supported by the work of \\citet{Schnee06}, who find that assuming a single temperature along each sightline does not give reliable column densities in the Perseus and Ophiuchus molecular clouds. \\item {\\it Cold Gas/Dust:\\/} If similar two-component models are correct and are applicable to other large clouds in the Galaxy, then the total mass of molecular gas in the Galaxy must be revised upwards by about 60\\%. \\item {Millimeter Continuum to $\\cOone$ Line Ratio as a Temperature Diagnostic:\\/} For the millimeter continuum to $\\cOone$ intensity ratio above some threshold, this ratio can serve as a diagnostic of the gas/dust temperature, even in the high-temperature limit. \\item {\\it X-Factor:\\/} An improved explanation for the X-factor is applicable.  This will be discussed in detail in \\citet{W05b}. \\end{enumerate}  Discussing the above results in depth would be premature at this point.  We must first examine the modeling and various systematic effects more carefully.  Even though we have tested the effects of some systematics on the model curves (i.e., by applying scale factors or changing the starting grid when fitting the models), we need to find and understand any inherent biases or flaws that the model fitting may have.  For example, the two-component models consistently find a $\\DT$ equal to or near 0$\\,$K, in spite of the scale factor applied to the data.  Is this because the dust-gas temperature difference really is zero?  Or is it simply because these models are biased towards $\\DT=0\\,$K despite its true value?  And what of the reliability of the other parameter values?  The component-1 density, for example, is supposed to be known to within a factor of a few of 600$\\,H_2$$\\unit cm^{-3}$, according to the simple two-component models (see Figure~\\ref{fig27}).  And yet, the two-subsample, two-component models suggest densities that are one to three orders of magnitude higher than this.  Clearly, the methods used thus far to examine the systematics are incomplete.  Consequently, Paper~II examines the results of running the models on simulated data and will answer the concerns raised here.  In addition to those concerns, there are other questions not yet addressed by the methods used here.  Specifically, what is the effect of the background subtractions used?  How will dust associated with HI affect the results?  Does changing the spectral emissivity index $\\beta$ appreciably affect the results?  Are there alternative kinds of models that would also explain the data?  The cloud positions modeled here only represent 26\\% of the area of the Orion clouds; how representative are the results of the clouds as a whole?  Paper~III examines these questions.  All these questions and concerns are addressed in Papers~II and III.  Paper~III also discusses the scientific implications of the results in more detail."
    },
    "0601/astro-ph0601235_arXiv.txt": {
        "abstract": "This work presents the main ultraviolet (UV) and far-infrared (FIR) properties of two samples of nearby galaxies selected from the GALEX ($\\lambda = 2315$\\AA, hereafter NUV) and IRAS ($\\lambda = 60\\mu$m) surveys respectively. They are built in order to get detection at both wavelengths for most of the galaxies. Star formation rate (SFR) estimators based on the UV and FIR emissions are compared. Systematic differences are found between the SFR estimators for individual galaxies based on the NUV fluxes corrected for dust attenuation and on the total IR luminosity. A combined estimator based on NUV and IR luminosities seems to be the best proxy over the whole range of values of SFR. Although both samples present similar average values of the birthrate parameter b, their star-formation-related properties are substantially different: NUV-selected galaxies tend to show larger values of $b$ for lower masses, SFRs and dust attenuations, supporting previous scenarios for the star formation history (SFH). Conversely, about 20\\% of the FIR-selected galaxies show high values of $b$, SFR and NUV attenuation. These galaxies, most of them being LIRGs and ULIRGs, break down the downsizing picture for the SFH, however their relative contribution per unit volume is  small in the local Universe. Finally, the cosmic SFR density of the local Universe is estimated in a consistent way from the NUV and IR luminosities. ",
        "introduction": "What is the best way to measure the SFR of galaxies on large scales and at different redshifts?  The possibility of estimating the SFR of a galaxy directly from the luminosity at a single wavelength would be a major advantage for anyone wanting to compute the SFR per unit volume at a given redshift. This quantity could be derived directly from the luminosity function (LF) at this wavelength and at this redshift. Under these conditions, large area surveys at single wavelengths might suffice.  The recent SFR of a galaxy is often measured from the light emitted by the young stars: given their short lifetimes their luminosity is directly proportional to the rate at which they are currently forming. The UV and FIR luminosities of star forming galaxies are both closely related to the recent star formation: most of the UV photons are originally emitted by stars younger than $\\sim 10^{8}$~yr, but many of these photons are re-processed by the dust present in galaxies and re-emitted at FIR wavelengths. Strictly speaking, neither of these fluxes can be used alone to estimate the SFR independently of the other one (e.g. Buat \\& Xu 1996; Hirashita et al. 2003; Iglesias-P\\'{a}ramo et al. 2004). Because of the previous lack of data at both wavelengths, attempts have been made using only the rest frame UV (Lilly et al. 1996; Madau et al. 1996; Steidel et al. 1999; and more recently Schiminovich et al. 2005 with GALEX data), or just FIR data (Rowan-Robinson et al. 1997; Chary \\& Elbaz 2001), but only a few authors have compared both (Flores et al. 1999, Cardiel et al. 2003). The SFR estimator based on the UV luminosity suffers from attenuation by dust and it has to be corrected  in order to properly trace the SFR: for UV-selected samples of galaxies the attenuation can reach more than 1~mag (e.g. Iglesias-P\\'{a}ramo et al. 2004; Buat et al. 2005). On the other hand the FIR emission is not free of problems because the dust can also be heated by old stars, and can be a non-negligible correction for many star forming galaxies (Lonsdale \\& Helou 1987; Sauvage \\& Thuan 1992). Neither of these two indicators taken alone is an accurate estimator of the SFR except perhaps for starburst galaxies where (a) the dust attenuation was found to follow a tight relation with the slope of the spectrum at UV wavelengths (Meurer et al. 1999), thus allowing one to estimate the dust attenuation with only information on UV fluxes, and (b) the contribution to the dust emission coming from old stars can be neglected (Sauvage \\& Thuan 1992). In the most general case, the best estimator of the SFR should contain combined information of the luminosities at both wavelength ranges (Hirashita et al. 2003). The UV and FIR fluxes are thus complementary for tracing star formation and it is well known that the FIR/UV ratio is a proper indicator of the dust attenuation (Buat et al. 1999; Witt \\& Gordon 2000; Panuzzo et al. 2003). Although other indicators of the recent SFR of galaxies have been extensively used in the literature, a detailed discussion of their quality as SFR tracers will not be discussed in this paper.  The GALEX mission (Martin et al. 2005a) is imaging the high-Galactic latitude sky at two UV wavelengths ($\\lambda=1530$\\AA, FUV; $\\lambda=2315$\\AA, NUV) and is providing the astronomical community with unprecedented data (both in quantity and quality). The UV data combined with existing FIR datasets (from the IRAS, ISO or Spitzer missions) now allow us to carry out detailed studies of the UV and FIR properties of galaxies, with special emphasis on the derivation of the dust attenuation and star formation activity in star forming galaxies.  With this purpose in mind we selected two samples of galaxies: one from the GALEX-All Imaging Sky (AIS) and the other from the IRAS PSCz (redshifts, infrared and optical photometry, and additional information for 18,351 IRAS sources, mostly selected from the Point Source Catalog) and FSC (Faint Source Catalog) for which UV and FIR fluxes are available. With these datasets in hand we undertaken a study of the properties related to their emission at these wavelengths. Both samples were extracted from the same region of the sky ($\\sim 600$~degrees$^{2}$, constrained by the status of the GALEX survey when this work was initiated). From the GALEX catalog we built a complete sample of galaxies down to $AB_{NUV}=16$~mag\\footnote{AB magnitudes are defined as: $AB_{\\nu} = -2.5 \\log f_{\\nu} -48.6$, where $f_{\\nu}$ is the monochromatic flux density expressed in erg~s$^{-1}$~cm$^{-2}$~Hz$^{-1}$.} and cross-correlated it with the IRAS database (FSC), allowing non-detections at 60$\\mu$m (fluxes lower than 0.2 Jy). The FIR-selected sample was built from the IRAS catalog (PSCz) with a limiting flux at 60$\\mu$m of 0.6~Jy as the only constraint. The resulting list of PSCz sources was cross-correlated with the GALEX database, allowing again non-detections in the NUV. Both the NUV and the 60$\\mu$m limits used to build the samples were chosen to allow for a very small number of non-detections  and to  sample the galaxy population over a large range of values of the dust attenuation. Besides an  analysis of the SFR, for  these NUV and FIR-selected samples of galaxies (chosen with well-defined selection criteria) can also be used to place  important constraints on  models designed to predict the statistical properties of galaxy populations. The attenuation of star light by interstellar dust, and its emission in the far infrared are usually computed very crudely in models of galactic evolution. Dust attenuation is usually deduced in such approaches from other quantities such as the mass gas and the metallicity (e.g. Guiderdonni \\& Rocca-Volmerange 1987; Devriendt \\& Guiderdoni 2000; Balland et al. 2003). The properties of large samples of galaxies observed both in the UV and FIR with clear selection criteria, such as the ones presented in this paper, provide an important statistical constraint for the calibration of the treatment of dust in such models.  The first paper in this series (Buat et al. 2005) based on these samples was mainly devoted to the dust attenuation properties. In the present work we discuss various aspects  relating to the NUV and 60$\\mu$m emission of our sample galaxies, including systematic differences in the 60$\\mu$m and NUV luminosities and we will focus on their star formation related properties. This paper is organized as follows: the samples are presented in section~2. The relation between NUV and 60$\\mu$m luminosities is discussed in section~3. Section~4 is devoted to the derivation of the SFR and to  a comparison of various estimators as well as to a discussion on the star formation activity related properties of the samples. The derivation of the local cosmic SFR density by different methods is discussed in section~5. The main conclusions are presented in section~6. Throughout this paper we adopt the following cosmological parameter set: $(h, \\Omega_0, \\lambda_0)=(0.7,0.3,0.7)$, where $h\\equiv H_0/100\\; (\\mbox{km\\,s}^{-1}\\mbox{Mpc}^{-1})$. ",
        "conclusions": "We have performed a detailed study of the star formation properties of two samples of galaxies selected on the basis of their NUV and FIR fluxes, which were found to be representative of the nearby Universe when compared to samples drawn from larger volumes. The main conclusions of this work are: \\begin{enumerate} \\item $L_{NUV}$ and $L_{60}$ are tightly correlated for the NUV-selected galaxies. The opposite holds for FIR-selected galaxies, which span a large range of $L_{60}$ for a given value of $L_{NUV}$ and show larger values of attenuation. Intrinsically bright galaxies are more luminous at FIR than at NUV wavelengths, including the UV luminous galaxies (UVLGs), and they show moderate to high attenuation. \\item The SFR deduced from the NUV fluxes, corrected for the dust attenuation ($SFR_{NUV}$), are not found to be consistent with those calculated using  the total dust emission ($SFR_{dust}$). Whereas $SFR_{NUV}$ is larger than $SFR_{dust}$ for galaxies with low attenuation ($A_{NUV} \\lesssim 1$~mag) the inverse is found for bright, but highly extinguished galaxies, mostly selected in IR: $SFR_{NUV}$ is likely to underestimate the actual SFR in these galaxies by a factor $\\sim 2$. A combined estimator based on UV and IR luminosities with a cirrus correction depending on the IR luminosity seems to be the best proxy over the whole range of values of SFR. As a practical recipe we found that $SFR_{tot}$ and $SFR_{NUV}$ yield similar results for $SFR_{tot} \\lesssim 15$~M$_{\\odot}$~yr$^{-1}$, whereas $SFR_{tot}$ and $SFR_{dust}$ are almost equivalent for $SFR_{tot} \\gtrsim 15$~M$_{\\odot}$~yr$^{-1}$. \\item NUV-selected galaxies follow a trend whereby low-mass galaxies show lower SFRs, low attenuation and higher values of $b$, indicating the existence of a dominant young stellar population. On the contrary, about 20\\% of the FIR-selected sample shows  high attenuation, high SFRs and also large values of $b$, most of them being LIRGs and/or ULIRGs. In spite of their discordant properties, these galaxies are not sufficiently abundant in the local Universe to question the downsizing picture for the SFH seen at $z = 0$ from optical surveys. \\item The cosmic SFR densities of the local Universe, estimated from the NUV and IR luminosities, are consistent to within 1-$\\sigma$, although the difference between the two values is large, when average corrections for the attenuation, UV escaping photons and IR cirrus component are applied. The sum of the individual contributions is quite consistent with the value obtained from the NUV luminosities corrected for average attenuation. A better knowledge of the cirrus contribution to $L_{IR}$ and of the bivariate LF is required in order to better understand the large differences found between the monochromatic estimators of the local SFR density. \\end{enumerate}"
    }
}